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Massive Stars: Life and Death

Dissertation

Presented in Partial Fulfillment of the Requirements for the Degree Doctor ofPhilosophy in the Graduate School of The Ohio State University

By

Jose Luis Prieto Katunaric

Graduate Program in Astronomy

The Ohio State University

2009

Dissertation Committee:

Professor Krzysztof Z. Stanek, Advisor

Professor Christopher S. Kochanek

Professor John F. Beacom

Copyright by

Jose L. Prieto

2009

ABSTRACT

Although small in number, massive stars are critical to the formation and

evolution of galaxies. They shape the interstellar medium of galaxies through

their strong winds and ultra-violet radiation, are a major source of the heavy

elements enriching the interstellar medium, and are the progenitors of core-collapse

supernovae and gamma-ray bursts, which are among the most energetic explosions

in the Universe and mark the death of a massive star. Still, our understanding of the

connection between massive stars and supernovae from observations is fairly limited.

In this dissertation, I present new observational evidence that shows the importance

of metallicity, mass-loss, and binarity in the lives and deaths of massive stars.

We investigate how the different types of supernovae are relatively affected by

the metallicity of their host galaxy. We take advantage of the large number of spectra

of star-forming galaxies obtained by the Sloan Digital Sky Survey and their overlap

with supernova host galaxies. We find strong evidence that type Ib/c supernovae are

occurring in higher-metallicity host galaxies than type II supernovae. We discuss

various implications of our findings for understanding supernova progenitors and

their host galaxies, including interesting supernovae found in low-metallicity hosts.

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We present the discovery of the progenitors of SN 2008S and the luminous

transient in NGC 300 in archival data obtained with the Spitzer Space Telescope.

They are deeply dust-enshrouded massive stars, with extremely red mid-infrared

colors compared to other massive stars, and relatively low bolometric luminosities

(≈ 5 × 104 L⊙). We discuss the implications of these findings for the evolution

and census of “low-mass” massive stars (i.e., ∼8 − 12 M⊙), and we connect it with

theoretical discussions of electron-capture supernovae near this mass range, explosive

birth of massive white dwarfs, and massive star outbursts.

We present a Spitzer low-resolution mid-infrared spectrum of the luminous

transient in NGC 300. The spectrum shows that the transient is very luminous in

the mid-infrared and most of the pre-existing progenitor dust survived the explosion.

Furthermore, the spectrum shows strong, broad emission features that are observed

in Galactic carbon-rich proto-planetary nebulae. These observations support our

conclusions of an explosive event on a massive carbon-rich AGB or post-AGB star

as the origin of the transient in NGC 300 and SN 2008S.

We present extensive ugriz Y HJKs photometry and optical spectroscopy of

SN 2005gj. These data show that SN 2005gj is the second possible case, after

SN 2002ic, of a thermonuclear explosion in a dense circumstellar environment. The

interaction of the supernova ejecta with the dense circumstellar medium is stronger

than in SN 2002ic.

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Finally, we present the discovery of a peculiar eclipsing binary in a variability

survey of the dwarf galaxy Holmberg IX undertaken with the Large Binocular

Telescope. The binary has a period of 271 days, and is composed of two yellow

supergiants that are overflowing their Roche lobes. Such systems must be rare, and

indeed we only note one similar system in the Small Magellanic Cloud. We propose

that these systems may be the progenitors of supernovae that appeared to have

yellow-supergiant progenitors.

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A mis padres, Alicia y Rodolfo.

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ACKNOWLEDGMENTS

I am very grateful of so many people that in different ways have helped me

these years during grad school and before. In particular, I don’t think I would have

made it without the continuous support and encouragement of my adviser, Kris

Stanek, my family, and Linda. First let me say Gracias to Kris, I will come back to

my family and Linda at the end. I feel very fortunate that I had the chance to meet

Kris and work with him over the last three years. Kris has shaped my way to see,

think about, and do astronomy and science in general. He has always been the most

supportive person, almost in unthinkable ways by (for example) letting me continue

working in supernova-related topics and encouraging me to pursue exciting ideas and

projects that came up along the way. We had many discussions about astronomy

and science, new ideas for research projects, and his yard (among other topics). I

enjoyed very much all of them. He also shared with me his vast experience and

knowledge of gamma-ray bursts, the distance scale, variable stars, and photometry

(to name just a few). His stories and experiences of other people working in the

field and how “not to do” science were always very enriching. I also enjoyed playing

ping-pong and sledding at his house. I will miss him much.

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I am also very grateful of John Beacom, Chris Kochanek, and Todd Thompson.

They were always happy to help, support, and encourage me. I learned a lot

from many extensive discussions we had, including Kris, about massive stars and

supernovae. Chris, with his amazing fundamental knowledge of astronomy and

physics, always had time to share his thoughts and help with papers and new ideas.

I appreciate very much the advice and help of John and Chris with my English

and scientific writing. I also thank Chris and John for being part of my defense

committee. I think it will be hard to find a person with more energy and excitement

about astronomy than Todd, and I hope I can at least have a fraction of that.

Most of the theoretical interpretation in this dissertation, particularly about the

dusty transients, were Todd’s ideas. Mi amigo Todd, thank you for sharing all that

knowledge about several areas in astrophysics and for so many discussions about

science and Espanol. I hope that Todd at least learned this from me, and doesn’t

forget it: Pablo Neruda was Chilean.

I thank Oleg Gnedin for advising me in my first project at OSU. It was hard to

work in a purely theoretical project on globular clusters, but with a lot of help, advice

and patience of Oleg, we could finally finish it and publish a paper on the results.

My big thanks to Darren DePoy for letting me work in the SDSS-II Supernova

Search, for sharing his observing experience and stories, and for organizing the

supernova follow-up campaigns with the MDM 2.4m telescope. Here I have to give

special thanks to all my fellow graduate students that went observing to MDM and

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participated in the 3 campaigns, to all the MDM staff for maintaining the telescopes

and instruments and also for their help, and to Rick Pogge for being always there to

help when something came up with the instrument, telescope or control system. It

wouldn’t have been possible to carry out this extensive follow-up project without

them. I thank the SDSS-II Supernova Search and the Carnegie Supernova Project

groups for contributing most of the photometric data for SN 2005gj. I learned a

lot from working in the SDSS-II Supernova Search, and I am especially grateful

to Masao Sako and Peter Garnavich. I thank Kris Sellgren for sharing part of her

experience and knowledge of infrared astronomy with me. Kris was very enthusiastic

to teach me about mid-infrared spectra of proto-planetary nebula and other infrared

sources.

Going all the way back to high school, I thank my physics teacher Jimena

Scarich, she was very important in my early science education and interest in

physics and astronomy. I would like to thank also my undergraduate adviser at

Universidad Catolica, Alejandro Clocchiatti, who introduced me to the exciting field

of supernovae, always supported me, and helped me get a position doing supernova

research at CTIO after undergrad. At Catolica, I also enjoyed very much working

with Felipe Barrientos, and I thank him for teaching me how to use Perl and PDL,

which had been very useful tools since. At CTIO, working with Nick Suntzeff and

Armin Rest was a really great experience. I am indebted to Nick and Armin (and

also others of the supernova group in Tololo, especially Kevin Krisciunas) for giving

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me the great opportunity to work in very exciting scientific projects during that year

and for introducing me to participate in two big scientific collaborations, Essence

and Supermacho.

The environment at the Astronomy Department was always nice and friendly,

and this made a big difference in the day-to-day life. I thank the Department

office staff, especially Kristy Scheckelhoff, for helping so much with all the paper

work. Also David Will and Michael Savage for so much help and assistance with

computers. I am also grateful of all the faculty for organizing morning coffee, I

enjoyed very much many discussions and learned quite a bit about very different

areas in astronomy. I also appreciate all the work put by faculty to make the

graduate courses useful and interesting (including order of magnitude) and I thank

them for that. I thank many conversations over lunches and dinners with fellow

graduate students and friends, especially Julio Chaname, Roberto Assef, and Subo

Dong. Outside work, I very much enjoyed spending time with Max Montenegro and

family, they always made me feel like at home.

I would like to thank my family, especially my parents to whom I would like to

dedicate this dissertation. Their endless love, care, guidance, and support without

bounds made me get to this point and allowed me to follow my dream of doing

astronomy, even though it took me so far from them. Muchas Gracias Papis. Gracias

tambien a mis hermanos Cristian, Eugenio, Rodolfo, y a los chicos. Finally, it’s hard

to describe in just a few words how lucky I feel to have met Linda and thank her

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enough for all her encouragement, support, love, patience, and care over the last

years. Thank you so much.

x

VITA

November 8, 1980 . . . . . . . . . . . . . . Born – Punta Arenas, Chile

2003 . . . . . . . . . . . . . . . . . . . . . . . . . . . B.S. Astronomy, P. Universidad Catolica de Chile

2003 – 2004 . . . . . . . . . . . . . . . . . . . . Research Assistant, Cerro TololoInter-American Observatory, La Serena, Chile

2004 – 2008 . . . . . . . . . . . . . . . . . . . . Graduate Research Associate,The Ohio State University

2008 – 2009 . . . . . . . . . . . . . . . . . . . . . Presidential Fellow, The Ohio State University

PUBLICATIONS

Research Publications

1. K. Krisciunas et al. (31 authors, incl. J. L. Prieto), “Optical and In-frared Photometry of the Type Ia Supernovae 1991T, 1991bg, 1999ek, 2001bt,2001cn, 2001cz, and 2002bo”, AJ, 128, 3034, (2004).

2. T. Matheson et al. (31 authors, incl. J. L. Prieto), “Spectroscopy ofHigh-Redshift Supernovae from the ESSENCE Project: The First 2 Years”, AJ, 129,2352, (2005).

3. K. Krisciunas et al. (32 authors, incl. J. L. Prieto), “Hubble Space Tele-scope Observations of Nine High-Redshift ESSENCE Supernovae”, AJ, 130, 2453(2005).

4. A. Rest et al. (23 authors, incl. J. L. Prieto), “Testing LMC Microlens-ing Scenarios: The Discrimination Power of the SuperMACHO MicrolensingSurvey”, ApJ, 634, 1103 (2005).

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5. A. Rest et al. (16 authors, incl. J. L. Prieto), “Light echoes from an-cient supernovae in the Large Magellanic Cloud”, Nature, 438, 1132, (2005).

6. K. Krisciunas, J. L. Prieto, P. M. Garnavich, J-L G. Riley, A. Rest, C.Stubbs, and R. McMillan, “Photometry of the Type Ia Supernovae 1999cc, 1999cl,and 2000cf”, AJ, 131, 1639, (2006).

7. S. Blondin et al. (37 authors, incl. J. L. Prieto), “Using Line Profiles toTest the Fraternity of Type Ia Supernovae at High and Low Redshifts”, AJ, 131,1648, (2006).

8. A. Clocchiatti et al. (29 authors, incl. J. L. Prieto), “Hubble Space Tele-scope and Ground-based Observations of Type Ia Supernovae at Redshift 0.5:Cosmological Implications”, ApJ, 642, 1, (2006).

9. M. Modjaz et al. (13 authors, incl. J. L. Prieto), “Early-Time Photome-try and Spectroscopy of the Fast Evolving SN 2006aj Associated with GRB 060218”,ApJL, 645, 21, (2006).

10. J. L. Prieto, A. Rest, and N. B. Suntzeff, “A New Method to Calibratethe Magnitudes of Type Ia Supernovae at Maximum Light”, ApJ, 647, 501, (2006).

11. K. Krisciunas et al. (12 authors, incl. J. L. Prieto), “The Type Ia Su-pernova 2004S, a Clone of SN 2001el, and the Optimal Photometric Bands forExtinction Estimation”, AJ, 133, 58, (2007).

12. K. Z. Stanek et al. (12 authors, incl. J. L. Prieto), “ “Anomalous” Op-tical Gamma-Ray Burst Afterglows Are Common: Two z 4 Bursts, GRB 060206and GRB 060210”, ApJL, 654, 21, (2007).

13. A. Garg et al. (21 authors, incl. J. L. Prieto), “Light Curves of TypeIa Supernovae from Near the Time of Explosion”, AJ, 133, 403, (2007).

14. S. A. Yost et al. (47 authors, incl. J. L. Prieto), “Exploring BroadbandGRB Behavior during γ-Ray Emission”, ApJ, 657, 925, (2007).

15. M. M. Phillips et al. (44 authors, incl. J. L. Prieto), “The Peculiar SN2005hk: Do Some Type Ia Supernovae Explode as Deflagrations ?”, PASP, 119, 360,(2007).

16. G. Miknaitis et al. (36 authors, incl. J. L. Prieto), “The ESSENCE Su-

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pernova Survey: Survey Optimization, Observations, and Supernova Photometry”,ApJ, 666, 674, (2007).

17. W. M. Wood-Vasey et al. (37 authors, incl. J. L. Prieto), “Observa-tional Constraints on the Nature of Dark Energy: First Cosmological Results fromthe ESSENCE Supernova Survey”, ApJ, 666, 694, (2007).

18. M. Hicken, P. M. Garnavich, J. L. Prieto, S. Blondin, D. L. DePoy, R.P. Kirshner, and J. Parrent, “The Luminous and Carbon-rich Supernova 2006gz: ADouble Degenerate Merger ?”, ApJL, 669, 17, (2007).

19. J. Frieman et al. (101 authors, incl. J. L. Prieto), “The Sloan DigitalSky Survey-II Supernova Survey: Technical Summary”, AJ, 135, 338, (2008).

20. M. Sako et al. (50 authors, incl. J. L. Prieto), “The Sloan Digital SkySurvey-II Supernova Survey: Search Algorithm and Follow-Up Observations”, AJ,135, 348, (2008).

21. J. L. Prieto et al. (17 authors), “LBT Discovery of a Yellow SupergiantEclipsing Binary in the Dwarf Galaxy Holmberg IX”, ApJL, 673, 59, (2008).

22. J. L. Prieto, K. Z. Stanek, and J. F. Beacom, “Characterizing Super-nova Progenitors via the Metallicities of their Host Galaxies, from Poor Dwarfs toRich Spirals”, ApJ, 673, 999, (2008).

23. S. Brady, J. R. Thorstensen, M. D. Koppelman, J. L. Prieto, P. M. Gar-navich, A. Hirschauer, and M. Florack, “The Eclipsing Cataclysmic Variable Lanning386: Dwarf Nova, SW Sextantis Star, or Both ?”, PASP, 120, 301, (2008).

24. M. Modjaz et al. (9 authors, incl. J. L. Prieto), “Measured Metallicitiesat the Sites of Nearby Broad-Lined Type Ic Supernovae and Implications for theSupernovae Gamma-Ray Burst Connection”, AJ, 135, 1136, (2008).

25. B. S. Gaudi et al. (22 authors, incl. J. L. Prieto), “Discovery of a VeryBright, Nearby Gravitational Microlensing Event”, ApJ, 677, 1268, (2008).

26. C. Zheng et al. (80 authors, incl. J. L. Prieto), “First-Year Spectroscopy for theSloan Digital Sky Survey-II Supernova Survey”, AJ, 135, 1766, (2008).

27. S. F. Anderson et al. (21 authors, incl. J. L. Prieto), “Two More Can-didate AM Canum Venaticorum (am CVn) Binaries from the Sloan Digital SkySurvey”, AJ, 135, 2108, (2008).

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28. A. Rest et al. (29 authors, incl. J. L. Prieto), “Spectral Identificationof an Ancient Supernova Using Light Echoes in the Large Magellanic Cloud”, ApJ,680, 1137, (2008).

29. J. L. Prieto et al. (10 authors), “Discovery of the Dust-Enshrouded Pro-genitor of SN 2008S with Spitzer”, ApJL, 681, 9, (2008).

30. A. Rest et al. (19 authors, incl. J. L. Prieto), “Scattered-Light Echoesfrom the Historical Galactic Supernovae Cassiopeia A and Tycho (SN 1572)”, ApJL,681, 81, (2008).

31. B. Dilday et al. (53 authors, incl. J. L. Prieto), “A Measurement of theRate of Type Ia Supernovae at Redshift z 0.1 from the First Season of the SDSS-IISupernova Survey”, ApJ, 682, 262, (2008).

32. A. C. Becker et al. (39 authors, incl. J. L. Prieto), “Exploring theOuter Solar System with the ESSENCE Supernova Survey”, ApJL, 682, 53, (2008).

33. X. Dai et al. (23 authors, incl. J. L. Prieto), “Go Long, Go Deep:Finding Optical Jet Breaks for Swift-Era GRBs with the LBT”, ApJL, 682, 77,(2008).

34. C. S. Kochanek, J. F. Beacom, M. D. Kistler, J. L. Prieto, K. Z. Stanek, T.A. Thompson, and H. Yuksel, “A Survey About Nothing: Monitoring a MillionSupergiants for Failed Supernovae”, ApJ, 684, 1336, (2008).

35. J. A. Holtzman et al. (32 authors, incl. J. L. Prieto), “The Sloan Digi-tal Sky Survey-II: Photometry and Supernova Ia Light Curves from the 2005 Data”,AJ, 136, 2306, (2008).

36. C. J. Grier et al. (17 authors, incl. J. L. Prieto), “The Mass of theBlack Hole in the Quasar PG 2130+099”, ApJ, 688, 837, (2008).

37. J. L. Prieto, and O. Y. Gnedin, “Dynamical Evolution of Globular Clus-ters in Hierarchical Cosmology”, ApJ, 689, 919, (2008).

38. S. Blondin, J. L. Prieto, F. Patat, P. Challis, M. Hicken, R. P. Kirshner,T. Matheson, and M. Modjaz, “A Second Case of Variable Na I D Lines in a HighlyReddened Type Ia Supernova”, ApJ, 693, 207, (2009).

39. J. C. Bird, K. Z. Stanek, and J. L. Prieto, “Using Ultra Long Period

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Cepheids to Extend the Cosmic Distance Ladder to 100 Mpc and Beyond”, ApJ,695, 874, (2009).

40. H. E. Bond, L. R. Bedin, A. Z. Bonanos, R. M. Humphreys, L. A. G.Berto Monard, J. L. Prieto, and F. M. Walter, “The 2008 Luminous OpticalTransient in the Nearby Galaxy NGC 300”, ApJL, 695, 154, (2009).

41. M.J. Valtonen et al. (41 authors, incl. J. L. Prieto), “Tidally InducedOutbursts in OJ 287 during 2005-2008”, ApJ, 698, 781, (2009).

42. C. Badenes, J. Harris, D. Zaritsky, and J. L. Prieto, “The Stellar Ances-try of Supernovae in the Magellanic Clouds - I. the Most Recent Supernovae in theLarge Magellanic Cloud”, ApJ, 700, 727, (2009).

FIELDS OF STUDY

Major Field: Astronomy

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Table of Contents

Abstract . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . ii

Dedication . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . v

Acknowledgments . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . vi

Vita . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . xi

List of Tables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . xx

List of Figures . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . xxii

Chapter 1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1

1.1 Massive Stars and Supernovae . . . . . . . . . . . . . . . . . . . . . . 1

1.2 Probing Massive Star Evolution and Supernova Connection . . . . . . 3

1.2.1 Indirect Studies: Host Galaxy Environments . . . . . . . . . . 4

1.2.2 Direct Studies: Identifying Transient-Progenitor Pairs . . . . . 5

1.3 Scope of the Dissertation . . . . . . . . . . . . . . . . . . . . . . . . . 7

Chapter 2 Characterizing Supernova Progenitors via the Metallicitiesof their Host Galaxies, from Poor Dwarfs to Rich Spirals . . . . . 10

2.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10

2.2 First Catalog: Supernova-Host Pairs with Known Host Metallicities(SAI ∩ SDSS-DR4) . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

2.2.1 Testing Supernova Trends with Metallicity . . . . . . . . . . . 20

2.2.2 Supernovae in Low-Metallicity Hosts . . . . . . . . . . . . . . 26

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2.3 Second Catalog: Supernova-Host Pairs with Unknown HostMetallicities (SAI ∩ SDSS-DR6) . . . . . . . . . . . . . . . . . . . . . 28

2.4 Discussion and Conclusions . . . . . . . . . . . . . . . . . . . . . . . 30

Chapter 3 Discovery of the Dust-Enshrouded Progenitor of theType IIn SN 2008S with Spitzer . . . . . . . . . . . . . . . . . . . . . . 54

3.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54

3.2 Searching for the Progenitor . . . . . . . . . . . . . . . . . . . . . . . 56

3.3 Beneath the Shroud . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60


Chapter 4 A New Class of Luminous Transients and A First Censusof Their Massive Stellar Progenitors . . . . . . . . . . . . . . . . . . 70

4.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

4.2 The Class . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75

4.2.1 SN 2008S . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76

4.2.2 NGC 300 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 77

4.2.3 M85 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79

4.2.4 SN 1999bw . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82

4.2.5 The Connection to Other Transients . . . . . . . . . . . . . . 83

4.3 Rates . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86

4.3.1 Observed Counts . . . . . . . . . . . . . . . . . . . . . . . . . 87

4.3.2 Arguments for Incompleteness & Some Implications . . . . . . 90

4.4 A First Census . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95

4.4.1 Catalog . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

4.4.2 The Color-Magnitude Diagram . . . . . . . . . . . . . . . . . 99

4.4.3 Spectral Energy Distributions . . . . . . . . . . . . . . . . . . 104

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4.4.4 Variability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

4.4.5 Other Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . 109

4.5 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

4.5.1 Numbers & Rates . . . . . . . . . . . . . . . . . . . . . . . . . 111

4.5.2 Connection to The Evolution of Massive Stars . . . . . . . . . 116

4.5.3 A More Complete Census . . . . . . . . . . . . . . . . . . . . 124

Chapter 5 A Spitzer/IRS Spectrum of the 2008 Luminous Transient inNGC 300: Connection to Proto-Planetary Nebula . . . . . . . . . . 137

5.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137

5.2 Spitzer Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . 139

5.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 140

5.3.1 Spectral Features . . . . . . . . . . . . . . . . . . . . . . . . . 140

5.3.2 Spectral Energy Distribution . . . . . . . . . . . . . . . . . . . 144

5.4 Discussion & Conclusions . . . . . . . . . . . . . . . . . . . . . . . . 149

5.4.1 Mid-IR Spectrum and SED of NGC 300-OT . . . . . . . . . . 149

5.4.2 NGC 300-OT and SN 2008S: Connection to Proto-PlanetaryNebulae . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 152

5.4.3 The Progenitors of NGC 300-OT and SN 2008S: MassiveCarbon-rich AGB/post-AGB stars ? . . . . . . . . . . . . . . 155

5.4.4 Progenitors and Transients: Concluding Remarks . . . . . . . 158

Chapter 6 A Study of the Type Ia/IIn Supernova 2005gj from X-rayto the Infrared . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 169

6.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 169

6.2 Photometry . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 173

6.2.1 SDSS and MDM . . . . . . . . . . . . . . . . . . . . . . . . . 173

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6.2.2 CSP . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 175

6.3 Spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 177

6.4 X-ray Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 179

6.5 Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 182

6.5.1 Optical light curves and colors . . . . . . . . . . . . . . . . . . 182

6.5.2 NIR light curves . . . . . . . . . . . . . . . . . . . . . . . . . 189

6.5.3 Bolometric light curve . . . . . . . . . . . . . . . . . . . . . . 190

6.5.4 Optical spectroscopy . . . . . . . . . . . . . . . . . . . . . . . 194

6.6 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 201

6.6.1 Structure of the CSM . . . . . . . . . . . . . . . . . . . . . . . 204

6.6.2 Rates, hosts galaxies and possible progenitors of SN 2002ic-likesupernovae . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 212

Chapter 7 LBT Discovery of a Yellow Supergiant Eclipsing Binary inthe Dwarf Galaxy Holmberg IX . . . . . . . . . . . . . . . . . . . . . 251

7.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 251

7.2 Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 253

7.3 Light Curve . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 254


Appendix A A. Extreme-AGB Star Variability . . . . . . . . . . . . . 267

Appendix B LBV Candidate Variability . . . . . . . . . . . . . . . . . . 270

Bibliography . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 270

xix

List of Tables

2.1 Supernova and host galaxy data . . . . . . . . . . . . . . . . . . . . 41













3.1 Spectral Energy Distribution of the Progenitor of SN 2008S . . . . . 69

4.1 MIR Catalog for 53,194 Point Sources in M33 . . . . . . . . . . . . . 134

4.2 Photometry for the 18 EAGBs in M33 . . . . . . . . . . . . . . . . . 135

4.3 Photometry for the 16 LBV Candidates in M33 from Massey etal. (2007) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 136

5.1 Features in the Spitzer/IRS Spectrum of NGC300-OT . . . . . . . . 166

5.2 Spectral Energy Distribution of the Progenitor of NGC 300-OT . . . 167

xx

5.3 Black-body Fits to the Transient and Progenitor SEDs . . . . . . . . 168

6.1 SDSS ugriz and CSP u′g′r′i′ photometry of comparison stars incommon in the field of SN 2005gj. . . . . . . . . . . . . . . . . . . . 235

6.1 SDSS ugriz and CSP u′g′r′i′ photometry of comparison stars incommon in the field of SN 2005gj. . . . . . . . . . . . . . . . . . . . 236

6.2 SDSS and MDM ugriz photometry of SN 2005gj . . . . . . . . . . . . 237

6.2 SDSS and MDM ugriz photometry of SN 2005gj . . . . . . . . . . . . 238

6.3 CSP u′g′r′i′ photometry of SN 2005gj . . . . . . . . . . . . . . . . . . 239

6.4 CSP Y JHKs photometry of SN 2005gj . . . . . . . . . . . . . . . . 240

6.5 Light-curve parameters for SN 2005gj . . . . . . . . . . . . . . . . . . 241

6.6 Spectroscopic observations of SN 2005gj . . . . . . . . . . . . . . . . 242

6.6 Spectroscopic observations of SN 2005gj . . . . . . . . . . . . . . . . 243

6.7 K-corrections of SN 2005gj . . . . . . . . . . . . . . . . . . . . . . . 244

6.8 Derived integrated luminosity and black-body fits. . . . . . . . . . . 245

6.8 Derived integrated luminosity and black-body fits. . . . . . . . . . . 246

6.9 Library of spectra used in SNID . . . . . . . . . . . . . . . . . . . . 247

6.9 Library of spectra used in SNID . . . . . . . . . . . . . . . . . . . . 248

6.10 Results of the Gaussian fits to Hα and Hβ features . . . . . . . . . . 249

6.10 Results of the Gaussian fits to Hα and Hβ features . . . . . . . . . . 250

7.1 Best-fit Binary Model Parameters. . . . . . . . . . . . . . . . . . . . 266

xxi

List of Figures

2.1 Metallicities of supernova host galaxies as a function of redshift andabsolute B magnitude from SDSS . . . . . . . . . . . . . . . . . . . . 36

2.2 Sample of metal rich and metal poor supernova host galaxies in SDSS 37

2.3 Cumulative fraction of oxygen abundances for supernova hosts in SDSS 38

2.4 Fraction of SN Ib/c to SN II as a function of host metallicity in SDSS 39

2.5 Cumulative fraction of projected separation between supernova and host 40

3.1 Pre-explosion images of SN 2008S from LBT and Spitzer . . . . . . . 66

3.2 Light curves of the progenitor of SN 2008S from Spitzer photometry . 67

3.3 Spectral energy distribution of the progenitor of SN 2008S . . . . . . 68

4.1 Mid-infrared color-magnitude diagram of M33 . . . . . . . . . . . . . 127

4.2 Same as Figure 4.1, but focused on the red and bright region of interest128

4.3 Images showing the reddest source in Figure 4.1 . . . . . . . . . . . . 129

4.4 Mid-infrared color-color diagram . . . . . . . . . . . . . . . . . . . . . 130

4.5 Spectral energy distributions of EAGB stars in M33 and candidate LBVs131

4.6 Light curves of two selected sources . . . . . . . . . . . . . . . . . . . 132

4.7 RMS variation of sources at 4.5 µm as a function of color . . . . . . . 133

5.1 Spitzer mid-infrared spectrum of NGC 300-OT . . . . . . . . . . . . . 161

5.2 Comparison of NGC 300-OT to type IIP supernovae . . . . . . . . . . 162

5.3 Comparison of NGC 300-OT to massive stars . . . . . . . . . . . . . 163

5.4 Comparison of NGC 300-OT to proto-planetary nebula . . . . . . . . 164

xxii

5.5 Spectral energy distribution of NGC 300-OT . . . . . . . . . . . . . . 165

6.1 r-band image of the field of SN 2005gj . . . . . . . . . . . . . . . . . 219

6.2 Spectra of SN 2005gj . . . . . . . . . . . . . . . . . . . . . . . . . . . 220

6.3 More spectra of SN 2005gj . . . . . . . . . . . . . . . . . . . . . . . . 221

6.4 ugrizY JHKs light curves of SN 2005gj . . . . . . . . . . . . . . . . . 222

6.5 Difference between synthetic g − r color from spectra and photometry 223

6.6 Colors of SN 2005gj as a function of time . . . . . . . . . . . . . . . . 224

6.7 Light curve of SN 2005gj in the optical compared with others . . . . . 225

6.8 Light curve of SN 2005gj in the near-IR . . . . . . . . . . . . . . . . . 226

6.9 Examples of black-body fits to the spectral energy distribution ofSN 2005gj . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 227

6.10 Bolometric light curve of SN 2005gj . . . . . . . . . . . . . . . . . . . 228

6.11 Comparison of spectra of SN 2005gj with other supernovae . . . . . . 229

6.12 Cross-correlation comparison of SN 2005gj with other spectra . . . . 230

6.13 Fits to the spectra of SN 2005gj . . . . . . . . . . . . . . . . . . . . . 231

6.14 Evolution of Balmer lines . . . . . . . . . . . . . . . . . . . . . . . . . 232

6.15 Line profiles of Hα and Hβ . . . . . . . . . . . . . . . . . . . . . . . . 233

6.16 Identification of lines in the spectrum of SN 2005gj . . . . . . . . . . 234

7.1 Light curve of eclipsing binary in Holmberg IX . . . . . . . . . . . . . 262

7.2 Color-magnitude diagrams of Holmberg IX with the binary . . . . . . 263

7.3 Light curve of eclipsing binary in the SMC . . . . . . . . . . . . . . . 264

7.4 Color-magnitude diagrams with the eclipsing binaries in Holmberg IXand SMC . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 265

A.1 Light curves of EAGB sources 1 . . . . . . . . . . . . . . . . . . . . . 268

A.2 Light curves of EAGB sources 2 . . . . . . . . . . . . . . . . . . . . . 269

xxiii

B.1 Light curves of LBVs 1 . . . . . . . . . . . . . . . . . . . . . . . . . . 271

B.2 Light curves of LBVs 2 . . . . . . . . . . . . . . . . . . . . . . . . . . 272

xxiv

Chapter 1

Introduction

1.1. Massive Stars and Supernovae

Stars are born, they live, evolve, and die. Initially composed mainly of

Hydrogen (∼ 70%) and Helium (∼ 25%), the life cycle of a star depends most

significantly on a single parameter: mass. Stars less massive than ∼ 8− 10 M⊙ burn

Hydrogen to Helium in the core during the main-sequence through thermonuclear

reactions, then Helium to Carbon/Oxygen (M ∼> 0.5 M⊙) in their post main-sequence

evolution. They end their lifes expelling their outer envelope in a planetary nebula

and leaving behind a Carbon/Oxygen white dwarf core of ∼< 1.4 M⊙. The lifetime of

a M ∼ 1 − 8 M⊙ star expands more than two orders of magnitude, from ∼ 40 Myr

(8 M⊙) to ∼ 10 Gyr (1 M⊙).

Stars more massive than ∼ 8 − 10 M⊙ burn past C/O in the core until they

form an Iron core of MChandra ≃ 1.4 M⊙. At this point the thermonuclear reactions

are no longer exothermic, neutrino losses dominate, and the core cannot support

itself against gravity. The core collapses in a fraction of a second, matter bounces

1

into the newly formed neutron star, and an energetic shock propagates out helped

by the production of ∼ 1053 erg in neutrinos. Eventually, the shock “breaks out”

through the envelope of the star and produces an energetic display of electromagnetic

radiation, that can last for months, marking the catastrophic death of a massive star

as a core-collapse supernova. Since supernovae have been discovered, observed, and

studied in our Galaxy and external galaxies, we know that this occurs in nature at

least in a fraction of massive stars, although theorists have mostly failed to produce

successful supernova explosions (e.g., Thompson et al. 2003). Stars more massive

than ∼ 25 M⊙ may end their life forming a black-hole instead of a neutron star

without an optically bright supernova.

Mass is not the only physical parameter that regulates the evolution of a star.

Stellar evolution theory (e.g., Meynet & Maeder 2003) predicts that metallicity,

rotation, and binarity are also important and can significantly affect the lives and

deaths of massive stars (e.g., Heger et al. 2003; Eldridge et al. 2008). In general

terms, metallicity is a source of opacity that affects the transport of photons through

the star and enhances radiatively-driven winds; rotation is a source of angular

momentum that changes the energetics and structure of the star; and the presence

of a close binary can change the evolution of the star through interactions and

mass-transfer.

Mass loss, which in turn is a complicated function of temperature, luminosity,

mass, metallicity, rotation, and binarity, plays a crucial role shaping the life of

2

massive stars and their final fate as supernovae. We know this from studies of

stellar populations in the Galaxy and nearby galaxies (e.g., Massey 2003), as well as

from the large variety of different classes of core-collapse supernovae that have been

identified (e.g., Filippenko 1997). The observational classification of core-collapse

supernovae, mainly based on the presence or absence of H and He in the optical

spectra obtained close to maximum light – type II-P, II-L, IIb, Ib, Ic, IIn – is

primarily an imprint of the mass-loss history of the progenitor star.

1.2. Probing Massive Star Evolution and Supernova

Connection

During their life, massive stars pollute the interstellar medium (ISM) with

metals through strong winds and inject energy into the ISM that affects the evolution

of subsequent generations of stars and galaxies (e.g., Massey 2003). At the time of

their death as a supernova, they inject still more energy into the ISM (∼ 1051 erg of

kinetic energy) and produce a large fraction of the elements heavier than Iron (e.g.,

Woosley et al. 2002).

Since their importance for astrophysics is large, testing the evolution of massive

stars and their connection to supernovae observationally is extremely important.

This is, however, difficult because massive stars are rare in any given galaxy and

3

supernovae explosions are even rarer (∼ 1 per 100 years in a galaxy like the Milky

Way).

1.2.1. Indirect Studies: Host Galaxy Environments

Statistical studies of the resolved stellar populations of massive stars in the

Milky Way and nearby galaxies have helped to constrain stellar evolution models.

For example, the ratio of the number of O-type stars to Wolf-Rayet stars in galaxies

with different metallicities and at different locations within the galaxy can give

important clues about mass-loss (e.g., Maeder & Meynet 1994; Massey 2003).

In a similar way, studying the host galaxy environment of extragalactic

supernovae discovered by different surveys offers an excellent opportunity to connect

the average properties of the (unresolved) massive star population (like metallicity)

with stellar deaths.

This kind of statistical study needs a large number of supernovae discoveries

with accurate positions, and only in the last ∼ 10 years supernova surveys

undertaken by amateur and professional astronomers have been sucessfull at finding

large numbers of supernovae (∼ 3300) out to cosmological distances. However,

studies of host galaxies of core-collapse supernovae have been lacking mainly because

of the overwhelming interest of the community in type Ia supernovae (exploding

white dwarfs rather than collapsing massive stars) to measure cosmological distances,

4

with some important exceptions (e.g., van Dyk 1992; van Dyk et al. 1996; Prantzos

& Boissier 2003; James & Anderson 2006; Anderson & James 2008; Modjaz et

al. 2008).

The large sky area covered by the Sloan Digital Sky Survey (SDSS), with

homogeneous imaging and spectra of thousands of galaxies, provides a unique

opportunity to study the host galaxies of many supernovae that overlap with the

survey galaxies. Taking advantage of the results of Tremonti et al. (2004) on the

Oxygen abundances of star-forming galaxies derived from SDSS spectra, in this

dissertation we present a statistical comparison of the metallicities of different

supernova types. These results can be directly connected to stellar evolution models

for massive stars.

1.2.2. Direct Studies: Identifying Transient-Progenitor

Pairs

One of the most powerful tests of stellar evolution theory for massive stars is

to observationally establish the causal mapping between different populations of

evolved stars that are observed in galaxies (e.g., massive AGB stars, blue-, yellow-,

and red-supergiants, Wolf-Rayet stars, Luminous Blue Variables) and their final

explosive deaths.

5

This connection has been firmly proven only for a small number of objects

(< 10), most notably in the case of the core-collapse supernova SN 1987A in the

Large Magellanic Cloud (LMC). SN 1987A, which had an unusual light curve and

spectral properties (e.g., Menzies et al. 1987), had a blue-supergiant progenitor

star with an estimated main-sequence mass of ∼ 20 M⊙ (e.g., West et al. 1987).

The progenitor star had been identified and cataloged in pre-explosion images

(Sk-69 202). The blue-supergiant nature of the progenitor of SN 1987A came as a

big surprise because stars of that mass were expected to explode as red-supergiants

(e.g., Arnett 1987).

In the last ∼ 10 years, the combination of successful supernova surveys,

and the increasing volume of deep archival imaging data of nearby (∼< 30 Mpc)

galaxies obtained primarily with the Hubble Space Telescope (HST), has allowed

new identifications of supernova progenitors in pre-explosion images (e.g., Smartt et

al. 2004; Li et al. 2007; Gal-Yam et al. 2007). Arguably, one of the most important

results has been the identification of red-supergiant stars with main-sequence masses

in the range 9 ∼< M (M⊙) ∼< 17 (e.g., Smartt et al. 2009) as the progenitors of

type II-P supernovae, the most common class of core-collapse supernovae (∼ 60% of

all).

However, the dearth of progenitors of type II-P supernovae with initial

main-sequence masses in the range ∼ 20− 25 M⊙ (“red supergiant problem”; Smartt

et al. 2009) was unexpected since red-supergiants with these masses are observed in

6

the stellar populations of our Galaxy and the Magellanic Clouds. This result might

indicate that more massive stars die quietly forming a black-hole (e.g., Kochanek

et al. 2008), that circumstellar dust is important at these masses hiding them from

optical observations, or that the most massive red-supergiants lead to supernova

types different from type II-P (e.g., Smith et al. 2009a).

In the main part of this dissertation, we present the discovery of the progenitors

of two supernova-like transients found in 2008 in the nearby galaxies NGC 6946

and NGC 300. The progenitors of these transients are unusual with respect to the

progenitors of other supernovae found before. They were undetected in deep optical

images, but bright in the mid-infrared, indicating the presence of circumstellar dust

that absorbs the optical-UV light of the central star and re-radiates it thermally

at mid-infrared wavelengths. Furthermore, the luminosities of the progenitors of

≃ 5 × 104 L⊙ puts them in the range of main-sequence masses of ∼ 8 − 10 M⊙, the

important transition region between low-mass and massive stars.

1.3. Scope of the Dissertation

The outline of this dissertation is as follows. In Chapter 2, we present a study

of the metallicities of supernova host galaxies in the SDSS galaxy catalog and

the implications of these findings for supernova progenitors. Chapter 2 has been

7

published in The Astrophysical Journal as J. L. Prieto et al. 2008, ApJ, v. 673,

p. 999.

In Chapter 3, we present the discovery of the progenitor star of the low-

luminosity SN 2008S in pre-explosion images obtained with the Spitzer Space

Telescope at mid-infrared wavelengths. Chapter 3 has been published in The

Astrophysical Journal Letters as J. L. Prieto et al. 2008, ApJL, v. 681, p. 9.

In Chapter 4, we present a comprehensive study of the new type of stellar

transients, uncovered by SN 2008S and the transient in NGC 300, and their dusty

progenitor stars. The rates of these transients and the mid-infrared properties of

their progenitors compared to massive stellar populations in nearby galaxies allows

us to put interesting constraints on the physical mechanism of the explosion and

their importance for stellar evolution. We have submitted a paper (T. A. Thompson,

J. L. Prieto, et al. 2009) to The Astrophysical Journal with the analysis and results

presented in Chapter 4.

In Chapter 5, we present a mid-infrared spectrum of the transient in NGC 300

obtained with the Spitzer Space Telescope and the analysis of the spectral energy

distribution that reveal a substantial fraction of circumstellar dust that survives the

explosion. The properties of this transient are consistent with the results presented

in Chapter 4 on the dust-enshrouded progenitor star. Chapter 5 has been submitted

for publication to The Astrophysical Journal as J. L. Prieto et al. 2009.

8

In Chapter 6, we present data and analysis of the light curves and spectra of

the unusual supernova 2005gj, discovered by the SDSS-II Supernova Survey. We

argue that this supernova was most likely the result of a thermonuclear explosion

in a white dwarf with a relatively massive AGB star companion or a single massive

AGB star. The data and results of this chapter have been submitted for publication

in The Astronomical Journal as J. L. Prieto et al. 2007.

Finally, in Chapter 7, we present the discovery of an eclipsing binary composed

of two massive yellow-supergiant stars in the nearby galaxy Holmberg IX. The data

for this work were obtained with the Large Binocular Telescope as part of a large

project to study massive stars and supernova progenitors in a complete sample of

nearby galaxies. Chapter 7 has been published in The Astrophysical Journal Letters

as J. L. Prieto et al. 2008, ApJL, v. 673, p. 59.

9

Chapter 2

Characterizing Supernova Progenitors via the

Metallicities of their Host Galaxies, from

Poor Dwarfs to Rich Spirals

2.1. Introduction

On general grounds, it is thought that metallicity will affect the endpoints of

stellar evolution, e.g., the relative outcomes in terms of different supernova types and

the observed properties of each. Metals are a source of opacity that affects supernova

progenitors (e.g., Kudritzki & Puls 2000) and also the supernova explosions

themselves (e.g., Heger et al. 2003). However, the hypothesized metallicity effects

have been rather difficult to measure directly. The number of supernova progenitors

that have been identified directly from pre-explosion imaging is small and limited to

core-collapse events (e.g., Hendry et al. 2006; Li et al. 2007). Previous works have

either used population studies with only observational proxies for metallicity (e.g.,

Prantzos & Boissier 2003) or have considered direct metallicity measurements with

only relatively small numbers of events (e.g., Hamuy et al. 2000; Gallagher et al.

2005; Stanek et al. 2006; Modjaz et al. 2008).

10

A new approach is now possible, which we employ in this paper, that takes

advantage of the large sample of well-observed and typed supernovae. Due to a

fortuitous match in coverage, many of these supernovae were in galaxies for which

the Sloan Digital Sky Survey (SDSS) has identified the host galaxies and measured

their oxygen abundances from emission lines in their spectra (Tremonti et al. 2004).

While these are central metallicities for the host galaxies, and are not measured for

each supernova site, they are much more directly connected to the latter than proxies

like the host luminosity. To further sharpen our tests, we compare the metallicity

distributions of the host galaxies of SN Ib/c and SN Ia to those of SN II, which are

taken as a control sample.

The progenitors of core-collapse supernovae (SN II and Ib/c) are massive stars,

either single or in binaries, with initial main sequence masses ∼> 8 M⊙ (e.g., Heger

et al. 2003). The presence of hydrogen in the spectra of SN II indicates that the

massive envelopes are retained by the progenitors, of which red supergiants are

probably the most common. However, SN Ib/c lack hydrogen (SN Ib) or both

hydrogen and helium (SN Ic) in their spectra, and are therefore thought to have

Wolf-Rayet (WR) stars as progenitors (see Crowther 2007 for a review). The latter

originate from the most massive stars, and have had their outer layers stripped off

by strong winds. Thus SN Ib/c are thought to have main sequence masses ∼> 30 M⊙,

which would make them ≃ (8/30)1.35 ≃ 20% of all core-collapse supernovae, assuming

a Salpeter slope in the high-mass end of the initial mass function.

11

Based on theoretical considerations, the effects of line-driven winds are expected

to introduce a metallicity dependence in the minimum mass necessary to produce

WR stars (e.g., Heger et al. 2003; Eldridge & Tout 2004; Vink & de Koter 2005),

which in turn can change the fractions of core-collapse supernovae that explode

as SN II and SN Ib/c. Due to the relative frequencies, SN Ib/c will be more

affected than SN II. These metallicity effects on the progenitor winds may strongly

affect the rate at which radioactive 26Al is expelled into the interstellar medium

before decaying (e.g., Prantzos 2004; Palacios et al. 2005), in which case the decays

contribute to the observed diffuse 1.809 MeV gamma-ray line emission from the

Milky Way (e.g., Diehl et al. 2006). While 26Al appears to originate in massive

stars, it is not yet known how much comes from the progenitors or the different

core-collapse supernova types (e.g., Prantzos & Diehl 1996; Higdon et al. 2004). For

the most massive stars, GRB progenitors in the collapsar model (e.g., MacFadyen &

Woosley 1999, Yoon & Langer 2005), the interplay between metallicity-dependent

mass loss through winds and rotation may be crucial (e.g., Hirschi et al. 2005). In

all cases, binary progenitors may be more complicated (e.g., Eldridge 2007b).

Prantzos & Boissier (2003) used the absolute magnitudes of galaxies as a proxy

for their average metallicities, from the luminosity-metallicity relationship, and

found that the number ratio of SN Ib/c to SN II increases with metallicity; they

argued that their result is consistent with stellar evolution models of massive stars

with rotation (e.g., Meynet et al. 2006). If so, then one would expect a more robust

12

signature if the host metallicities were known directly. Ideally, in the latter approach,

one would use the metallicities as measured from follow-up spectra obtained at the

supernova sites, but this is difficult in practice. This approach of using measured as

opposed to estimated metallicities was used by Stanek et al. (2006) (with compiled

results from the literature) to study nearby long-duration GRBs with subsequent

supernovae, finding that all of them had very low metallicity environments and that

this appeared to be key to forming powerful GRB jets, and by Modjaz et al. (2008)

to study nearby broad-lined SN Ic (without GRBs), finding in contrast that the

metallicities of these environments were much higher. The main caveats associated

with these results are the low statistics, five and twelve events, respectively. We try

to combine the virtues of these two approaches, with higher statistics and mostly

direct metallicity measurements.

The likely progenitors of SN Ia are white dwarfs, forming from stars with

initial main-sequence masses ∼< 8 M⊙, which accrete mass from a companion

(single-degenerate model) until they reach the Chandrasekhar mass (≃ 1.4 M⊙) and

produce a thermonuclear explosion that completely disrupts the star (e.g., Whelan

& Iben 1973). During the accretion process, white dwarfs could have strong winds

when the accretion rate reaches a critical value (e.g., Hachisu et al. 1996), which

would allow them to burn hydrogen steadily and grow in mass. At low metallicities

([Fe/H] ∼< −1), SN Ia may be inhibited through the single-degenerate channel

(Kobayashi et al. 1998), as the white dwarf wind is thought to be weak and the

13

system passes through a common envelope phase before reaching the Chandrasekhar

mass. Metallicity also affects the CNO abundances of white dwarfs, which can

affect the production of 56Ni in the explosion, and therefore the peak luminosities

of SN Ia (e.g., Umeda et al. 1999; Hoflich et al. 2000; Timmes et al. 2003; Ropke

et al. 2006). Studies of the integrated metallicities of nearby SN Ia hosts (Hamuy

et al. 2000; Gallagher et al. 2005) have shown that metallicity does not seem to be

the main factor regulating their peak luminosities, which is consistent with some

theoretical models (e.g., Podsiadlowski et al. 2006). Instead, the age of the stellar

population where SN Ia progenitors originate seems to be very important: prompt

(SN Ia explode ∼ 108 yr after star formation) and delayed (SN Ia explode > 109 yr

after star formation) components were suggested to explain the high rates of SN Ia

in actively star-forming galaxies (late type spirals and irregulars) compared with

SN Ia in old, elliptical galaxies (e.g., Mannucci et al. 2005; Scannapieco & Bildsten

2005; Neill et al. 2006).

In this work, to our knowledge for the first time, we compare the directly

measured oxygen abundances of the hosts of SN Ib/c and SN Ia with SN II. We use

the Sternberg Astronomical Institute (SAI) supernova catalog and match it with the

SDSS-DR4 catalog of oxygen abundances of a large sample of star-forming galaxies

from SDSS. Using the supernova classifications presented in the literature, we can

separate the sample according to different supernova types and make statistical

comparisons of the metallicity distributions of their host galaxies. We also investigate

14

some individual cases in metal-poor environments that are especially interesting and

which can be used to test the strong predictions made by some theoretical models.

We create a second catalog by matching the positions of all supernovae with images

from SDSS-DR6, independent of the host galaxy association. This allows us to

investigate significantly fainter SNe hosts, and we identify some even more extreme

hosts for follow-up observations. To enable their further use in other studies, we

make both catalogs available online, and will update them regularly.

2.2. First Catalog: Supernova-Host Pairs with Known

Host Metallicities (SAI ∩ SDSS-DR4)

We use the SAI supernova catalog1 (Tsvetkov et al. 2004) to obtain the main

properties of supernovae (name, classification, RA, DEC, redshift) and their host

information when available (galaxy name, RA, DEC, redshift). The SAI catalog

is a compilation of information about supernova discoveries, obtained mainly from

reports in the International Astronomical Union Circulars (IAUC), which include

the coordinates and classification of the supernovae from the IAUCs and also basic

information about the host galaxies in the cases where the galaxies can be identified

in online galaxy catalogs (e.g., HyperLEDA, NED and SDSS). The version of the

catalog we use contains 4,169 entries2, of which we have selected 3,050 supernovae

1http://www.sai.msu.su/sn/sncat/

2Version updated on June 15, 2007.

15

discovered between 1909 and 2007 classified as SN Ia, II, and Ib/c, including their

sub-types. Supernovae in the catalog with no classification or only classified as

Type I are not considered for further analysis since we want to be able to distinguish

between SN Ia and the core-collapse types SN Ib/c.

Tremonti et al. (2004) determined metallicities for a sample of star-forming

galaxies in the SDSS Data Release 2 (SDSS-DR2; Abazajian et al. 2004) from their

spectra. Here we use a larger sample of 141,317 star-forming galaxies (excluding

AGN) from the SDSS-DR4 (Adelman-McCarthy et al. 2006), with metallicities

derived in the same consistent fashion, and which are available online3. The

metallicities are derived by a likelihood analysis which compares multiple nebular

emission lines ([O II], Hβ, [O III], Hα, [N II], [S II]) to the predictions of the hybrid

stellar-population plus photoionization models of Charlot & Longhetti (2001). A

particular combination of nebular emission line ratios arises from a model galaxy that

is characterized by a galaxy-averaged metallicity, ionization parameter, dust-to-metal

ratio, and 5500A dust attenuation. For each galaxy, a likelihood distribution for

metallicity is constructed by comparison to a large library of model galaxies. We use

the median of the oxygen abundance distributions in this paper. The metallicities

derived by Tremonti et al. (2004) are essentially on the Kewley & Dopita (2002)

abundance scale (∆[12 + log (O/H)] < 0.05 dex; Ellison 2006). For further reference

in this paper, we call this galaxy metallicity catalog SDSS-DR4.

3http://www.mpa-garching.mpg.de/SDSS/DR4/

16

We restrict the initial sample of galaxies to 125,958 by applying two of the cuts

that Tremonti et al. (2004) used for their final cleaned sample: (1) the redshifts

of the galaxies have to be reliable by SDSS standards; and (2) Hβ, Hα, and

[N II] λ6584 should be detected at > 5σ confidence, and [S II] λλ6717,6731 and

[O III] λ5007 should at least have detections. While in our analysis we directly

compare nebular oxygen abundance within the SDSS-DR4 catalog for the supernova

hosts, when referring to “Solar metallicity,” we adopt the Solar oxygen abundance

of 12 + log (O/H) = 8.86 (Delahaye & Pinsonneault 2006).

We cross-matched the SAI catalog with the galaxy metallicity catalog

SDSS-DR4 using a matching radius of 60′′ (∼ 48 kpc at z = 0.04). We used the

coordinates of the host galaxies in the cases where they are known and identified

in the SAI catalog, and the supernovae coordinates were used otherwise. We also

required that the redshifts reported in the SAI catalog, which were taken from

galaxy catalogs and the IAUCs, to be consistent within 20% with the redshifts of the

closest galaxy from the SDSS catalog that passed the proximity cut. After selecting

the supernovae that passed the proximity and redshift criteria, we visually inspected

the SDSS images around the galaxies to identify the ones that were wrongly selected

as hosts (e.g., close galaxy pairs). The number of supernovae that passed all these

cuts is 254 in total: 95 SN Ia, 123 SN II, and 36 SN Ib/c. There were some galaxies

that hosted more than one supernova: five galaxies had three supernovae each

(NGC 1084, NGC 3627, NGC 3631, NGC 3938, and NGC 5457) and 15 galaxies

17

had two supernovae (NGC 2532, NGC 2608, NGC 3627, NGC 3780, NGC 3811,

NGC 3913, NGC 4012, NGC 4568, NGC 5584, NGC 5630, NGC 6962, UGC 4132,

UGC 5695, IC 4229, and MCG +07-34-134).

In Table 2.1 we present the final matched sample of supernovae and host galaxy

metallicities from SDSS-DR4, as well as the absolute MB magnitudes of the galaxies

obtained from the HyperLEDA database and SDSS. The absolute magnitudes

for SDSS galaxies were calculated using Petrosian gr magnitudes transformed to

B magnitudes using the transformation of Lupton (2005), corrected by Galactic

extinction (Schlegel et al. 1998) and internal extinction to a face-on geometry (Tully

et al. 1998), and k-corrections (Blanton et al. 2003). To calculate the absolute

magnitudes, we use a flat cosmology with H0 = 70 km s−1 Mpc−1, ΩM = 0.3,

ΩΛ = 0.7. The typical 1σ uncertainties in the oxygen abundances are 0.05 dex

at 12 + log (O/H) > 8.5, and 0.15 dex at 12 + log (O/H) < 8.5. Our estimated

uncertainty in the absolute magnitudes of the hosts is ∼ 0.3 mag, calculated from a

sub-sample of galaxies in the catalog with reliable absolute magnitudes from SDSS

and HyperLEDA.

Our first catalog, SAI ∩ SDSS-DR4, is available online4 and will be updated

as new supernovae are discovered with host galaxy metallicities in the SDSS-DR4

catalog. It includes the information presented in Table 2.1, as well as images around

the supernovae obtained from SDSS-DR6.

4http://www.astronomy.ohio-state.edu/∼prieto/snhosts/

18

Figure 2.1 shows the distribution of metallicities as a function of redshift

and MB of the supernova host galaxies, as well as the distribution of star-forming

galaxies in the SDSS-DR4 catalog. The apparent “stripes” in the plots, regions

with very few oxygen abundance measurements, are an effect of the grid of model

parameters (metallicity, ionization parameters, attenuation, etc.) used to calculate

the metallicities (see Brinchmann et al. 2004 for details). As can be seen, the redshift

distribution of supernovae varies for different types, with the median redshifts of the

samples at z = 0.014 (II), 0.018 (Ib/c), and 0.031 (Ia). This is a combination of

several effects. First, SN Ia supernovae are, on average, ∼ 2 mag brighter at peak

luminosity than core-collapse events (Richardson et al. 2002), therefore, they can

be found at larger distances in magnitude-limited surveys. Second, the local rate of

core-collapse supernovae in late-type galaxies is ∼ 3 times higher than the SN Ia rate

(Cappellaro et al. 1999; Mannucci et al. 2005). Finally, the great interest in SN Ia as

cosmological distance indicators makes most of the supernovae searches concentrate

their limited spectroscopic follow-up resources on likely Type Ia supernovae (as

determined by their light curves).

As shown in Figure 2.1, the distribution of host galaxy metallicities follows

the distribution of galaxies from SDSS, with a wide range spanning ∼ 1.4 dex

(7.9 < 12 + log (O/H) < 9.3). However, there appear to be significant differences

present between the hosts of different supernovae types. In particular, most of

the SN Ib/c hosts are concentrated in the higher metallicity/luminosity end of the

19

distribution (12 + log (O/H) ∼> 8.7), while the metallicities of SN II and SN Ia hosts

are more evenly distributed and appear to be tracing each other fairly well.

Figure 2.2 shows a mosaic of SDSS-DR6 (Adelman-McCarthy et al. 2008)

images5 of the host galaxies with the highest and lowest metallicities in the sample,

including two supernovae of each type. This figure shows the wide range of host

galaxy environments present in the sample, from big spirals (e.g., SN 2000dq,

SN 2004cc, SN 2005bc, SN 2005mb, SN 2002cg, and SN 2006iq) to small dwarfs (e.g.,

SN 1997bo, SN 2006jc, SN 2004hw, SN 1998bv, SN 2007I, and SN 2007bk), and

that all types of supernovae can be found in metal-rich and metal-poor star-forming

galaxies.

2.2.1. Testing Supernova Trends with Metallicity

Is the tendency of SN Ib/c hosts towards higher metallicity, compared with

SN II and SN Ia, clearly seen in Figure 2.1, a real physical effect? To answer this

question we identify and try to reduce some of the biases present in the sample.

The supernova sample studied in this work is far from homogeneous. The

supernovae have been discovered by a variety of supernova surveys, including

amateur searches that look repeatedly at bright galaxies in the local universe,

professional searches that look at a number of cataloged galaxies to a certain

5http://cas.sdss.org/dr6/en/tools/chart/chart.asp

20

magnitude limit (e.g., LOSS), and professional searches that look at all the galaxies

in a given volume (e.g., SDSS-II, The Supernova Factory), among others. The

host galaxies of supernovae discovered by amateur searches tend to have higher

metallicities due to the luminosity-metallicity relation (see Figure 2.1), while the

metallicities of galaxies observed by professional searches span a wider range.

As an example of a possible bias in the supernovae in our catalog, we note that

the median metallicity decreases by ∼ 0.1 dex for the hosts of supernovae discovered

between 1970 and 2007. Ideally, all the supernovae for the current study would be

selected from galaxy-impartial surveys. However, the numbers of different supernova

types found by such surveys in our catalog are still small (especially core-collapse

events), and do not allow a statistical comparison (see the discussion in Modjaz

et al. 2008).

Another bias present in the galaxy data that we use is the so-called aperture

bias (Kewley et al. 2005; Ellison 2006). The SDSS spectra are taken with a fixed

fiber aperture of 3′′ (2.4 kpc at z = 0.04). Since galaxies have radial metallicity

gradients (e.g., Zaritsky et al. 1994), for nearby galaxies we are, on average, only

measuring the higher central metallicity, while for more distant galaxies we are

covering a larger fraction of the galaxy light with the SDSS fiber. This effect also

depends on galaxy luminosity, as for dwarf galaxies the fiber will cover a larger

fraction of the total light than in large spirals. Kewley et al. (2005) find a mean

difference of ∼ 0.1 dex, although with a large scatter (0.1 − 0.3 dex), between the

21

central and integrated metallicities of a sample of ∼ 100 galaxies of all types (S0-Im)

in the redshift range z = 0.005 − 0.014.

In order to reduce these and other biases, we limit the comparison of supernova

types to host galaxies in the redshift range 0.01 < z < 0.04, where there are 115

supernovae. In this pseudo volume-limited sample, the median redshifts of the 58

SN II (0.020), 19 SN Ib/c (0.021) and 38 SN Ia (0.024) hosts are consistent, while the

number of galaxies in each sub-sample still allows us to make a meaningful statistical

comparison. By using a small redshift slice we are, effectively, reducing the aperture

biases when comparing the galaxy metallicity measurements, such that they are now

comparable to or smaller than the statistical errors in the metallicity determination.

We made additional checks of relative biases between supernova types in our

redshift-limited sample. First, the ratios of the numbers of SN Ib/c and SN Ia to the

total number of core-collapse supernovae are reasonably consistent with the ratios

obtained from the local supernovae rates (e.g., Cappellaro et al. 1999; Mannucci

et al. 2005). Second, the fact that the SN–host separation distributions for SN Ia

and SN II agree, particularly at small radii (see below), suggests that our supernova

samples are not biased (relatively, one supernova type to another) by obscuration

effects.

We compare the metallicity distributions of the hosts of SN Ib/c and SN Ia to

SN II, which are taken as the control sample. Given that SN II are the most common

22

type of supernovae (e.g., Mannucci et al. 2005) and that they come from massive

stars from a wide range of masses that explode in all environments, presumably

independent of metallicity, they are effectively giving us the star-formation-rate

weighting of the luminosity-metallicity (or mass-metallicity) relationship for

star-forming galaxies. It would be of interest to test if indeed the SN II rates are

independent of metallicity, but this is outside the scope of the current paper.

Figure 2.3 shows the cumulative distribution of metallicities for hosts of

different supernova types in the redshift ranges z < 0.04 and 0.01 < z < 0.04. Two

important results are immediately apparent:

• The metallicities of SN Ib/c hosts tend to be higher than those of SN II hosts,

• The SN Ia and SN II hosts have very similar metallicity distributions.

Kolmogorov-Smirnov (KS) tests between the metallicity distributions of

different supernova types in the redshift range 0.01 < z < 0.04 strengthen these

findings. The KS probabilities of two host metallicity samples being drawn from the

same distribution are: 5% (II–Ib/c), 3% (Ia–Ib/c) and 56% (II–Ia). We obtain a

similar result if we compare the mean metallicities of the host samples: 8.94±0.04

(SN II), 8.94±0.04 (SN Ia) and 9.06±0.04 (SN Ib/c), where the errors are the RMS

of similar-sized samples obtained using bootstrap resampling.

23

The metallicity distribution of the SDSS-DR4 star-forming galaxies in our

redshift range, weighted only by galaxy counts, is also shown in Figure 2.3. This

should not be used in any comparisons, as it does not take into account the weighting

with star formation rate or the supernova and galaxy selection criteria. We take all

of these into account by only making relative comparisons between supernova types.

If we restrict the sample of SN Ib/c to only SN Ic and broad-lined Ic in the same

redshift range, leaving out supernovae classified as SN Ib/c and SN Ib, the difference

in metallicity distributions of the hosts of SN II and SN Ic+hypernovae (13 SN)

becomes smaller, with a KS probability of 19%. If only the three supernova classified

as SN Ib (SN 2003I, SN 2005O, and SN 2005hl) in the pseudo volume-limited sample

are not considered, then the KS probability of SN II and SN Ic+hypernovae+SNIb/c

being drawn from the same sample is 10%.

In Figure 2.4 we show the number ratio of SN Ib/c to SN II as a function

of the metallicities of the host galaxies. This ratio is very important because the

rates of core-collapse SNe are expected to change as a function of the progenitor

mass and metallicity and, therefore, it can help to put constrains on massive stellar

evolution models (e.g., Eldridge 2007b). We have calculated the ratio in bins of

equal number of SN II+SN Ib/c, with 11 SNe per bin, to do a direct comparison

with the results of Prantzos & Boissier (2003). The small statistics compared with

Prantzos & Boissier (2003), who used the absolute magnitudes of the hosts as a

proxy for the average metallities through the luminosity-metallicity relationship,

24

is reflected in the large errors of the ratio. The large error bars do not allow us

to put constraints in progenitor models, however, the general trend observed in

the cumulative distribution (see Figure 2.3) is confirmed with the number counts:

SN Ib/c are more common at higher metallicities compared with SN II. Our results

are consistent with those of Prantzos & Boissier (2003).

Figure 2.5 shows the cumulative distributions of projected host-supernova

distances for the reduced sample of 115 SNe used to compare the host metallicities

(0.01 < z < 0.04). Clearly, the SN Ib/c in the sample are found more towards the

centers of their hosts when compared with SN II and SN Ia (e.g., van den Bergh 1997;

Tsvetkov et al. 2004; Petrosian et al. 2005), which have similar distributions to each

other (as also in Figure 2.3). A galactocentric concentration of SN Ib/c and their

progenitors may be important for the angular distributions of diffuse gamma-ray

line emission from the Milky Way. Besides the 1.809 MeV line from 26Al, the 0.511

MeV line from positron annihilation is poorly understood, in terms of its high flux

and very strong central concentration (e.g., Casse et al. 2004; Knodlseder et al.

2005; Beacom & Yuksel 2006). Since the SN Ib/c are found at small separation, the

central galaxy metallicities determined by the SDSS should be representative of the

local environments of the supernovae. Taking into account the existence of negative

metallicity gradients in increasing galactocentric radii, the local metallicities of the

SN II and SN Ia, if anything, are even lower than deduced from the SDSS central

25

metallicities. The tendency for SN Ib/c to prefer higher metallicity relative to SN II

and SN Ia is probably even stronger than shown in Figure 2.3.

2.2.2. Supernovae in Low-Metallicity Hosts

Even though we have shown that there is a strong preference of SN Ib/c for

high-metallicity environments, compared with SN II and SN Ia, there are four

SN Ib/c with relatively metal-poor host galaxies (12 + log (O/H) < 8.6). These

events, and also some SN Ia found in low-metallicity dwarfs, made us investigate more

carefully a number of individual cases. We found that among the lowest-metallicity

host galaxies in the sample, there were supernovae that stood out because of their

unusual properties (all shown in Figure 2.2).

SN 2006jc: Peculiar SN Ib/c supernova with strong He I lines in emission in

the spectrum (Crotts et al. 2006), thought to arise from the interaction

of the supernova ejecta with a He-rich circumstellar medium (Foley et al.

2007; Pastorello et al. 2007). Its host galaxy, UGC 4904 at z = 0.006, is a

low luminosity, blue, and relatively low-metallicity starburst (MB = −16.0,

12 + log (O/H) = 8.5). Interestingly, the host galaxy of SN 2002ao at z = 0.005

(in UGC 9299), another peculiar SN Ib/c with spectral properties very similar

to SN 2006jc (Benetti et al. 2006a) that is also present in our first catalog, has

26

low metallicity compared with the majority of the SN Ib/c hosts, and shares

similar morphological properties with the host of SN 2006jc.

SN 2007I: Broad-lined Ic (or hypernova) with a spectrum similar to SN 1997ef

(Blondin et al. 2007a) at z = 0.022. Its host galaxy is a star-forming,

low-metallicity dwarf (MB = −16.7, 12 + log (O/H) = 8.3), unlike other

broad-lined Ic supernovae observed in higher-metallicity galaxies (Modjaz et al.

2008), and somewhat similar to the host galaxies of long GRBs associated with

supernovae (Stanek et al. 2006; Fruchter et al. 2006); however, see a detailed

discussion in Modjaz et al. (2008). The other four broad-lined Ic supernovae in

our sample that have been reported in the literature are: SN 1997ef (Iwamoto

et al. 2000), 2004ib (Sako et al. 2005), 2005ks (Frieman et al. 2005), 2006qk

(Frieman et al. 2006).

SN 2007bk: Type Ia supernova with a spectrum similar to the slow

decliner/luminous SN 1991T (Dennefeld et al. 2007) at z = 0.032. The

host galaxy is a low metallicity/luminosity dwarf, with MB = −18.2 and

12 + log (O/H) = 8.3, similar to the Large Magellanic Cloud. The supernova

was found very far from the center of its dwarf host, at a projected separation

of ∼ 9 kpc. The magnitude of the supernova at discovery (R = 16.7, Mikuz

2007) and the phase obtained from the spectrum (+50 days, Dennefeld

et al. 2007, although S. Blondin finds equally good matches with templates

at +30 days, private communication), imply that this was a very luminous

27

Type Ia event. If the reported discovery magnitude and spectral phases are

accurate, SN 2007bk was ∼ 0.5− 1.5 mag brighter than SN 1991T at the same

phase after maximum light.

2.3. Second Catalog: Supernova-Host Pairs with

Unknown Host Metallicities (SAI ∩ SDSS-DR6)

The existence of supernovae with unusual properties among the most metal-

poor, low-luminosity galaxies in the first catalog prompted us to investigate a much

larger sample of supernovae. We constructed a second catalog with images around

the positions of supernovae using SDSS, matching the SAI catalog with SDSS-DR6.

We included the redshifts obtained from the SAI catalog to produce images in

physical units around the supernovae. The total number of matches is 1225 for

supernovae at z < 0.3. This catalog is also available online, with the first catalog

described earlier in §2.2.

This extended second catalog (SAI ∩ SDSS-DR6) does not have information

about metallicities or luminosities of the hosts. It is a visual tool that can be used to

explore the environments around supernovae found in the SDSS area, independent

of the host galaxy association. Identification of the supernovae hosts and their

integrated properties obtained from SDSS will be added in a future study.

28

Visually inspecting the images of the second catalog, we identified a number

of supernovae in what appear to be very faint galaxies, and which are likely to be

low-luminosity, metal-poor galaxies not present in the first catalog. Some examples

in the redshift range 0.01 ∼< z ∼< 0.05 are (supernova types shown in parentheses):

SN 1997ab (II), SN 1997az (II), SN 2001bk (II), SN 2003cv (II), SN 2004gy (II),

SN 2004hx (II), SN 2005cg (Ia), SN 2005gi (II), SN 2005hm (Ib), SN 2005lb

(II), SN 2006L (IIn), SN 2006bx (II), SN 2006fg (II), 2007bg (Ic), 2007bu (II),

and 2007ce (Ic). In this incomplete sample, which was selected by noting some

especially low-luminosity galaxies, the SN Ia/SN II ratio is lower than expected.

Similarly, in our catalog of hosts with known metallicities, SN Ia may also be

relatively underabundant at the lowest host luminosities and metallicities, as shown

in Figure 2.1. We caution that the small statistics make these only hints, and we

discuss these issues further below.

One of the most interesting supernovae in our second catalog is SN 2007bg, a

recently discovered broad-lined Ic (Quimby et al. 2007; Harutyunyan et al. 2007)

at z = 0.03, which has an extremely faint galaxy coincident with the position of

the supernova. Using photometry and images from SDSS-DR6, we estimate the

luminosity of the apparent host galaxy to be MB ≈ −12, most likely a very metal-

poor dwarf (12 + log (O/H) ∼ 7.5, or ∼ 1/20 solar; see the metallicity-luminosity

relationship extended to dwarf galaxies by Lee et al. 2006). Due to the extremely

low luminosity of that galaxy, in fact one of the lowest-luminosity supernova hosts

29

ever seen, and also fainter than most if not all GRB hosts (see e.g., Fruchter et al.

2006), this event may represent the missing link between broad-lined SN Ic and

GRBs. This event is therefore an excellent candidate for a search for an off-axis

GRB jet in radio (Soderberg et al. 2006) and possibly other wavelengths. Follow-up

spectroscopic observations and deep photometry to determine the metallicity of the

host and study the supernova environment are strongly encouraged in this and other

cases of very low-luminosity SN hosts.

2.4. Discussion and Conclusions

We find that SN Ib/c tend to be in high-metallicity host galaxies, compared to

SN II, our control sample that traces the underlying star formation rates. This is

the first time that such a trend has been found using the directly-measured oxygen

abundances of the supernova host galaxies. This confirms and greatly strengthens

an earlier result of Prantzos & Boissier (2003), which found a similar result using the

absolute magnitudes of the host galaxies as an indirect estimate of their metallicities

through the luminosity-metallicity relationship. This can be interpreted in relative

supernova rates: the ratio of SN Ib/c to SN II increases with increasing metallicity

and hence also cosmic age. We also find that SN Ib/c are consistently found towards

the centers of their hosts compared with SN II and SN Ia, which had been also found

in previous studies (e.g., van den Bergh 1997; Tsvetkov et al. 2004; Petrosian et al.

30

2005). This suggests that direct measurements of metallicities at the explosion sites,

as opposed to the central host metallicities used here, would reveal an even stronger

effect, due to the radial metallicity gradients observed in spiral galaxies. The local

metallicities of SN Ib/c would be less reduced with respect to the central metallicities

than SN II and SN Ia, which would widen the separation seen in Figure 2.3.

The tendency towards high metallicity of SN Ib/c environments compared

to those of SN II supports, in general terms, theoretical models of the effects of

metallicity in stellar evolution and the massive stars that are core-collapse supernova

progenitors (e.g., Heger et al. 2003; Meynet et al. 2006; Eldridge 2007b; Fryer et al.

2007). Also, models of stellar evolution that include rotation, from Meynet et al.

(2006), predict that at high metallicity Wolf-Rayet stars will earlier enter the WC

phase when they still are rich in helium, and that these stars would explode as

SN Ib. The fact that we do see both SN Ib and SN Ic in hosts at high metallicity

should not be taken as inconsistent with these models, mainly because the number of

supernovae is small and the sample has not been homogeneously selected. There is

an indication, although not statistically significant, that SN Ib may be more common

in higher metallicity environments than SN Ic and broad-lined SN Ic in our sample.

The agreement between the metallicity distributions of the hosts of SN II and

SN Ia shows that their hosts are sampling a wide range of properties of star-forming

galaxies, from the relatively metal-poor dwarfs to metal-rich grand design spirals.

Using models of white dwarf winds in the framework of single-degenerate progenitors

31

of SN Ia (Hachisu et al. 1996), Kobayashi et al. (1998) made a strong prediction

that SN Ia would not be found in low metallicity environments, such as dwarf

galaxies and the outskirts of spiral galaxies. However, we do observe SN Ia in metal

poor dwarfs (e.g., SN 2004hw, SN 2006oa, and SN 2007bk, with host metallicities

between ∼ 0.2 and 0.5 solar) and at large projected distances (> 10 kpc) from their

star-forming hosts (e.g., SN 1988Y, SN 1992B, SN 1993I, SN 2001bg, SN 2002gf,

SN 2004ia, SN 2004ig, SN 2005ms, SN 2006fi, and SN 2006gl). There are also

extreme cases that have been pointed out in previous studies, like the low-luminosity

dwarf (MB ≈ −12.2) host galaxy of the luminous and slow decliner SN 1999aw

(Strolger et al. 2002), which is most likely very metal-poor (12 + log (O/H) ∼ 7.5,

or ∼ 1/20 solar; see Lee et al. 2006). Also, SN 2005cg was found in a dwarf with

subsolar gas metallicity (Quimby et al. 2006).

We do not find a statistically significant low-metallicity threshold in the

metallicities of SN Ia compared with SN II hosts, as predicted from theory by

Kobayashi et al. (1998) for single-degenerate progenitors of SN Ia with winds.

However, there is a preference for finding more SN II in very faint galaxies compared

with SN Ia in our second catalog, which is suggestive of a luminosity or metallicity

threshold for the main channel that produces SN Ia. This will have to be explored

in the future with a larger sample that includes good luminosity information for the

hosts and actual metallicities measured from spectra. If no metallicity threshold

is found in larger samples, it means that the models and predictions of white

32

dwarf winds will have to be revisited. This would have implications for modeling

and understanding of galactic chemical evolution that include the effects of white

dwarf winds to shut down SN Ia at low metallicities (e.g., Kobayashi et al. 2007).

Interestingly, modeling the X-ray spectra of supernova remnants from probable

SN Ia explosions in our Galaxy, the LMC and M31, Badenes et al. (2007) did not

find the strong effects of white dwarf winds predicted from theory.

On the other hand, independent of the existence (or not) of a mechanism

that can shut down the production of SN Ia in low-metallicity environments, we

have noted examples of SN Ia that explode in low-metallicity dwarf galaxies, like

SN 2007bk. Also, supernova remnants from probable SN Ia have been identified in

the LMC (e.g., Hughes et al. 1995) and SMC (e.g., van der Heyden et al. 2004). Is this

SN Ia population dominated by a different kind of progenitors, like double-degenerate

mergers, compared to the main progenitor channel? Is the expected trend between

progenitor metallicities and peak-luminosity starting to appear as we extend the

sample to even lower metallicity hosts? It is suggestive that the small number of

SN Ia in low-metallicity hosts, like SN 2007bk, SN 2005cg and SN 1999aw, were all

luminous events compared with normal SN Ia. Also, the very luminous SN Ia events

that have spectral signatures of a strong ejecta-CSM interaction, like SN 2005gj, are

mostly associated with low-luminosity, and most likely low-metallicity, hosts (Prieto

et al. 2007). Is low metallicity necessary to produce this extreme class of SN Ia?

33

Detailed comparison studies of the observational properties of supernovae in these

extreme environments are encouraged.

In the course of this work, we have prepared two new catalogs that should be

useful for other studies. We used the SAI supernova catalog and the SDSS-DR4

sample of metallicities of star-forming galaxies from Tremonti et al. (2004) to produce

a catalog of supernovae hosts with metallicities derived in a consistent fashion.

From this first catalog, we found several interesting core-collapse (e.g., SN 2002ao,

SN 2006jc, and SN 2007I) and SN Ia events (e.g., SN 2007bk) in low-metallicity

galaxies. We constructed a second catalog by matching the SAI supernova catalog

with images obtained from SDSS-DR6. The second catalog does not contain

information on host metallicities, but it can be used to investigate the environments

of supernovae independent of the host association. In that second catalog, we

found several examples of core-collapse supernovae in faint galaxies. One of most

interesting cases is SN 2007bg, a broad-lined SN Ic in an extremely low-luminosity

and very likely low-metallicity host. These catalogs will allow researchers to select

interesting candidates for further follow-up observations. Also, as more homogeneous

light curve and spectroscopic data become available for supernovae in the first

catalog, this will allow us to test possible correlations between supernova properties

and the metallicities of their hosts, which may turn out to be crucial for improving

our understanding of the nature of different supernova explosions. Another possible

34

use of our catalogs is for systematically characterizing the morphologies of supernova

hosts.

We stress the great importance of galaxy-impartial surveys for finding and

studying the properties of all supernovae types. Some very interesting and potentially

informative supernovae have been found in very low-luminosity, low-metallicity

galaxies, hosts which are not included in supernova surveys based on catalogs

of normal galaxies. These unusual supernovae and hosts may help probe the

relationship between the SN Ib/c and SN II core-collapse supernova types, the

progenitors of SN Ia as well as the possible correlations between observed SN Ia

properties and host metallicities, the supernova-GRB connection (e.g., Stanek et al.

2003) and its possible metallicity dependence (e.g., Stanek et al. 2006; Modjaz et al.

2008), and also to test the consistency between the cosmic stellar birth and death

rates (e.g., Hopkins & Beacom 2006). As we pointed out in §2.2.1, presently the

comparison of host metallicities using supernovae discovered by galaxy-impartial

surveys is limited by their small numbers, especially for core-collapse events, since

SN Ia receive much more attention when decisions about spectroscopic follow-up are

made. This is also true for the study of their observational properties (e.g., light

curves and spectra). However, in order to better understand all types of cosmic

explosions and put further constraints on the predictions of stellar evolution theory,

a larger effort on other supernovae types is greatly needed.

35

Fig. 2.1.— Metallicities of supernova host galaxies from SDSS-DR4 as a functionof redshift and absolute B magnitude. The symbols distinguish different supernovatypes: SN II (triangles), SN Ib/c (squares) and SN Ia (circles). The dots in the left

panel are 125,958 star forming galaxies in SDSS-DR4 with reliable metallicity andredshift measurements. The dots in the right panel are a subsample of 86,914 starforming galaxies (z > 0.005) selected from the main SDSS-DR4 galaxy sample.

36

IaII Ib/c

SN2000dq (II)SDSSJ020820.18+011103.5 NGC4568

SN2004cc (Ic)NGC5698 SN2005bc (Ia)

SDSSJ213933.98+102900SN2006iq (Ia)

UGC10415 SN2002cg (Ic)SN2005mb (II)

NGC4963

SDSSJ121952.54+074352.3 SN1997bo (II)

UGC04904SN2006jc (Ib/c)

SDSSJ235117.28+011339.7SN2004hw (Ia)

SDSSJ152845.00+585200.1SN2007bk (Ia)

SDSSJ115913.13−013616.2SN2007I (Ic)

SDSSJ103825.73+474236.7 SN1998bv (II)

Fig. 2.2.— SDSS color images of the most metal rich (top six images) and most metalpoor (bottom six images) host galaxies in our sample. We show two galaxies of eachsupernova type. The images are centered on the position of the supernova explosion,marked with a cross, with North-up and East-left. They have the same physical sizeof 30 kpc at the distance of each galaxy.

37

Fig. 2.3.— Cumulative fraction with the oxygen abundance of the supernova hostgalaxies. The number of host galaxies of each supernova type is indicated in eachpanel and the lines correspond to: SN II (solid), SN Ib/c (dot-dashed) and SN Ia(dotted). The left panel includes host galaxies with redshifts z < 0.04, while theright panel includes host galaxies with redshifts 0.01 < z < 0.04. The thick line inthe right panel shows the cumulative distribution of the 15,116 star-forming galaxiesfrom SDSS-DR4 in the same redshift range.

38

Fig. 2.4.— Number ratio of SN Ib/c to SN II as a function of metallicity of thehost galaxies. The open circles are the values obtained with our sample from directlymeasured central metallicities, and the filled squares are the results from Prantzos &Boissier (2003) using absolute magnitudes as a proxy to host metallicities. The errorbars are obtained from Poisson statistics. The solid line shows the predicted ratiofrom the binary models of Eldridge (2007b); the dashed line is from the models ofsingle stars with rotation of Maeder & Meynet (2004); and the long-dashed line isfrom the single star models of Eldridge (2007b).

39

Fig. 2.5.— Cumulative fraction plot of the projected separation between thesupernova and its host for the reduced sample in the redshift range 0.01 < z < 0.04.The lines correspond to: SN II (solid), SN Ib/c (dot-dashed) and SN Ia (dotted).

40

SN Name Typea RA Dec Host Name Distanceb zc MB 12 + log(O/H)

(deg) (deg) (kpc) (mag) (dex)

1909A II: 210.5129 54.4661 NGC5457 . . . 0.00082 –21.03 9.12

1920A II: 128.8156 28.4754 NGC2608 3.6 0.00727 –20.21 9.12

1936A IIL 184.9830 5.3522 NGC4273 6.0 0.00783 –20.57 9.14

1951H II: 210.9804 54.3614 NGC5457 . . . 0.00082 –21.03 9.12

1961U IIL 178.2366 44.1504 NGC3938 . . . 0.00263 –20.01 9.20

1963J Ia 177.6596 55.3504 NGC3913 . . . 0.00318 –18.00 9.01

1963P Ia 41.5088 –7.5807 NGC1084 . . . 0.00436 –20.58 8.89

1964A IIPec 170.2256 53.1945 NGC3631 . . . 0.00387 –21.00 9.20

1964L Ib 178.2046 44.1293 NGC3938 . . . 0.00263 –20.01 9.20

1965L II 170.2423 53.1851 NGC3631 . . . 0.00387 –21.00 9.20

1966B IIL 191.9371 4.3256 NGC4688 . . . 0.00325 –17.64 8.77

1966E II: 183.4414 13.4149 NGC4189 6.8 0.00702 –20.20 9.14

1967C Ia 162.1044 12.5453 NGC3373 . . . 0.00433 –19.67 9.02

1969C Ia 175.3243 47.6925 NGC3811 3.1 0.01057 –20.76 9.06

1970G IIL 210.7535 54.2424 NGC5457 . . . 0.00082 –21.03 9.12

1971G Ia 183.0500 13.2378 NGC4165 4.8 0.00624 –18.33 8.94

1971I Ia 198.9567 41.9875 NGC5055 . . . 0.00131 –21.20 9.06

1971K II: 175.3082 47.6861 NGC3811 7.3 0.01057 –20.76 9.06

1971S II: 20.5505 0.9557 NGC0493 8.5 0.00789 –20.65 8.83

1973R IIP 170.0481 12.9987 NGC3627 . . . 0.00193 –21.17 9.16

(cont’d)

Table 2.1. Supernova and host galaxy data

41

Table 2.1—Continued



1976B Ib: 186.5508 13.1115 NGC4402 . . . 0.00079 –18.73 9.05

1978H II 174.8486 56.2669 NGC3780 3.2 0.00841 –20.97 9.16

1979B Ia 177.6808 55.3592 NGC3913 . . . 0.00318 –18.00 9.01

1983W Ia 170.1356 57.7820 NGC3625 1.4 0.00647 –19.60 8.97

1985G IIP 187.1683 9.2611 NGC4451 . . . 0.00258 –17.69 9.06

1985H II 161.6339 63.2304 NGC3359 . . . 0.00339 –20.53 8.74

1987C IIn 127.5056 52.6926 UGC04438 6.2 0.01415 –20.47 9.24

1988C Ia 114.4125 41.9499 UGC03933 9.4 0.01970 –21.20 9.06

1988L Ib 211.5913 50.7298 NGC5480 2.4 0.00637 –19.91 9.18

1988M II 187.9206 3.9231 NGC4496B 3.9 0.01511 –19.71 9.16

1988Q II 248.3362 34.8058 SDSSJ163320.66+344825.9 3.5 0.03511 –19.61 8.74

1988R Ia 207.0593 54.7982 MCG+09–23–009 4.5 0.02548 –20.48 8.97

1988Y Ia 41.2288 –8.4079 SDSSJ024456.18–082411.9 14.5 0.02948 –19.61 8.92

1989B Ia 170.0581 13.0054 NGC3627 . . . 0.00193 –21.17 9.16

1989C IIn 146.9395 2.6267 UGC05249 0.4 0.00625 –19.10 8.72

1989U II 148.2191 42.8622 UGC05295 18.7 0.01603 –20.23 9.17

1990B Ib 189.1410 11.2416 NGC4568 2.1 0.00744 –21.66 9.24

1990E IIP 39.8713 –8.1371 NGC1035 . . . 0.00429 –19.23 9.02

1991C II 173.1695 5.3421 SDSSJ113240.27+052035.9 4.4 0.02814 –18.05 8.73

1991G IIP 181.3702 50.5366 NGC4088 . . . 0.00262 –20.28 9.11

(cont’d)

42




1991J II 198.1567 12.5843 NGC5020 15.9 0.01097 –20.64 9.16

1991L Ib/c 250.3127 39.2924 MCG+07–34–134 5.1 0.03054 –19.86 8.97

1991T IaPec: 188.5425 2.6657 NGC4527 6.7 0.00529 –21.53 9.12

1991aj Ia 247.4371 41.7786 MCG+07–34–084 1.8 0.03139 –18.61 8.88

1992B Ia 169.0415 55.4919 SDSSJ111608.45+552925.1 15.7 0.05760 –20.30 9.10

1992bt IIP 174.8323 56.2711 NGC3780 4.0 0.00841 –20.97 9.16

1993E II 145.9340 49.2989 SDSSJ094345.02+491748.1 6.0 0.02510 –18.41 8.77

1993I Ia 188.6779 9.0022 MCG+02–32–144 13.6 0.04295 –21.13 9.14

1993N IIPec 157.4431 13.0208 UGC05695 2.5 0.00979 –19.00 8.96

1994J Ia 150.4032 54.5824 SDSSJ100136.65+543458.8 2.7 0.05632 –20.12 9.02

1994N II 157.4455 13.0209 UGC05695 2.3 0.00979 –19.00 8.96

1994P II 179.8378 52.7164 UGC06983 . . . 0.00361 –18.37 8.73

1994Q Ia 252.4631 40.4322 SDSSJ164951.17+402600.2 2.6 0.02946 –19.70 8.98

1994ak IIn 138.5061 40.1060 NGC2782 9.5 0.00853 –20.90 8.94

1995I II 200.2494 3.5989 SDSSJ132059.90+033550.6 8.2 0.07565 –19.92 8.76

1995J II 188.1598 63.8859 UGC07700 8.8 0.00989 –19.00 8.63

1995R Ia 208.3858 –1.1917 UGC08801 6.9 0.02374 –21.28 9.06

1995T Ia 336.8017 –9.4957 SDSSJ222712.65–092942.0 4.6 0.05640 –20.28 8.95

1995V II 41.6115 –0.4988 NGC1087 1.9 0.00507 –20.58 9.10

1995ah II 4.7955 15.1073 SDSSJ001911.01+150622.7 1.1 0.01456 –16.67 8.27

(cont’d)

43




1996ak II 198.0184 46.1940 NGC5021 12.2 0.02864 –21.86 8.85

1996an IIP 41.5035 –7.5723 NGC1084 . . . 0.00436 –20.58 8.89

1996aq Ic 215.5947 –0.3900 NGC5584 2.3 0.00553 –19.83 9.03

1996bg II 13.2182 –0.0753 SDSSJ005252.42–000427.8 6.8 0.11571 –21.89 9.06

1996bu IIPec 170.2471 53.2023 NGC3631 . . . 0.00387 –21.00 9.20

1997X Ib/c 192.0596 –3.3331 NGC4691 . . . 0.00372 –19.58 9.06

1997bo II 184.9694 7.7305 SDSSJ121952.54+074352.3 0.9 0.01257 –16.91 7.94

1997bs IIn 170.0594 12.9721 NGC3627 . . . 0.00193 –21.17 9.16

1997bz Ia 170.6061 1.1893 SDSSJ112225.46+011122.2 0.4 0.02980 –18.48 8.94

1997cs IIPec 228.4158 2.8957 SDSSJ151338.52+025342.4 14.2 0.03690 –19.95 8.88

1997ef IcPec 119.2618 49.5612 UGC04107 6.0 0.01168 –20.13 9.14

1997ei Ic 178.7499 58.4907 NGC3963 3.3 0.01063 –21.07 9.19

1998S IIn 176.5257 47.4821 NGC3877 . . . 0.00303 –20.37 9.18

1998ab IaPec 192.1968 41.9245 NGC4704 8.4 0.02725 –21.05 9.16

1998aq Ia 179.1075 55.1282 NGC3982 . . . 0.00395 –19.83 9.13

1998bv IIPec 159.6058 47.7092 HS1035+4758 0.6 0.00526 –15.25 8.11

1998dk Ia 3.6340 –0.7364 UGC00139 1.6 0.01319 –20.19 9.09

1998dl II 41.5061 –7.5736 NGC1084 . . . 0.00436 –20.58 8.89

1999ab II 156.2527 53.8769 SDSSJ102500.07+535230.0 5.6 0.03215 –19.72 8.65

1999ap II 127.4477 4.6274 SDSSJ082947.33+043735.8 2.7 0.03901 –19.03 8.41

(cont’d)

44




1999ay II: 221.1827 58.9284 SDSSJ144444.23+585544.4 3.1 0.04379 –18.56 8.69

1999bg II: 181.0304 62.5003 IC0758 . . . 0.00426 –18.13 8.81

1999br IIPec 195.1742 2.4961 NGC4900 . . . 0.00322 –19.04 9.18

1999bw Ib: 154.9450 45.5264 NGC3198 . . . 0.00222 –20.48 9.10

1999cb Ia 246.4517 40.3423 SDSSJ162549.27+402042.7 8.3 0.02904 –20.45 8.85

1999ce Ia 172.2149 57.1344 SDSSJ112851.62+570803.1 1.5 0.07748 –20.07 8.84

1999eh Ib 137.3859 33.1213 NGC2770 2.4 0.00645 –20.73 9.00

1999gb IIn 122.5571 33.9583 NGC2532 7.1 0.01751 –21.77 9.13

1999gk II 190.9663 –0.5487 NGC4653 9.6 0.00875 –20.29 9.04

2000I II 145.1731 11.8857 NGC2958 4.5 0.02175 –21.54 9.18

2000L II 153.7576 65.1343 UGC05520 5.1 0.01101 –19.83 8.94

2000cv Ia 183.9259 61.8900 SDSSJ121540.80+615323.6 4.2 0.02026 –19.68 9.19

2000db II 178.4156 47.8629 NGC3949 . . . 0.00285 –19.93 8.77

2000de Ib 186.3007 54.5081 NGC4384 1.8 0.00833 –19.66 8.89

2000dq II 32.0874 1.1836 SDSSJ020820.18+011103.5 10.0 0.04232 –21.14 9.27

2001D II 176.2124 –1.6051 IC0728 9.3 0.02840 –21.65 9.20

2001K II 168.4838 12.3017 IC0677 1.5 0.01080 –20.55 9.08

2001W II 250.3129 39.2957 MCG+07–34–134 12.3 0.03054 –19.86 8.97

2001ad IIb 261.0101 58.9978 NGC6373 10.0 0.01103 –19.23 8.83

2001ae II 200.6090 –2.4236 IC4229 9.4 0.02320 –20.93 9.19

(cont’d)

45




2001ag II 143.6292 46.4607 MCG+08–18–009 3.5 0.02649 –19.82 9.10

2001ai Ic 205.4140 55.6681 NGC5278 5.1 0.02597 –21.02 9.10

2001bg Ia 128.8286 28.4683 NGC2608 4.5 0.00727 –20.21 9.12

2001bp Ia 240.5385 36.7189 SDSSJ160208.90+364313.8 12.6 0.09484 –20.53 9.10

2001dr II 195.6629 50.4443 NGC4932 12.2 0.02350 –21.04 8.77

2001dy II 256.2476 23.1684 SDSSJ170459.90+231008.8 4.3 0.03007 –20.21 9.17

2001em Ib/c 325.5986 12.4975 NGC7112 4.6 0.01951 –20.63 9.16

2001er Ia 148.3361 42.8458 UGC05301 2.0 0.01615 –19.81 8.70

2001fb II 2.5256 –0.4384 SDSSJ001006.62–002609.6 6.9 0.03215 –20.37 8.81

2001fg Ia 318.1885 –0.8767 SDSSJ211245.45–005232.3 3.1 0.03164 –18.27 8.75

2001im IIb 31.0700 0.6531 SDSSJ020416.94+003909.7 3.5 0.07501 –20.05 8.85

2001km Ia 137.9099 –0.7150 SDSSJ091138.38–004254.0 0.1 0.07022 –19.73 8.95

2001kt Ia 244.3058 48.4744 SDSSJ161713.39+482827.8 0.3 0.10373 –20.50 9.15

2002I Ia 200.6118 –2.4178 IC4229 5.2 0.02320 –20.93 9.19

2002ao Ic: 217.3989 –0.0155 UGC09299 2.4 0.00517 –18.31 8.57

2002at II: 173.0915 0.8024 NGC3720 3.3 0.01997 –21.09 9.20

2002ca II 203.1313 1.8450 UGC08521 3.6 0.01090 –19.50 9.06

2002cb IIn 196.0998 47.5981 MCG+08–24–034 7.1 0.02956 –20.32 9.07

2002ce II 128.3421 29.5336 NGC2604A 3.7 0.00694 –19.04 8.98

2002cg Ic 247.2520 41.2834 UGC10415 2.0 0.03189 –21.58 9.29

(cont’d)

46




2002ck Ia 236.7531 –0.9904 UGC10030 12.4 0.02975 –21.41 9.23

2002ea IIn 32.1045 14.3480 NGC0820 1.7 0.01477 –21.19 9.15

2002ew II 313.6271 –0.1406 SDSSJ205430.46–000820.8 3.1 0.02992 –19.28 8.74

2002fs Ia 343.5037 14.6568 SDSSJ225400.65+143924.5 2.5 0.03828 –18.41 8.56

2002gb Ia 340.8388 –0.1147 SDSSJ224321.26–000654.7 2.6 0.07392 –18.62 8.61

2002gf Ia 313.8245 –0.0726 SDSSJ205518.46–000424.4 14.7 0.08622 –21.05 9.14

2002gr Ia 36.8051 0.8994 SDSSJ022713.44+005356.0 6.3 0.09102 –19.19 8.72

2002ha Ia 311.8274 0.3127 NGC6962 8.3 0.01447 –21.58 8.94

2002hg II 159.2942 12.6536 NGC3306 1.5 0.00994 –19.91 8.86

2002hn Ic 122.5623 33.9554 NGC2532 2.0 0.01751 –21.77 9.13

2002ii Ia 330.4684 –1.1734 SDSSJ220152.75–011018.8 13.2 0.10059 –21.30 9.16

2002in II 326.0724 –0.2421 SDSSJ214417.38–001431.5 0.7 0.07553 –18.01 8.38

2002iq II 338.0294 0.9156 SDSSJ223207.06+005456.0 1.6 0.05575 –18.44 8.29

2002jo Ia 219.5684 40.4539 NGC5708 2.0 0.00904 –20.33 9.00

2002ln II: 249.8539 41.7914 SDSSJ163925.01+414737.0 20.0 0.14199 –21.01 9.09

2003I Ib 141.8728 3.9293 IC2481 3.9 0.01778 –20.59 9.06

2003L Ic 165.8014 11.0772 NGC3506 4.2 0.02119 –21.93 9.12

2003S Ia 205.3312 55.6768 SDSSJ134119.63+554040.1 3.0 0.03982 –20.63 8.94

2003ci II 167.5993 4.8266 UGC06212 10.3 0.03034 –21.69 9.18

2003cn II 196.9044 –0.9472 IC0849 11.7 0.01808 –20.56 9.11

(cont’d)

47




2003da II 140.9482 42.1802 UGC04992 3.5 0.01385 –19.88 8.87

2003dg Ib/cPe 179.3832 –1.2538 UGC06934 1.5 0.01840 –20.73 8.98

2003ds Ic 154.7431 46.4543 MCG+08–19–017 2.4 0.02993 –19.76 8.95

2003dt Ia 311.8232 0.3119 NGC6962 10.8 0.01447 –21.58 8.94

2003du Ia 218.6492 59.3344 UGC09391 2.2 0.00635 –17.83 8.65

2003eh Ia 167.1014 3.4964 MCG+01–29–003 4.4 0.02539 –21.36 9.18

2003ej II 189.7962 0.7251 UGC07820 10.2 0.01697 –20.34 9.01

2003ez Ia 190.7995 –2.5011 SDSSJ124312.22–023003.2 5.1 0.04800 –21.04 9.19

2003ke IIn 146.4707 34.6837 MCG+06–22–009 4.8 0.02043 –20.90 9.10

2003la II 157.5842 61.2635 MCG+10–15–089 3.2 0.03071 –20.39 9.12

2003ld II 3.9660 16.0893 UGC00148 2.1 0.01386 –19.56 8.78

2004C Ic 171.8738 56.8801 NGC3683 2.7 0.00570 –20.04 9.06

2004F IInPec 49.4742 –7.2953 NGC1285 3.7 0.01752 –21.49 9.24

2004G II 218.3392 4.4471 NGC5668 5.7 0.00527 –19.94 8.72

2004I II 40.8801 0.3085 NGC1072 3.2 0.02600 –21.75 9.06

2004ak II 127.4482 48.7730 UGC04436 16.8 0.02402 –20.82 9.11

2004aq II 179.6135 10.0179 NGC4012 4.2 0.01388 –20.50 8.97

2004at Ia 164.6883 59.4867 MCG+10–16–037 5.9 0.02203 –19.00 8.68

2004bk Ia 204.3697 4.1031 NGC5246 5.3 0.02306 –20.73 9.20

2004bl Ia 183.8003 –3.4384 MCG+00–31–042 4.9 0.01739 –20.66 8.60

(cont’d)

48




2004bn II 163.1294 7.2281 NGC3441 5.4 0.02173 –21.61 9.18

2004cc Ic 189.1433 11.2424 NGC4568 2.3 0.00744 –21.66 9.24

2004cj Ia 155.7088 11.7030 SDSSJ102250.11+114210.8 0.2 0.10191 –20.86 9.16

2004cm II 211.8540 55.1031 NGC5486 . . . 0.00453 –18.27 8.41

2004dt Ia 30.5532 –0.0976 NGC0799 4.7 0.01944 –20.74 9.08

2004eb II 262.1012 57.5460 NGC6387 2.5 0.02853 –20.58 8.82

2004el II 224.9659 54.6183 MCG+09–25–004 5.1 0.02627 –19.77 8.93

2004ey Ia 327.2825 0.4442 UGC11816 4.0 0.01583 –19.96 9.10

2004fc II 27.7660 –9.7019 NGC0701 0.5 0.00622 –19.79 9.01

2004hw Ia 357.8223 1.2276 SDSSJ235117.28+011339.7 1.2 0.05993 –17.53 8.23

2004hy II 316.5177 1.2158 SDSSJ210604.63+011258.0 6.4 0.05806 –19.26 8.69

2004ia Ia 34.5096 –0.5590 SDSSJ021802.06–003334.9 10.9 0.14370 –21.77 9.06

2004ib Ic 40.2350 –0.1801 SDSSJ024056.35–001045.1 3.5 0.05621 –19.33 8.50

2004ie Ia 330.4441 1.2367 SDSSJ220147.09+011412.8 7.5 0.05133 –19.13 8.86

2004ig Ia 1.4658 –0.9959 SDSSJ000551.98–005946.7 10.1 0.18300 –20.93 9.11

2005H II 32.4105 –10.1454 NGC0838 0.8 0.01284 –20.27 9.12

2005J II 179.6189 10.0192 NGC4012 5.2 0.01388 –20.50 8.97

2005O Ib 160.5763 –0.3781 NGC3340 2.6 0.01814 –21.25 9.12

2005Y II 24.4054 0.0416 UGC01159 1.8 0.01640 –19.19 8.84

2005ab II 190.2716 50.3823 NGC4617 12.9 0.01490 –21.47 9.11

2005ad II 37.1227 –1.1389 NGC0941 4.8 0.00534 –19.06 8.75

(cont’d)

49




2005ay II 178.2003 44.1051 NGC3938 . . . 0.00263 –20.01 9.20

2005bb II 194.2995 –1.7048 UGC08067 1.6 0.00913 –19.44 8.88

2005bc Ia 219.3127 38.4564 NGC5698 2.3 0.01210 –20.60 9.31

2005bk Ic 240.5710 42.9154 MCG+07–33–027 3.5 0.02445 –20.64 9.09

2005ci II 218.6870 48.6722 NGC5682 1.1 0.00759 –18.75 8.69

2005cr Ia 185.5717 12.3970 SDSSJ122216.99+122348.8 1.5 0.02139 –18.62 8.90

2005dp II 216.9026 41.2542 NGC5630 2.5 0.00891 –20.37 8.74

2005eh Ia: 327.4350 0.6576 SDSSJ214944.35+003924.8 5.9 0.12488 –20.05 8.79

2005em IIb: 48.4488 –0.2436 SDSSJ031347.86–001435.9 1.2 0.02517 –16.04 8.27

2005en II 119.8026 32.9166 UGC04132 3.2 0.01730 –22.21 9.23

2005eo Ic 119.8083 32.9218 UGC04132 10.2 0.01730 –22.21 9.23

2005gp Ia 55.4970 –0.7827 SDSSJ034159.35–004658.4 2.8 0.12660 –19.63 9.00

2005hk IaPec 6.9620 –1.1979 UGC00272 4.5 0.01301 –19.47 8.70

2005hl Ib 313.8325 0.5430 SDSSJ205519.27+003226.2 5.2 0.02319 –20.52 9.18

2005ho Ia 14.8504 0.0026 SDSSJ005924.12+000009.7 0.5 0.06283 –19.57 8.69

2005if Ia 52.5536 –0.9746 SDSSJ033012.89–005828.2 0.5 0.06706 –19.89 9.01

2005ij Ia 46.0886 –1.0629 SDSSJ030421.26–010347.2 1.2 0.12446 –21.17 8.87

2005ip II 143.0267 8.4457 NGC2906 2.3 0.00678 –20.29 9.13

2005ir Ia 19.1824 0.7946 SDSSJ011643.87+004737.0 5.4 0.07635 –20.26 9.09

2005kb II 12.7112 0.8536 SDSSJ005051.62+005104.3 4.8 0.01526 –18.76 8.38

2005ks Ib/c 324.4857 –0.0325 SDSSJ213756.52–000157.7 1.8 0.09866 –19.58 8.91

(cont’d)

50




2005ku Ia 344.9275 –0.0137 SDSSJ225942.67–000049.0 0.9 0.04541 –20.07 8.88

2005lj Ia 29.4293 –0.1794 SDSSJ015743.10–001045.8 1.6 0.07774 –19.78 8.88

2005lm Ia 3.7705 0.3551 SDSSJ001504.87+002118.4 0.9 0.08466 –18.41 8.70

2005mb II 196.4686 41.7164 NGC4963 10.2 0.02384 –21.19 9.25

2005mf Ic 137.1764 44.8143 UGC04798 7.3 0.02675 –20.45 9.05

2005mn Ib 57.3268 –0.6921 SDSSJ034918.34–004129.4 2.3 0.04737 –20.17 8.73

2005ms Ia 132.3098 36.1300 UGC04614 23.2 0.02523 –20.93 9.06

2006E Ia 208.3688 5.2062 NGC5338 . . . 0.00270 –16.76 8.94

2006am IIn 216.9052 41.2598 NGC5630 2.0 0.00891 –20.37 8.74

2006ar Ia 159.3781 65.0161 SDSSJ103731.99+650105.9 5.1 0.02249 –20.33 9.09

2006ck Ic 197.4184 –1.0492 UGC08238 5.4 0.02443 –20.91 8.99

2006ct Ia 182.4876 47.0955 SDSSJ120956.71+470545.6 2.4 0.03145 –19.57 8.71

2006db IIn 178.9112 44.3837 SDSSJ115538.31+442301.8 2.0 0.02308 –17.54 8.33

2006ed II 225.9610 42.1135 UGC09684 2.5 0.01687 –20.21 9.02

2006fe Ia 313.0382 –0.5111 SDSSJ205209.10–003039.2 1.6 0.07049 –20.48 9.10

2006fi Ia 334.9594 0.0244 SDSSJ221950.56+000125.2 19.2 0.23062 –21.98 8.96

2006fo Ic 38.1620 0.6175 UGC02019 2.4 0.02074 –20.41 9.02

2006fq IIP 5.0034 –0.6250 SDSSJ002000.77–003731.4 2.1 0.06790 –20.10 8.95

2006fs Ia 317.4958 0.4088 SDSSJ210958.97+002431.0 1.4 0.09914 –21.29 8.95

2006fu Ia: 357.7850 –0.7464 SDSSJ235108.38–004447.8 3.3 0.19852 –20.50 8.87

(cont’d)

51




2006fy Ia 351.6675 –0.8403 SDSSJ232640.14–005026.2 2.1 0.08265 –19.90 8.72

2006gb Ia 359.8188 –1.2504 SDSSJ235916.46–011502.5 5.4 0.26607 –21.58 9.13

2006gl Ia 16.4588 0.1448 SDSSJ010549.99+000843.4 11.3 0.26535 –20.78 9.02

2006gx Ia 42.0587 –0.3470 SDSSJ024814.09–002048.5 2.5 0.18063 –19.77 8.84

2006ha Ia 344.6429 15.1738 IC1461 2.2 0.03084 –20.76 9.20

2006iq Ia 324.8906 10.4849 SDSSJ213933.98+102900.6 8.8 0.07885 –20.83 9.23

2006iv IIb 177.0515 54.9874 UGC06774 3.2 0.00810 –18.82 8.53

2006iw II 350.3312 0.2597 SDSSJ232119.13+001532.8 3.4 0.03078 –18.73 8.86

2006ix II 359.8085 –0.3110 SDSSJ235913.99–001839.2 1.0 0.07565 –19.51 8.70

2006jc Ib/cPe 139.3367 41.9091 UGC04904 1.6 0.00555 –15.98 8.48

2006ju Ia: 351.1625 –0.7183 SDSSJ232438.93–004304.4 5.1 0.14864 –21.54 9.15

2006kh II 27.2994 –0.6052 SDSSJ014911.88–003618.4 0.8 0.05964 –18.18 9.03

2006kn II 314.5217 0.9005 SDSSJ205805.21+005402.2 1.1 0.12020 –20.10 8.43

2006kq Ia 318.9024 –0.3214 SDSSJ211536.50–001918.1 5.0 0.19827 –21.36 9.06

2006lb Ia 49.8675 –0.3180 SDSSJ031928.18–001904.8 0.7 0.18191 –20.30 8.98

2006nd Ia 341.2461 –1.0066 2MASXJ22445879–0100229 8.2 0.12876 –21.08 9.08

2006ns II 323.3843 0.7848 SDSSJ213332.23+004705.3 0.5 0.11991 –20.53 8.63

2006oa Ia 320.9289 –0.8435 SDSSJ212342.96–005035.2 1.6 0.06256 –17.95 8.45

2006pc II 26.0441 –0.1559 SDSSJ014411.09–000916.1 9.7 0.05532 –19.17 8.72

2006pt Ia 36.8174 –0.3935 SDSSJ022716.08–002335.5 7.4 0.29883 –20.90 9.07

2006qk Ic: 336.3849 0.1542 SDSSJ222532.38+000914.8 0.4 0.05831 –17.89 8.95

(cont’d)

52




2006rz Ia 56.5282 0.3897 SDSSJ034606.86+002325.1 1.4 0.03089 –18.35 8.55

2006ss IIb 215.1144 35.1952 NGC5579 6.7 0.01200 –19.62 8.95

2006te Ia 122.9291 41.5547 SDSSJ081143.49+413318.0 3.7 0.03155 –20.43 9.10

2007F Ia 195.8128 50.6188 UGC08162 5.9 0.02362 –20.07 9.05

2007I Ic 179.8048 –1.6053 SDSSJ115913.13–013616.2 1.3 0.02156 –16.73 8.34

2007S Ia 150.1302 4.4073 UGC05378 3.4 0.01385 –20.05 9.07

2007aa II 180.1154 –1.0810 NGC4030 . . . 0.00495 –20.76 9.24

2007af Ia 215.5876 –0.3938 NGC5584 6.1 0.00553 –19.83 9.03

2007av II 158.6798 11.1938 NGC3279 . . . 0.00464 –19.35 9.02

2007be IIP: 189.5277 –0.0309 UGC07800 6.7 0.01251 –19.76 9.05

2007bk Ia 232.1899 58.8702 SDSSJ152845.00+585200.1 8.7 0.03214 –18.17 8.27

Note. — Complete table is available online.

aSupernova classification in the SAI Catalog. The types followed by a colon indicate a provisional classification in the

SAI Catalog.

bProjected SN-Host distance.

cRedshift of the host galaxy from SDSS-DR4.

53

Chapter 3

Discovery of the Dust-Enshrouded Progenitor

of the Type IIn SN 2008S with Spitzer

3.1. Introduction

Over the last ∼ 20 years, several significant milestones have been reached in

the pre-explosion detection of core-collapse supernova progenitors. These began

with the “peculiar” type II-P supernova 1987A in the Large Magellanic Cloud

(e.g., Menzies et al. 1987), where a cataloged ∼ 20 M⊙ blue supergiant star was

identified as the progenitor (Sk −69 202; e.g., West et al. 1987). Next came the

transition type IIb 1993J in M81, with a progenitor identified as a red supergiant

in a binary system (e.g., Podsiadlowski et al. 1993; Maund et al. 2004). During

the last decade, analyses of pre-explosion archival optical imaging of nearby galaxies

obtained (mainly) with the Hubble Space Telescope have convincingly shown red

supergiants with masses 8 M⊙ ≤ M ≤ 20 M⊙ to be the typical progenitors of

type II-P supernovae (e.g., Smartt et al. 2004; Li et al. 2007), the most common

core-collapse supernovae. Curiously, the progenitors of nearby type Ib/c supernovae,

thought to result from very massive (∼> 20 M⊙) stars with strong winds that end

54

their lives as Wolf-Rayet stars, have evaded optical detection (e.g., Crockett et al.

2008).

The rarest and most diverse class of core-collapse supernovae are the type IIn

(Schlegel 1990), which represent ∼ 2 − 5% of all type II supernovae (e.g., Capellaro

et al. 1997). Their optical spectra, dominated by Balmer lines in emission, and

slowly declining light curves show clear signatures of interactions between the

supernova ejecta and a dense, hydrogen-rich circumstellar medium (e.g., Filippenko

1997). Mainly due to their low frequencies, high mass loss rates, and the massive

circumstellar envelopes generally required to explain the observations, some luminous

type IIn supernovae have been associated with the deaths of the most massive

stars (e.g., Gal-Yam et al. 2007; Smith 2008 and references therein). Recently,

evidence for this association has increased with the report of a very luminous source

in pre-explosion images of the type IIn SN 2005gl (Gal-Yam et al. 2007) and the

discovery of an LBV eruption two years before the explosion of SN 2006jc (Pastorello

et al. 2007). On the other hand, some low luminosity type IIn have been associated

with the super-outbursts of LBVs like η Carinae (e.g., van Dyk et al. 2000; van Dyk

et al. 2006).

The appearance of the type IIn SN 2008S in the nearby galaxy NGC 6946

(d ≃ 5.6 Mpc; Sahu et al. 2006) was fortuitous, since a massive stellar progenitor

would be relatively easy to find. However, pre-explosion images serendipitously

obtained from the Large Binocular Telescope revealed nothing at the position of

55

SN 2008S, allowing us to put stringent limits on the optical emission. In this paper,

we report the discovery of an infrared point source coincident with the site of

SN 2008S using archival Spitzer Space Telescope data. The Spitzer mid-IR detection,

and deep optical non-detections, of the progenitor are the tell-tale signs of a ∼ 10 M⊙

star obscured by dust. We describe the available data in §3.2, our analysis in §3.3,

and our conclusions in §3.4.

3.2. Searching for the Progenitor

NGC 6946 is quite a remarkable galaxy, giving birth to (at least) nine SNe

in the last century. The latest event discovered in NGC 6946 is SN 2008S, found

on February 1.79 UT at ∼ 17.6 mag (Arbour & Boles 2008) and located 52′′ West

and 196′′ South of the nucleus of NGC 6946. It was spectroscopically classified

as a likely young type IIn supernova from the presence of narrow Balmer lines in

emission, highly reddened by internal extinction with a measured Na D absorption

equivalent width of 5 A (Stanishev et al. 2008). Steele et al. (2008) later reported

that it had a peculiar spectrum due to the presence of narrow emission lines from

the [Ca II] 730 nm doublet, Ca II infrared triplet, and many weak Fe II features.

The spectral properties and low peak luminosity led Steele et al. (2008) propose

that SN 2008S was a supernova impostor such as SN 1997bs (van Dyk et al. 2000).

56

Accurate coordinates are needed in order to search for the progenitor in

pre-explosion images. Fortunately, Swift started monitoring SN 2008S with UVOT

and XRT shortly after the discovery. We retrieved the UVOT ubv optical images

obtained on Feb. 4.8, 6.0, and 10.5 (UT) from the Swift archive. We used WCSTools

v3.6.7 (Mink 1999) and the USNO-B astrometric catalog (Monet et al. 2003) to

obtain astrometric solutions for the images. The mean coordinates of SN 2008S are

α = 20h34m45.s37, δ = 6005′58.3′′ (J2000.0), with rms uncertainties of σα = 0.′′5 and

σδ = 0.′′3.

The Large Binocular Telescope (Hill et al. 2006) obtained deep optical images

of NGC 6946 on 19−21 May 2007, 225 days before discovery, during Science

Demonstration Time using the LBC/Blue camera (Ragazzoni et al. 2006; Giallongo

et al. 2008). We combined the 12×300 sec images obtained using the U filter (seeing

1.′′0), and the 4 × 300 sec images obtained using the B and V filters (seeing 1.′′5).

After finding an astrometric solution for the combined images using the USNO-B

catalog (σα ≃ σδ = 0.′′2), we do not detect a source at the position of SN 2008S (see

Fig. 3.1). After calibrating the images using ancillary optical data obtained by the

Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al. 2003) and Swift,

we obtain 3σ upper limits on the progenitor magnitudes of U > 25.8, B > 25.9 and

V > 26.0, which correspond to absolute magnitudes MU > −4.8, MB > −4.3, and

MV > −3.8, correcting for AV = 1.1 mag of Galactic extinction (Schlegel et al.

1998). The upper limits are calculated using aperture photometry from the standard

57

deviation of the sky at the SN position using a 10 pixel (2.′′2) diameter aperture.

We correct these values with aperture corrections estimated using bright stars.

The ∼ 0.2 mag uncertainties in the 3σ upper limits are due to the uncertainties in

the aperture corrections and the standard deviation of the sky (which is estimated

from the rms variations in the standard deviation measured in equal-sized apertures

placed in the background around the SN position). Welch et al. (2008) reported 3σ

upper limits from pre-explosion Gemini/GMOS observations of V > 24.0, R > 24.5

and I > 22.9. These correspond to absolute magnitudes MV > −6.8, MR > −5.2,

and MI > −6.5, correcting for Galactic extinction.

Such a deep non-detection led us to investigate IRAC (3.6 − 8.0 µm; Fazio

et al. 2004) and MIPS (24 − 160 µm; Rieke et al. 2004) images obtained by the

SINGS Legacy Survey in 2004. We astrometrically calibrated the images in the

same way as the optical images from Swift and LBT. We detect a point source

at α = 20h34m45.s35, δ = 6005′58.0′′ in the 4.5, 5.8, and 8.0 µm IRAC bands

(see Fig 3.1), with rms uncertainties σα = 0.′′5, σδ = 0.′′2. This is consistent with

the position of SN 2008S given the estimated uncertainties, and thus likely to be

the progenitor. The source is not detected at 3.6, 24, or 70 µm. We estimate a

probability of random coincidence given the uncertainty in the SN position (0.′′5) of

0.8% (0.02%) from the density of 4.5 micron sources (with [3.6]−[4.5] > 1.5 mag)

detected within a 1′ radius of the SN position.

58

We searched the Spitzer archive for all the programs that have observed

NGC 6946. Observations by the SINGS survey (PID: 159), and two programs

(PIs: Meikle, Sugerman, Barlow) monitoring the type II-P SNe 2002hh and 2004et

(PID: 230, 20256, 30292, 30494) provide a 2.5-year baseline (June 2004 − January

2007) of IRAC and MIPS observations prior to the discovery of SN 2008S. We used

aperture photometry (a 2 pixel extraction radius with aperture corrections) in the

flux-calibrated images provided by the Spitzer Science Center to derive light curves

for the progenitor. Fig. 3.2 shows the flux density as a function of time in the 4.5,

5.8, and 8.0 µm bands starting from June 2004. There is no sign of variability at the

∼ 10% level. The non-detection at 3.6, 24, and 70 µm in single and stacked images

allows us to place useful upper limits on these fluxes.

Finally, we searched the Chandra archive to determine if the progenitor was

an X-ray source. All five ACIS-S observations of NGC 6946 include the location of

SN 2008S. These observations include a 60 ks exposure in 2001, a 30 ks exposure

in 2002, and 3 × 30 ks exposures in 2004. No source is detected at the supernova

position in any of these images. We set a 3σ upper limit on the flux of the progenitor

of fX < 3 × 10−15 erg cm−2 s−1 (LX < 1037 erg s−1) in the broad X-ray band

(0.5−8 keV), which rules out a bright X-ray binary as the progenitor. This flux

limit corresponds to 20 counts in the longest exposure. Table 3.1 summarizes the

detections and 3σ upper limits on the progenitor fluxes.

59

3.3. Beneath the Shroud

The measured fluxes and upper limits in the mid-IR bands are shown in Fig. 3.3.

The shape of the spectral energy distribution (SED) suggests thermally-radiating

dust as the source of the emission. We derive a best-fit single-temperature blackbody

of T ≃ 440 K, with a luminosity of Lbol ≃ 3.5 × 104 L⊙ (d = 5.6 Mpc; Sahu et al.

2006), which implies a blackbody radius1 RBB ≃ 150 AU. This luminosity points to

a ∼ 10 M⊙ star at the end of its life (e.g., Meynet & Maeder 2003). The 3σ upper

limit at 70 µm further limits the total luminosity of the dust-enshrouded source and

the geometry of obscuring dust distribution.

As shown in Fig. 3.3, a blackbody yields a relatively poor fit to the data

(χ2 ≈ 4.9 per d.o.f.). The inability of a single-temperature blackbody to

accommodate the data follows primarily from the rapid change in the SED implied

by the 3.6 µm upper limit and the 4.5 µm detection. Radiation transport calculations

using DUSTY (Ivezic & Elitzur 1997) were performed as a sanity check. Using a

central incident blackbody with T = 3000 − 20000 K we calculated the emergent

spectrum from a spherical dusty shell extended over approximately one decade

in radius. As expected, the best correspondence with the data is obtained for a

total optical depth at 8.0 µm of order unity, although the precise value depends

1These values would change to Lbol ≃ 8 × 104 L⊙ and RBB ≃ 230 AU if we assume an extreme

distance to NGC 6946 of 8.5 Mpc, which is the 3σ upper limit of the distance used by Li et al.

(2005; 5.5±1.0 Mpc).

60

on the assumed radial gradient of the density, the radial extent of the obscuring

medium, and the mixture of grain types. Although a detailed investigation of the

dust properties is beyond the scope of this paper, we note that the strong evolution

in the SED between 3.6 and 4.5 µm may signal the need for relatively large grains

(e.g., Ivezic & Elitzur 1996).

We can estimate the mass of obscuring gas and dust by assuming that the

medium is marginally optically thick at 8.0 µm. Setting τλ ≈ 1 ≈ κλρRBB,

and assuming ρ = M/(4πR3BB/3), we find that M ∼ 10−3κ−1

λ, 10 M⊙, where

κλ, 10 = κλ/10 cm2 g−1 is a typical value for the Rosseland-mean dust opacity for

gas at ∼ 440 K (e.g., Semenov et al. 2003). This suggests a gas density on the scale

RBB of n ∼ 3 × 107 cm−3. We also estimate a minimum mass loss rate from the

progenitor of Mmin = 4πR2BBρcg ∼ 10−5 M⊙ yr−1, where cg ∼ 2 km s−1 is the gas

sound speed in the medium on the scale of RBB.

The lack of variability in the mid-IR fluxes (see Fig. 3.2) limits the expansion

velocity of the photosphere. Given our estimated temperature and luminosity,

keeping the mid-IR fluxes constant to within ∼ 10% over the ∼ 103 days covered

by the observations means that the dust photosphere cannot be expanding by more

than ∼ 10 km s−1, which is below the escape velocity of 13 km s−1 for a 10 M⊙ star

at the estimated photospheric radius of 150 AU. This is further evidence that the

dust is part of a relatively steady, massive wind rather than an explosively-expelled

dust shell.

61


Our pre-explosion detection of the progenitor of the type IIn SN 2008S is, to

the best of our knowledge, the first in the mid-IR. The Spitzer observations suggest

an enshrouded star with a mass of ∼ 10 M⊙, buried in ∼ 10−3 M⊙ of gas and dust.

If SN 2008S was a real supernova explosion, this is direct evidence that relatively

low-mass stars can end their lives as type IIn SNe when they have a sufficiently

dense CSM from a massive wind, as proposed by Chugai (1997b). If this event was

the luminous outburst of an LBV, it presents evidence for low-luminosity, low-mass

LBVs that have not been observed before2. These conclusions about the relatively

low mass hold even if the identification of the progenitor with the Spitzer source is

incorrect. In this case, we know the total extinction from the colours of the SN (see

below). As shown in Fig. 3.3, our optical limits with this extinction correspond to

mass limits of ∼< 12 M⊙ for red supergiants and ∼< 20 M⊙ for blue supergiants3.

Interestingly, we see luminous dust-enshrouded stars in the Milky Way and the

LMC whose physical properties match well the observed properties of the progenitor

of SN 2008S. van Loon et al. (2005, and references therein) studied the properties

2The lowest-mass LBVs known have initial masses of 20 − 25 M⊙ and luminosities > 105 L⊙

(e.g., Smith et al. 2004; Smith 2007).3We obtain an upper limit in the absolute optical magnitude of the progenitor of MV ∼> −7.1

if we assume an upper limit on the extinction estimate from the SN color (AV ≃ 3.5 mag; bluest

black-body possible) and an extreme distance to NGC 6946 of 8.5 Mpc.

62

(T⋆, Tdust, Lbol, M) of dust-enshrouded AGB stars and red supergiants in the LMC

using mid-IR observations. These are M-type stars with effective temperatures

∼ 2500 − 3800 K, which have strong winds with high (gas + dust) mass loss rates

(M ∼ 10−6 − 10−3 M⊙ yr−1), and warm dust emission from their dusty envelopes

(200 K < Tdust < 1300 K). Due to these similarities, we conclude that the progenitor

of SN 2008S was likely a dust-enshrouded AGB (core-collapse produced from electron

capture in the O-Ne-Mg core, e.g., Eldridge et al. 2007; Poelarends et al. 2008) or

red supergiant like the ones observed in the LMC.

Although the detection and physical properties of the progenitor are the main

results of this study, we can also try to understand something about the progenitor

and explosion mechanism from the supernova itself. The classification spectrum

of SN 2008S is similar to the published spectrum of SN 1997bs (van Dyk et al.

2000), which showed narrow Balmer lines in emission and many weaker Fe II lines

(V. Stanishev, priv. comm.; Steele et al. 2008). SN 2003gm had photometric and

spectroscopic characteristics similar to SN 1997bs (Maund et al. 2006). Since both

of these were faint (MV ∼ −14 mag) compared with the typical absolute magnitudes

at maximum of type II SNe (MV ∼ −16 to −18 mag), it is still debated whether

they were intrinsically faint explosions or super-outbursts of LBVs.

The early optical photometry obtained with Swift also indicates that SN 2008S

was a low-luminosity object, with MV ∼ −14 mag after correcting for the total

extinction along the line of sight. We estimate the total extinction for RV = 3.1

63

to be AV ≈ 2.5 mag from the observed color B − V ≃ 0.8 mag and assuming a

typical intrinsic temperature of ∼ 10000 K at this early phase of the evolution.

This value is roughly consistent with the estimated reddening obtained from the

reported equivalent width of the Na D absorption feature (2.5 < AV < 7.8; based

on Turatto et al. 2002). This implies the presence of significant internal extinction

with AV ≃ 1.4 mag after correcting for AV (Gal) = 1.1 mag. Although the light from

the supernova likely destroyed the dust that obscured the progenitor to significantly

beyond the blackbody scale of ∼ 150 AU, the existence of internal extinction in the

supernova light curve implies a more tenuous dusty obscuring medium on larger

scales. In fact, the rare detection of the [Ca II] 730 nm doublet in emission by Steele

et al. (2008) may provide direct, and independent, evidence for a significant amount

of dust in the CSM that was destroyed by the UV-optical flash (e.g., Shields et al.

1999). The future spectra and light curves of SN 2008S, optical as well as radio and

X-ray, should further probe the environment as they show signs of interactions with

the progenitors’s wind.

The field of supernova forensics has advanced rapidly in recent years, with

∼ 10 SN progenitors now known (e.g., Smartt et al. 2004; Li et al. 2007). Moving

forward, several groups are obtaining the data required to more fully characterize

the progenitors of future nearby SNe (e.g., Kochanek et al. 2008). We note that the

discovery of the progenitor of SN 2008S itself would not have been possible only

64

few years ago without Spitzer. Future multi-wavelength surveys of the local

universe are thus encouraged in order to catch other unexpected stellar phenomena,

potentially even before they occur.

65

LBT V IRAC 3.6 µm m

IRAC 5.8 IRAC 8.0µm µm MIPS 24 m

IRAC 4.5 µ

µ

E

N

Fig. 3.1.— Pre-supernova images (30′′ × 30′′) of the site of SN 2008S. We showthe LBT/LBC optical non-detection of the progenitor and the images obtained withSpitzer by the SINGS project at 3.6, 4.5, 5.8, 8.0, and 24 µm. The progenitor isclearly detected at 4.5, 5.8, and 8.0 µm. The circle in each panel has a radius of2′′ and is centered on the position of the supernova, corresponding to 4 times theastrometric uncertainty of 0.′′5. The dark line in the LBT image is bleeding from asaturated star.

66

Fig. 3.2.— Flux densities at 4.5, 5.8, and 8.0 µm as a function of time (in days beforethe discovery) for the progenitor of SN 2008S. The solid line in each panel shows themean for each band and the dashed lines show the rms deviations of ±3.3, 12.2, and13.0 µJy, respectively.

67

Fig. 3.3.— The spectral energy distribution of the progenitor of SN 2008S from Spitzer

observations. Detections are shown as open squares at 4.5, 5.8, and 8.0 µm. Upperlimits (3σ) from the combined images at 3.6 and 24 µm are also indicated. The solidline is the best-fit blackbody with T = 440 K. We also show the UBV upper limits(3σ) from LBT and the RI limits from Welch et al. (2008). The measured fluxesare not extinction corrected. The dotted line shows a reddened blackbody with theluminosity (104.5 L⊙) and effective temperature (103.5 K) of a 12 M⊙ red supergiant.The dashed line shows a reddened blackbody with the approximate temperature(104.2 K) and luminosity (105 L⊙) of the blue supergiant progenitor of SN 1987A,which has similar properties to the lowest luminosity LBVs observed (e.g., Smith2007). The models were reddened with AV = 2.5 mag, the total extinction estimatedfrom the colors of SN 2008S.

68

λ λFλ Source

(10−17 W m−2)

0.3-8 keV < 0.3 Chandra/ACIS-S

0.36 µm < 0.07 LBT/LBC-Blue



0.64 µm < 0.22 Welch et al. (2008)

0.80 µm < 0.63 Welch et al. (2008)

3.6 µm < 0.45 Spitzer/IRAC

4.5 µm 1.47 ± 0.22 Spitzer/IRAC



24 µm < 1.20 Spitzer/MIPS

70 µm < 40 Spitzer/MIPS

Table 3.1. Spectral Energy Distribution of the Progenitor of SN 2008S

69

Chapter 4

A New Class of Luminous Transients and A

First Census of Their Massive Stellar

Progenitors

4.1. Introduction

Identifying the progenitors of core-collapse supernovae, the outbursts of

Luminous Blue Variables (LBVs), and other massive-star transients is essential for

understanding the physics, demographics, variability, evolution, and end-states of

massive stars. The problem of identifying the progenitors of bright transients from

massive stars is difficult, and traditionally limited to serendipitous archival imaging of

nearby galaxies in the optical and near-infrared (e.g., van Dyk et al. 2003; Smartt et

al. 2004; Li et al. 2007; see the extensive summary in Smartt et al. 2009). Progenitor

searches are complemented by statistical studies of supernova environments within

their host galaxies, which provide indirect evidence for associations between certain

types of supernovae and broad classes of progenitors (e.g., James & Anderson 2006;

Kelly et al. 2008; Prieto et al. 2008b; Anderson & James 2008).

70

A much more direct method for understanding the relation between types

of massive stars and their transients is to catalog all of the massive stars in the

local universe (D ∼< 10 Mpc) before explosion. While surveys for bright optical

transients in the local universe are well-developed (e.g., Li et al. 2001), a fairly

complete census of the massive stars in nearby galaxies has only recently been

proposed and undertaken (Massey et al. 2006; Kochanek et al. 2008). Despite the

technical challenges required by the depth, area, and cadence of the observations,

these surveys are critical for our understanding of the one-to-one correspondence

between massive stars and their end-states, whether they are successful or failed

explosions (e.g., Kochanek et al. 2008). The long-term promise of these surveys

is to produce a catalog within which the characteristics of progenitors of future

supernovae are listed pre-explosion, as in the case of SN 1987A (West et al. 1987;

Menzies et al. 1987). They will provide an essential mechanism for understanding

the direct causal mapping between individual progenitor types and their transients

(see, e.g., Gal-Yam et al. 2007).

Here, we describe a new link in this causal mapping. The discovery by

Prieto et al. (2008c) and Prieto (2008d) of the dust-obscured progenitors of the

luminous outbursts in NGC 6946 (SN 2008S; Arbour & Boles 2008) and in NGC 300

(Monard 2008) with Spitzer, opens up qualitatively new possibilities in the study

of the connection between massive stars and their explosions. We show that these

discoveries allow us to make a strong — and perhaps unprecedented — connection

71

between a dust-enshrouded sub-population of massive stars and a new relatively

common class of bright transients.

The argument presented in this paper can be summarized in four points:

1. The transients SN 2008S and NGC 3001 constitute a class. Both have peak

absolute V -band magnitude MV ≈ −14± 1 (≈ 2− 3 magnitudes fainter than normal

core-collapse supernovae; e.g., Richardson et al. 2002), strong evidence for internal

extinction in their spectra, narrow emission lines (similar to low-luminosity Type IIn

supernovae; e.g., Filippenko 1997), and progenitors that are optically obscured and

deeply dust-embedded (dust-reprocessed emission giving blackbody temperatures

∼< 500 K), and with bolometric luminosities of ∼ 5 × 104 L⊙. In addition, they show

little infrared variability on a few-year baseline. The details of this unique class

of progenitor-transient pairs and its members, as well as a comparison with other

classes of optical transients are presented in §4.2. An in-depth search for analogs to

the progenitors of SN 2008S and NGC 300 in M33 is presented in §4.4.

2. Transients of this type are relatively common with respect to core-collapse

supernovae. A total of ≈ 22 core-collapse supernovae or supernova-like transients

1Throughout this paper we denote “the transient in NGC 300” as “NGC 300” (e.g., “the

progenitor of NGC 300”) unless we specifically refer to the host galaxy.

72

have been discovered within ≈ 10 Mpc since 1999. Sixteen were confirmed

supernovae, two were LBV eruptions, one was a Type IIn supernova (SN 2002bu

in NGC 4242, D ≈ 8 Mpc; Puckett & Gauthier 2002), whose relatively low peak

magnitude (MV ≈ −15; Hornoch 2002) suggests a close similarity with the remaining

two bright transients (see §4.2.5), which are of primary interest in this paper: SN

2008S (D ≈ 5.6 Mpc) and NGC 300 (D ≈ 1.9 Mpc), whose physical nature is

uncertain. They may be either true (but optically sub-luminous) supernovae, or

a new class of massive star eruptions. Taken at face value, these numbers imply

that the rate of SN 2008S-like transients is of order ∼ 10% of the supernova rate.2

Because of incompleteness, the true rate is likely higher. We discuss the frequency

of these events in detail in §4.3.

3. The progenitors of this class are extremely rare among massive stars at any

moment, in any star-forming galaxy. Although the bolometric luminosities of

the progenitors of SN 2008S and NGC 300 are unremarkable for massive stars

(∼ 5 × 104 L⊙), their colors put them in a class consisting of less than 10−4 of all

massive stars ([3.6] − [4.5]µm color ∼> 2.0 and ≈ 2.7, respectively). In a mid-infrared

(MIR) color-magnitude diagram (see Figs. 4.1 & 4.2), these progenitors lie at the

extremum of the AGB sequence in both luminosity and MIR color. We refer to them

as “extreme AGB” (EAGB) stars throughout this work. Because of their relatively

2Throughout this paper we use “supernova rate” to mean the core-collapse supernova rate unless

we specifically mention the contribution from Type Ia supernovae.

73

low bolometric luminosities, they are not η Carinae, cool hypergiant, or classical

LBV analogs (see, e.g., Smith 2008). They are thus distinct from the “supernova

impostors,” produced by bright outbursts of optically-luminous LBVs (e.g., SN

1997bs; van Dyk et al. 2000, 2003). In §4.4 we present results from a comprehensive

survey of M33 for massive stars with properties similar (in bolometric luminosity,

obscuration, and variability) to the progenitors of SN 2008S and NGC 300. We find

remarkably few. We compare with MIR surveys of the LMC, SMC, and NGC 300.

4. Conclusion: A large fraction of all massive stars undergo a dust-enshrouded phase

within ∼< 104 yr of explosion. This is the most natural explanation for the facts of

points (2) and (3) above. If these transients have a rate comparable to the supernova

rate (∼ 20%; point 2), then the timescale for the obscured phase is determined by the

ratio of the number of dust-obscured massive stars relative to the entire population

(∼< 10−4; see §4.4, §4.5) times the average lifetime of a massive star (∼ 107 yr).

Importantly, from the rarity of SN 2008S-like progenitors alone (∼ 10−4 of massive

stars; point 3), one would naively expect a comparable fraction of supernovae to

have progenitors of this type, if the dust-obscured phase occurs at a random time in

the life of a massive star. However, the relative frequency of these explosions (point

2) shows that this phase must come in the last ∼< 104 yr, just before explosion. Thus,

there must be a causal relation between the occurrence of the highly dust-enshrouded

phase and eruption. These points, together with a discussion of the theory of

74

the evolution of massive stars, the potential connection with electron-capture

supernovae, white dwarf birth, and other hypotheses for the physical mechanism of

2008S-like explosions, as well as a call for a more comprehensive Spitzer survey for

analogous sources within D ≈ 10 Mpc, are presented in §4.5.

4.2. The Class

We start by listing the objects we view as likely to represent this new class of

SN 2008S-like transients and progenitors. The two objects that define the class —

SN 2008S (§4.2.1) and NGC 300 (§4.2.2) — are unique among transient-progenitor

pairs. The progenitors are relatively low-luminosity, have low variability, and are

deeply dust-embedded on ∼ 100 AU scales. The transients are low-luminosity, with

spectra exhibiting both narrow Balmer lines (similar to low-luminosity IIn’s and

impostors), and [Ca II] in emission, and have rapidly decaying lightcurves compared

to IIP supernovae. The transient in M85 (§4.2.3) and SN 1999bw (§4.2.4) may also

be members of the same class, but we cannot confirm the existence of a dust-obscured

progenitor similar to SN 2008S/NGC 300. In §4.2.5, we contrast this class with

other peculiar outbursts, such as the supernova impostors and low-luminosity Type

IIP supernovae, and we note a number of other transients that are not excluded as

members of the SN 2008S-like class.

75

4.2.1. SN 2008S

SN 2008S in NGC 6946 (D ≈ 5.6 Mpc; Sahu et al. 2006) was discovered

February 1.79 UT (Arbour & Boles 2008). Because of the presence of narrow

Balmer lines (FWHM ≈ 1000 km s−1) it was initially classified as a young Type IIn

supernova. Stanishev et al. (2008) reported strong Na D absorption with equivalent

width of 5 A, indicating a high degree of internal extinction. On the basis of its

relatively low luminosity (MV ≃ −14) and peculiar spectrum (including the presence

of strong and narrow [Ca II] 730 nm doublet in emission), Steele et al. (2008)

proposed that SN 2008S was a supernova impostor such as SN 1997bs (van Dyk

et al. 2000). Botticella et al. (2009) present the late-time lightcurve of SN 2008S,

which shows evidence for a power-law time dependence with a slope indicative of

being powered by the radioactive decay 56Co. Following the suggestion by Prieto

et al. (2008c) and the discussion presented in §4.5, Botticella et al. argue that SN

2008S may have been an electron-capture supernova.

Prieto et al. (2008c) used a deep pre-explosion archival Large Binocular

Telescope (LBT) image of NGC 6946 to put stringent upper limits on the optical

emission from the site of the supernova (3σ limits of MU > −4.8, MB > −4.3,

and MV > −3.8). The failure to detect a progenitor in the optical led Prieto et

al. (2008c) to examine archival Spitzer IRAC imaging of NGC 6946 from the SINGS

Legacy Survey (Kennicutt et al. 2003) for the progenitor. A point source at the

76

location of SN 2008S was detected at 4.5, 5.8, and 8µm, but undetected at 3.6µm

and 24µm, leading to a lower limit on the [3.6] − [4.5] color of ∼> 2. The best-fitting

blackbody temperature to the SED was ≈ 440 K.3 The integrated luminosity was

≈ 3.5 × 104 L⊙, consistent with a relatively low-mass massive-star progenitor with a

zero-age main sequence (ZAMS) mass of M ≈ 10 M⊙. The progenitor SED is shown

in the left panel of Figure 4.5. Simple arguments suggest that the obscuration was

circumstellar, with an optical depth at visual wavelengths larger than a few on a

physical scale of order 150 AU (Prieto et al. 2008c). The explosion itself was also

serendipitously observed by Spitzer 5 days after discovery (Wesson et al. 2008), with

an infrared luminosity ∼50 times larger than the progenitor, suggesting a substantial

amount of dust-reprocessing (see Fig. 2 of Prieto et al. 2008e). The lack of variability

of the progenitor on a ∼ 3 yr timescale argues that the obscuring medium was either

a continuous wind with a steady photosphere or a (implausibly?) slowly (∼< 10 km

s−1) expanding shell.

4.2.2. NGC 300

A luminous optical transient (MV ≃ −13) in the nearby galaxy NGC 300

(D ≈ 1.9 Mpc; Gieren et al. 2005) was discovered by Monard (2008), and reported

by Berger & Soderberg (2008). The latter used archival imaging with ACS/WFC

onboard the Hubble Space Telescope (HST) to put very tight limits on the optical

3A blackbody provides a rather poor fit to the SED, perhaps indicating an interesting grain size

distribution in the obscuring material (Prieto et al. 2008c).

77

emission (see Fig. 4.5). These limits led them to suggest that the NGC 300 transient

was analogous to the outburst V838 Monocerotis (Bond et al. 2003).

Similar to the case of SN 2008S, Prieto (2008d) discovered a deeply-embedded

source in archival Spitzer imaging (PI: R. Kennicutt) at the location of the transient.

The source was detected in all IRAC bands, as well as in MIPS 24µm, and had

a [3.6] − [4.5] color of ≈ 2.7. The SED implies a ≈ 330 K blackbody (as with SN

2008S, a blackbody is a fairly poor fit to the SED) with a bolometric luminosity

of ≈ 5.6 × 104 L⊙. This finding confirms the massive stellar origin of the NGC 300

transient, and is consistent with a relatively low-mass massive star (see Fig. 4.5).

Importantly, depending on the details of stellar models for ZAMS masses in the

range of ∼ 10 M⊙, at fixed final bolometric luminosity, the inferred initial progenitor

mass may be multiply-valued and a luminosity of ≈ 5.6 × 104 L⊙ can imply a ∼ 5,

∼ 8 or ∼ 11 M⊙ progenitor (see Smartt et al. 2009, their Fig. 2; §4.5.2).4

The luminosity and blackbody temperature of the progenitor of NGC 300

suggest an obscuring medium with physical scale of order 300 AU. The deep limits

on the optical emission from the progenitor with HST suggest an optical depth at V

considerably larger than unity (∼ 8 − 10).5

4Prieto (2008d) assumed standard singly-valued masses, and given the measured luminosity of the

progenitor of NGC 300 quoted (∼ 105 L⊙) they inferred an initial progenitor mass of ∼ 15− 20 M⊙.5Assuming a spherical homogeneous medium and a Galactic dust-to-gas ratio, this optical depth

at V requires a total mass of obscuring material of roughly ∼ 2 × 10−3 M⊙ on a scale of 300 AU.

78

The fact that the transient in NGC 300 and SN 2008S were similar, both in

their luminosities (both relatively faint with respect to typical supernovae with

MV ≈ −14) and spectra (with narrow Balmer lines and strong [Ca II] 730 nm

doublet in emission; Bond et al. 2008), and that their progenitors were similar

(highly dust-obscured, relatively modest bolometric luminosities) led Prieto (2008d)

to propose that NGC 300 and SN 2008S share a common origin: the explosion —

whether supernova or outburst — of a massive star enshrouded in its own dust.

As we show below, the fact that their progenitors are so rare among massive stars

implies that just two events (although, see §4.2.3 & 4.2.4) are sufficient to define a

class.

4.2.3. M85

The optical transient in the Virgo galaxy M85 (NGC 4382) was discovered in

early 2006 by the Lick Observatory Supernova Search Team (KAIT; Li et al. 2001)

and discussed in Kulkarni et al. (2007) and Pastorello et al. (2007) and may also

be a member of the class defined by SN 2008S and NGC 300. The transient had

peak R-band absolute magnitude of ≈ −12 with a ∼ 80 − 100 day plateau, similar

to low-luminosity Type IIP supernovae. Optical limits constrain the progenitor to

be less than ≈ 7 M⊙ or highly obscured (Ofek et al. 2008; see also Pastorello et

al. 2007).6 The low optical luminosity, plateau, and redward spectral evolution of the

6Ofek et al. (2008) quote a limiting absolute magnitude at F850LP(z) with HST of > −6.2.

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transient led Kulkarni et al. (2007) to propose that it was analogous to the outburst

V838 Mon. In contrast, Pastorello et al. (2007) argued that it was a low-luminosity

Type IIP supernova (see §4.2.5).

Prieto et al. (2008e) again searched archival Spitzer imaging and discovered

an infrared source at the site of the optical transient, but, by chance, taken 8.8

days before the discovery in the optical by KAIT. The source was detected in all

IRAC bands and was associated with the transient itself, and not the progenitor,

since archival Spitzer imaging from 2004 (PID 3649; PI P. Cote) does not reveal a

point source at the location of the M85 transient. Using these images, we derive

3σ upper limits of 4 × 105 L⊙ and 2 × 105 L⊙ at 3.6µm and 4.5µm, respectively,

for the progenitor. The transient was also detected in the mid-IR seven months

after the initial discovery by Rau et al. (2007), the fluxes having decreased by a

factor of ≈ 5 over that time. The bright infrared transient discovered by Prieto

et al. (2008e) is adequately fit by a blackbody with temperature of ≈ 800 K with

luminosity ≈ 2 × 106 L⊙. The optical and NIR photometry of Kulkarni et al. (2007)

indicate a second component to the SED with a blackbody temperature of ≈ 3900 K

and with a luminosity of ≈ 5 × 106 L⊙.

The cooler re-radiated dust emission arises from a region of order 300− 400 AU

in physical scale. Assuming that the optical emission did not vary in the ≈8.8

days between the IR discovery and the optical discovery, the ratio of the power in

these two blackbody components implies that the optical depth at V (τV ) is less

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than unity. The physical scale of the obscuring medium indicates that it is likely

circumstellar, and the result of a mass-loaded wind. In addition, the luminosity of

the transient suggests that any pre-existing dust within ∼ 100 − 200 AU would have

been destroyed during the explosion. Given the fact that the optical depth to the

source scales as r−1 in a freely-expanding wind, it is not implausible that τV to the

progenitor was a factor of ∼ 10− 20 larger before explosion. These estimates suggest

a pre-explosion obscuring medium similar in its gross properties to the SN 2008S

and NGC 300 progenitors.

The M85 transient also showed narrow Balmer lines in emission, as well as

some Fe II lines, similar to SN 2008S and NGC 300. Because of the strong evidence

for obscuration of the progenitor, as evidenced by the bright IR transient, and the

similarity of the spectra, Prieto et al. (2008e) proposed that these outbursts share

a common origin and that their obscured progenitors may give rise a new class of

2008S-like transients.

We emphasize that because the character of the progenitor is not known (except

for the optical and IR limits), the connection to SN 2008S- and NGC 300-like events

is plausible rather than certain. Nevertheless, if the M85 transient was associated

with an embedded massive star, the IR limits we derive are consistent with the

luminosities derived for the 2008S and NGC 300 progenitors.

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4.2.4. SN 1999bw

The Lick Observatory Supernova Search reported in April 1999 the discovery

of a possible supernova in the galaxy NGC 3198 (Li 1999). The optical spectrum of

the transient, dominated by narrow Balmer lines in emission (Garnavich et al. 1999;

Filippenko et al. 1999), and its low V -band absolute magnitude at maximum of

≃ −13 (D ≈ 13.7 Mpc; Freedman et al. 2001) led Li et al. (2002) to propose that

this transient was an LBV-like outburst. Like SN 2008S and NGC 300, its spectrum

showed [Ca II] in emission. Additionally, an infrared source coincident with the

optical position of the transient was detected in archival Spitzer imaging obtained

with IRAC by the SINGS Legacy Survey five years after the discovery of SN 1999bw

(Sugerman et al. 2004). The source was detected in all IRAC bands and the SED

was well-fit by a 450 K blackbody with an integrated luminosity of ≈ 1.4 × 105 L⊙,

which translates into a blackbody scale of ∼ 300 AU.7 We have checked archival

IRAC images obtained in December 2005 (PID 20320; PI B. Sugerman), 1.5 yr after

the detection in the SINGS images, and we confirm that the MIR source is indeed

the transient, since the fluxes have declined by a factor of more than 3 during this

time.

The combination of a low optical luminosity at maximum, an optical spectrum

dominated by narrow Balmer lines in emission, the presence of [Ca II] emission, and

a luminous infrared emission detected with Spitzer, make SN 1999bw similar to SN

7The luminosity and blackbody scale have been adjusted to the distance employed here.

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2008S, NGC 300, and the transient in M85. However, as in the case of M85, we

emphasize that because there is no information on the progenitor we cannot be sure

that SN 1999bw was of the same class as SN 2008S and NGC 300.

4.2.5. The Connection to Other Transients

As we discuss in detail in §4.4, perhaps the primary distinguishing characteristic

of this class of transients is their deeply embedded progenitors. Since we are unable

to confirm the presence of such progenitors for the M85 transient or SN 1999bw, we

are unable to make a direct analogy with SN 2008S and NGC 300, and instead rely

on the fact that the transients themselves provides strong evidence for obscuration

on few-hundred AU scales.

In our effort to understand which cataloged transients might belong to the

SN 2008S/NGC 300-like class we have examined archival imaging of many recent

supernovae, as well as archetypal peculiar supernovae, including supernova impostors,

LBV outbursts, and low-luminosity Type IIP supernovae. Here, we provide a brief

discussion in an effort to orient the reader.

As we mentioned in §4.1 (point 2), the low-luminosity Type IIn SN 2002bu is

an interesting transient that may be a member of the class defined by SN 2008S

and NGC 300. We checked archival Spitzer data of the host galaxy NGC 4242 (PID

69; PI G. Fazio) taken 2 years after discovery and we find a bright infrared point

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source detected in all IRAC bands within 0.′′6 of the supernova position. Also, two

epochs of MIPS data (PID 40204; PI R. Kennicutt), obtained in 2008, 6 years after

explosion, and separated by just 6 days, reveal a 24µm source at the position of

the supernova. This is qualitatively similar to the case of M85 and SN 1999bw, but

because only a single post-explosion IRAC epoch exists, and because the two MIPS

epochs are separated by such a short time, we are unable to definitively confirm that

the MIR source is associated with SN 2002bu.

SN 1997bs is an intriguing example of an object that does not fit into this class,

although its peak absolute magnitude, lightcurve, color, and some spectral features

are comparable to SN 2008S and NGC 300. In this case, a luminous un-obscured

progenitor has been identified in the optical (MV ≈ −7) and the transient itself has

been argued to be the outburst of an LBV (van Dyk et al. 2003). Of interest is the

fact that no object has been subsequently identified in the optical at the site of the

transient (van Dyk 2005). We have checked archival Spitzer data of the host galaxy

(NGC 3627) obtained in 2004 (∼ 7 years after discovery) by SINGS and we do not

detect a bright MIR source at the site of SN 1997bs, in contrast with SN 2008S,

M85, and SN 1999bw. The event SN 2003gm is also interesting in this context

since it had photometric and spectroscopic evolution similar to SN 1997bs, and also

showed an optically luminous progenitor (MV ≈ −7.5; Maund et al. 2006). The fact

that both SN 1997bs and 2003gm had bright unobscured progenitors is our primary

reason for excluding them from the class defined by SN 2008S and NGC 300.

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Historical LBV eruptions in nearby galaxies that have been initially classified

as supernovae are also worth mentioning here. These include SN 1954J (e.g., Smith

et al. 2000) and SN 2002kg (e.g., Maund et al. 2006; van Dyk et al. 2006) in NGC

2403, SN 1961V (e.g., Humphreys 2005) in NGC 1058, and SN 2000ch in NGC 3432

(Wagner et al. 2004). As in the cases of SN 1997bs and SN 2003gm, a very important

common difference between these objects and SN 2008S or NGC 300 is that they all

had optically luminous progenitors with absolute magnitudes ∼< −6, consistent with

originating from very massive stars. These transients also have other properties that

are not consistent with SN 2008S-like explosions: (1) their peak absolute magnitudes

range between −18 ∼> MV ∼> −9, (2) the transient timescales vary widely from a few

days to years, and (3) they are not luminous MIR sources in archival Spitzer data.

Other optical transients classified as sub-luminous Type IIP supernovae that

might potentially fall into the class defined by SN 2008S and NGC 300 include SN

1994N, SN 1997D, SN 1999eu, 1999br, 2001dc, 2003Z, and 2005cs (e.g., Pastorello

et al. 2004, 2006). We note, however, that there are fundamental differences in the

spectra of low-luminosity Type IIP SNe compared with SN 2008S-like transients.

In particular, sub-luminous Type IIP SNe show Balmer lines with strong P-Cygni

absorption profiles and velocities of a few thousand km s−1, as is observed in more

luminous Type IIP SNe. This stands in sharp contrast with the Balmer lines in

SN 2008S-like transients, which are fairly narrow (FWHM ∼ 1000 km s−1) and

which do not show strong P-Cygni absorption features. In this way, SN 2008S-like

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transients most closely resemble the spectra of low-luminosity Type IIn SNe and

LBV outbursts.

In addition to the very interesting case of 2002bu, there were five other

low-luminosity transients classified as “impostors” or “unknown” that could have

been LBV outbursts, but for which no progenitor has been identified, and which

might be 2008S-like: NGC 4656, SN 2001ac, 2006bv, 2006fp, and 2007sv (see §4.3).

However, as we have emphasized and as the sources discussed above imply, the

properties of the progenitor cannot be deduced from the character of the optical

outburst alone (e.g., contrast SN 1997bs and SN 2003gm with SN 2008S). Thus,

in order to understand the causal mapping between progenitor and explosion, a

census of the progenitors must first be completed. What is clearly needed is a

comprehensive survey for bright MIR sources in all nearby galaxies (∼< 10 Mpc) with

(warm) Spitzer, analogous to the survey proposed by Kochanek et al. (2008) in the

optical. In the next section we discuss our search for deeply-embedded progenitors

in M33. We discuss a more complete census in §4.5.

4.3. Rates

The absolute rate of transients analogous to NGC 300 and SN 2008S is

uncertain. Current samples of supernovae over the last decade in the local volume

within 10, 20, and 30 Mpc allow us to make only a rough estimate of the true

rate. A systematic transient search in the local volume is crucial to solidify these

86

numbers. Nevertheless, we estimate that 2008S-like transients occur with a frequency

equivalent to ∼ 20% of the Type II supernova rate. We discuss the observed rates in

§4.3.1 below, and then we enumerate several arguments suggesting that the sample

of SN 2008S-like transients may be highly incomplete in the local universe (§4.3.2).

4.3.1. Observed Counts

D ∼< 10 Mpc

In addition to SN 2008S and NGC 300, ≈ 20 other core-collapse supernovae

or supernova-like transients have been discovered within ≈ 10 Mpc since 1999.8

Sixteen were confirmed supernovae; they are SN 1999em, 1999ev, 1999gi, 1999gq,

2002hh, 2002ap, 2003gd, 2004am, 2004dj, 2004et, 2005af, 2005at, 2005cs, 2007gr,

2008bk, and 2008ax. Two were bona fide LBV eruptions (SN 2000ch, Wagner et

al. 2004; 2002kg, Schwartz et al. 2003; Weis & Bomans 2005; Maund et al. 2006; van

Dyk et al. 2006). One was a Type IIn supernova potentially of the SN 2008S class

(SN 2002bu in NGC 4242, D ≈ 8 Mpc; Puckett & Gauthier 2002; Hornoch 2002).

Finally, the transient in NGC 4656 also had some of the spectral characteristics of

low-luminosity IIn supernovae (e.g., narrow Hα in emission), but only reached an

absolute magnitude of ≈ −11.5, (Rich et al. 2005; Elias-Rosa et al. 2005).

Taken at face value, with no correction for incompleteness, these numbers

imply that 2/22 ≈ 9% or 3/22 ≈ 14% (including SN 2002bu) of all optically bright

8Throughout the discussion here we exclude Type Ia supernovae.

87

transients are SN 2008S-like. Removing the two bona fide LBV outbursts (2000ch

and 2002kg) for comparison with the supernova sample proper, if SN 2008S and

NGC 300 are supernovae, they represent ≈ 10% and ≈15% (again, with SN 2002bu)

of the sample.

D ∼< 20 Mpc

A similar exercise can be carried out within the larger volume of ∼ 20 Mpc.

With a peak absolute magnitude of ∼ −14, SN 2008S-like transients would have

an apparent magnitude of 17.5 at D = 20 Mpc, without including a correction for

extinction intrinsic to the transient. In fact, SN 2008S and NGC 300 had AV ≈ 1.2

(Prieto et al. 2008c) and AV ∼ 0.3 − 1.2 (Bond et al. 2008), respectively.

Using the Smartt et al. (2009) compilation, we find 29 IIP, 4 IIb, 15 Ib/c, 2 IIn

(1998S & 2002bu), and 2 IIL supernovae in the last decade. There are 6 classified

as “LBV eruptions/impostors” (1999bw, 2000ch, 2002kg, 2003gm, 2007sv, and NGC

46569), but only 2000ch and 2002kg have strong evidence for an LBV progenitor.

Whether the yellow supergiant progenitor of 2003gm survived the explosion, and

hence whether 2003gm was in fact a supernova, has not yet been definitively

established (Maund et al. 2006). Similarly, although SN 2007sv, which reached

an absolute magnitude of ≈ −14, bears some similarity to 1997bs, the nature of

its progenitor has not been established (Duszanowicz et al. 2007; Harutyunyan et

9Note that NGC 4656 is not included in Smartt et al. (2009).

88

al. 2007). Finally, there are 3 events whose nature is classified in Smartt et al. (2009)

as unknown: M85, NGC 300, and SN 2008S.

With no correction for incompleteness, and taking only NGC 300 and SN 2008S,

this compilation implies an overall rate of 2/61 ≈ 3.3% within 20 Mpc. Including

M85 and 2002bu in the sample of 2008S analogs doubles the rate. Including 1999bw,

2007sv, and NGC 4656 brings the overall observed rate of SN 2008S-like transients

to ∼ 10% within D ∼< 20 Mpc.

D ∼< 30 Mpc

With a limiting magnitude of ≈ 18 − 19, amateur and professional supernova

surveys could in principle find SN 2008S-like transients with MV ≈ −14 to a

luminosity distance of ∼ 25 − 40 Mpc (again assuming no extinction). Objects of

interest discovered over the last ∼ 10 years with faint absolute magnitudes in this

distance range, and with the spectral characteristics of IIn supernovae analogous to

SN 2008S include 2001ac, 2006bv, and 2006fp.

A total of 92 core-collapse supernovae and 7 LBV eruptions appear in the

recent compilation of Smartt et al. (2009). Three events — M85, NGC 300, and

SN 2008S — are classified as “unknown.” Of the 7 “LBV eruptions,” only 2000ch,

2002kg, and 2003gm have LBV-like progenitors. The remainder — 1999bw, 2001ac,

2006fp, and 2007sv — have little or no progenitor information, spectra that resemble

IIn’s, and may be SN 2008S analogs. Given the number of confirmed NGC 300

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and SN 2008S-like analogs (just 2) and those suspected of belonging to this class

(M85, 1999bw, 2002bu), as well as those tradiationally labelled “LBV outbursts,”

but with no strong confirmation (4656, 2001ac, 2006bv, 2006fp, 2007sv) the overall

rate ranges from 2/102 ≈ 2% to 10/102 ≈ 10% when measured within 30 Mpc.

4.3.2. Arguments for Incompleteness & Some

Implications

To summarize the previous subsections, conservatively taking NGC 300 and SN

2008S as the only examples of their type ever observed (that is, excluding all other

low-luminosity transients), the observed rate is ∼ 9%, ∼ 3%, and ∼ 2% within 10,

20, and 30 Mpc volumes, respectively, with respect to all bright optical transients

when averaged over the last 10 years.

It is difficult to estimate the degree of uncertainty in these numbers since the

surveys that find local supernovae are a combination of professional and amateur,

with complicated and unquantitified selection functions for transient identification.

Most surveys responsible for transient discoveries in the local universe do not

have detailed descriptions of completeness in the literature. On the contrary, the

large majority of the transients in the local volume are discovered by amateurs.10

Nevertheless, it is possible to make an estimate of completeness that gives a sense of

how large the correction to the rate of SN 2008S-like transients may be.

10Of the 13, 17 and 13 local (D ∼< 30 Mpc) supernovae (including Type Ia’s) discovered in 2006,

2007, and 2008, only 1, 4 and 4 were discovered by professional searches.

90

As an example, B. Monard typically quotes ±0.2 magnitude (unfiltered)

photometric errors on discovery observations of SNe detected at ≈ 16−17 magnitude

(e.g., Monard 2006). This photometric error translates to a signal-to-noise ratio of

S/N ≈ 5. In order to estimate a lower limit on the incompleteness, we can compare

this value for S/N with the mean detection efficiency for Type Ia supernovae in the

SDSS-II Supernova Survey, which employs a well-tested photometric pipeline that

uses difference imaging to subtract off the host galaxy (see Dilday et al. 2008). Their

Figure 7 shows that for S/N ≈ 5, the detection efficiency is ≈ 0.5 in Sloan gri. The

detection efficiency drops to ∼ 0.1− 0.2 for S/N ≈ 2, corresponding to a photometric

error of ±0.4 magnitudes.

Momentarily ignoring the difference in the shape of Type Ia lightcurves with

respect to those of 2008S and NGC 300, the results of Dilday et al. (2008) imply that

surveys achieving limiting magnitudes of ≈ 18 − 19 are of order ∼ 10% complete for

SN 2008S-like transients with MV ≈ −14 at 25 − 40 Mpc (S/N ∼ 1 − 2). Thus, a

factor of ∼ 10 correction should be applied to the 30 Mpc sample in §4.3.1 from the

Smartt et al. (2009) catalog for the true rate of low-luminosity IIn supernovae like

SN 2008S. Within 20 Mpc the correction for incompletness is likely a factor of ∼ 5,

and within 10 Mpc the incompleteness is probably closer to a factor of ∼ 2. Similar

corrections should be applied to the observed rate of true LBV eruptions.

There is another argument for incompleteness at the factor of ∼ 2 level within

10 Mpc for SN 2008S-like transients. As summarized by Horiuchi et al. (2008) (their

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Section II.B), the observed rate of supernovae of all types within 30 Mpc yields a ratio

of Type Ia supernovae to core-collapse supernovae that is significantly in excess of the

cosmic ratio measured at high redshift (0.5 ∼< z ∼< 1; Dahlen et al. 2008). Indeed, the

ratio within 30 Mpc is large enough that we would expect to have seen several Type

Ia supernovae within 10 Mpc in the last 10 years. Yet, none have been found. This

fact implies that the ratio of Ia to core-collapse supernovae has been over-estimated

within 30 Mpc because normal core-collapse supernovae are intrinsically fainter than

Type Ia’s. Thus, the sample of normal core-collapse supernovae is incomplete within

30 Mpc at the factor of ∼ 1.5 − 2 level, even though these objects typically have

peak absolute visual magnitudes of MV ≈ −16 to −18. Naively, analogs to SN

2008S with MV ≈ −14 would be ∼ 6 times more incomplete than normal supernovae

at D ≈ 30 Mpc. Of course, the actual incompleteness correction depends on the

overall extinction correction for the transient population and on the cadence of the

observations since the lightcurve declines much faster for 2008S-like transients than

for supernova of type IIP. This implied nearly order-of-magnitude incompleteness

correction at D ≈ 30 Mpc for 2008S-like events strongly indicates that the sample is

incomplete at order unity within D ≈ 10 Mpc.

Additionally, a plot of the discovery rate of all supernovae (Ia’s included) within

30 Mpc over the last 10 years shows an increasing trend, super-Poisson variance, and

strong dependence on the results and observing strategy of a single survey (LOSS;

Li et al. 2000). There is also an asymmetry between the rate of discovery in the

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northern and the southern sky in excess of a simple extrapolation of star formation

from the catalog of Karachentsev et al. (2004). Finally, there is an unquantified

bias against small star-forming galaxies in the local universe. These points further

solidify the case that the normal core-collapse supernova rate is incomplete, which

implies that we are missing 2008S analogs in abundance in the local universe.

Taking yet another angle on the question of overall rate, we may consider the

a posteriori statistics of the events SN 2008S and NGC 300 themselves. Taking

the overall supernova rate as ∼ 1 − 2 yr−1 within 10 Mpc implies a probability of

∼ 4 − 8 × 10−4 of seeing two events in a single year if the overall rate of 2008S-like

transients is 2% of the supernova rate. We consider this uncomfortably small.

Similarly, if we were to take the true SN 2008S-like transient rate to be ∼ 100%

of local supernova rate, we would be forced to explain the fact that only ∼ 2 such

events have been seen within 10 Mpc in the last 10 years. Given the discussion of

incompleteness above, an overall rate of ∼ 20% with respect to supernovae gives a

reasonable chance of seeing two in one year and of seeing only a handful on a 10 year

baseline.

Taken together, these arguments imply that the sample of transients in the

local universe when averaged over the last 10 years is highly incomplete. We suggest

that the incompleteness correction is a factor of ∼ 2, ∼ 5, and ∼ 10 for MV ≈ −14

transients at 10 (≈ 16.0 mag), 20 (≈ 17.5 mag), and 30 Mpc (≈ 18.4 mag), even

before accounting for the potentially higher average extinction of these transients

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relative to normal supernovae. These estimates are consistent with the Richardson

et al. (2002), who conclude that low-luminosity supernovae with MB ∼> −15 may

constitute more than 20% of the overall supernova population (for related arguments,

see Schaefer 1996; Hatano, Fisher, & Branch 1997; Pastorello et al. 2004). There are

several immediate implications:

1. The true rate of SN 2008S-like transients is ∼ 20% of the core-collapse

supernova rate. However, we emphasize that lower and higher values at the

factor of ∼ 2 level are not excluded until a more thorough census has been

made.

2. The true rate of massive star eruptions (LBV-like and otherwise) is similarly

incomplete. The observed rate of “LBV eruptions” within 10 and 20 Mpc in

§4.3.1 and §4.3.1 implies that they are ∼ 1 − 3 times more common than SN

2008S-like transients. Thus, the true rate of massive LBV eruptions is likely

∼ 20 − 60% of the core-collapse supernova rate.11

3. The observed rate of core-collapse supernovae is incomplete at the factor of

∼ 2 level for D ∼< 30 Mpc.

11We thank the anonymous referee for pointing this out.

94

4.4. A First Census

Because of the implied frequency of events similar to SN 2008S and NGC 300

(§4.1, point 2) and the interesting character of their progenitors, we searched for

analogous sources in archival Spitzer imaging of nearby galaxies. Our goal was to

identify the underlying sub-population of massive stars from which these progenitors

emerge, to characterize their properties and frequency, and to catalog them for

future study.

The key characteristics of the progenitors of SN 2008S and NGC 300 are that

they are optically-obscured and deeply embedded, with very red MIR colors, that

their bolometric luminosities are indicative of relatively low-mass massive stars

(L ≈ 4 × 104 and ≈ 5.6 × 104 L⊙, respectively), and that the several epochs on

NGC 6946 revealed that the progenitor of SN 2008S was not highly variable in the

≈ 3 years before explosion (Prieto et al. 2008c). We discuss the variability of the

progenitor of NGC 300 in §4.4.4 below based on just two pre-explosion epochs.

Although we are only able to derive a lower limit to its RMS variation at 4.5µm,

like the progenitor of SN 2008S, we find that it too is consistent with a low level of

variability.

For a first census, we searched for the deepest archival Spitzer observations of

a nearby relatively massive bright star-forming galaxy, with already extant optical

catalogs. The Triangulum galaxy M33 is a perfect test case. It has an absolute

95

B-band magnitude of MB ≈ −19.2, a distance of ≈ 0.96 Mpc (distance modulus

µ = 24.92; Bonanos et al. 2006), and it has extensive optical (e.g., Hartman et

al. 2006; Massey et al. 2006), Hα (Massey et al. 2007), and MIR and FIR imaging

(McQuinn et al. 2007). A similarly rich dataset exists for several other local galaxies,

including the Magellanic clouds (e.g., Blum et al. 2006; Bolatto et al. 2007) and M31

(e.g., Mould et al. 2008). An analysis similar to that described below will be the

subject of future work (but, see §4.4.5).

In this Section, we present the MIR color-magnitude and color-color diagrams

for all cataloged point sources in M33 obtained from multi-epoch archival

Spitzer/IRAC (Fazio et al. 2004) observations (PI R. Gehrz; PID 5). This dataset

allows us to search for and identify stars analogous to the progenitors of SN 2008S

and NGC 300. We also present a variability study for both the reddest sources we

find in M33 (all extreme-AGB “EAGB” stars) and for the optically-selected LBV

candidates from Massey et al. (2007), which are detected in the MIR imaging. We

find that the population of the reddest sources — those most likely to be true

analogs of the progenitors of SN 2008S and NGC 300 — is completely distinct from

the population of LBV candidates in the primary metrics of color, magnitude, and

variability. Indeed, we find very few sources with the properties of the SN 2008S and

NGC 300 progenitors.

Section 4.4.1 describes our procedure for extracting point sources and the

resulting catalog. In §4.4.2 we present the color-magnitude diagram for M33 and

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we discuss the reddest sources vis a vis the optically-selected LBV candidates from

Massey et al. (2007). In §4.4.3 and §4.4.4 we discuss their SEDs and variability,

respectively. Finally, in §4.4.5 we discuss a preliminary search for similar sources in

NGC 300, the LMC, and the SMC.

4.4.1. Catalog

We coadded six epochs of mid-IR imaging of M33 obtained between Jan. 9,

2004 and Feb. 4, 2006 with IRAC (3.6-8.0µm; see McQuinn et al. 2007 for details

of the observing program). We produced the coadds in all four IRAC channels

from the flux calibrated mosaics provided by the Spitzer Science Center (post-BCD

data). Our final mosaics cover an area of ∼ 33′ × 33′ (1650 × 1650 pixels, with

1.′′2/pix) centered on M33, this is approximately within R25 (≃ 35′; de Vaucouleurs et

al. 1991). We performed source detection and PSF fitting photometry on the coadds

using the DAOPHOT/ALLSTAR package (Stetson 1992). The PSF magnitudes

obtained with ALLSTAR were transformed to Vega-calibrated magnitudes using

simple zero point shifts derived from aperture photometry (using a 12′′ radius),

performed in the original images, of 10 − 20 bright and isolated stars in each

band. We estimate errors in the photometric transformations to the Vega system of

0.04 mag in 3.6 µm, 0.05 mag in 4.5 µm, 0.07 mag in 5.8 µm, and 0.07 mag in 8.0 µm

channel. For the 3.6 and 4.5µm bands, the detection limits (3σ) in the coadds are

[3.6] ≃ 18.9 mag and [4.5] = 18.2 mag, respectively. The sample becomes incomplete

97

at [3.6] ≃ [4.5] > 17.1, approximately 0.5 magnitudes deeper than McQuinn et

al. (2007). A total of ≈ 80, 000 sources are detected in either 3.6 or 4.5µm. We

cross-matched the two catalogs using a 0.5 pixel (0.′′6) matching radius to obtain a

final catalog with ≈ 53, 200 individual sources detected at both 3.6 and 4.5µm (see

Table 4.1).

There are several reasons for producing new point source and variability (see

§4.4.4) catalogs, given the already existing catalog from McQuinn et al. (2007): (1)

because we are specifically looking for objects with extreme colors, we wanted to

be able to relax the criterion for point source detection in all IRAC bands (in the

Figures that follow, all sources of interest are detected with 3σ confidence); (2) we

wanted to be able to derive our own upper limits in each band for the same reason;

(3) we wanted full IRAC SEDs, whereas the catalog of McQuinn et al. (2007) does

not provide data at 5.8µm; (4) we wanted to combine the images used by McQuinn

et al. (2007) with a sixth archival epoch (PI R. Gehrz; PID 5); (5) for the reddest

stars, we wanted to derive full six-epoch lightcurves for a more complete measure of

variability. The resulting MIR color-magnitude diagram looks different from that in

McQuinn et al. (2007). Of the 18 sources we discuss extensively below, only 2 were

detected at both 3.6 and 4.5 µm in the catalog of McQuinn et al. (2007).

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4.4.2. The Color-Magnitude Diagram

The primary result of this effort on M33 was the production of the MIR

color-magnitude diagram (CMD), shown in Figure 4.1, which shows the [3.6] − [4.5]

color for all the sources detected at both 3.6 and 4.5µm as a function of absolute

magnitude at 4.5µm, M4.5. The dashed line marks the 3σ completeness limit in this

plane. The Spitzer colors and magnitudes of the progenitors of SN 2008S and NGC

300 are shown for comparison (filled triangle and filled square, respectively; Prieto et

al. 2008c, Prieto 2008d). Note that there are remarkably few objects inhabiting the

bright and (very) red region of the CMD. Among the ∼ 5× 104 massive stars in M33

(see §4.5), ≈ 18, 186, and 567 have both M4.5 < −10 and [3.6] − [4.5] color larger

than 1.5, 1.0, and 0.7, respectively, which correspond to blackbody temperatures of

≈ 500, 700, and 1000 K, respectively. A total of 2264 point sources are detected with

[3.6] − [4.5] ≤ 0.7 and M4.5 < −10.

Figure 4.2 shows an expanded view of the brightest MIR sources. In addition

to the progenitors of SN 2008S and NGC 300, we include several well-studied LBVs

(e.g., η-Carina), and a number of cool hyper-giants (VY CMa, NML Cyg, and

Var A in M33) for comparison. The MIR magnitudes and colors for η-Carina,

VY CMa, and NML Cyg were synthesized from ISO spectra (Sloan et al. 2003).

The magnitudes of M33 Var A were obtained from Humphreys et al. (2006). Also

included are the 16 sources matched between the catalog of MIR sources presented

99

here and the LBV sample of Massey et al. (2007), obtained from narrow-band Hα

imaging, using a 0.′′6 matching radius. The larger circles within the dotted lines show

the EAGB stars, which we discuss in detail below.

The primary point of Figures 4.1 and 4.2 is to show that there are very few

massive stars with the colors and MIR luminosities of the progenitors of SN 2008S

and NGC 300. Although it is difficult to identify a quantitative criterion for inclusion

in the class defined by the progenitors of SN 2008S and NGC 300, it is clear from

Figure 4.1 that the number of analogs in color and magnitude is very small with

respect to the total number of massive stars in M33. For example, if to be included

as an analog to the SN 2008S and NGC 300 progenitors we require that M4.5 be

brighter than or equal to the progenitor of NGC 300 and we require that the color

be redder than the lower limit on the progenitor of SN 2008S, we find a single

source. If we require M4.5 < −10 and color redder than SN 2008S, we find just two

sources. Casting the net more widely, for the purpose of having a sample larger

than one or two objects and in an effort to be conservative, we use M4.5 < −10 and

[3.6] − [4.5] > 1.5 to identify a sample of 18 EAGB stars (large open circles). We

discuss the spectra and variability properties of these sources in §4.4.3 and §4.4.4.

Our choice of the cuts M4.5 < −10 and [3.6] − [4.5] > 1.5 to identify objects

of interest is somewhat arbitrary. Because our argument in this paper relies on the

fact that analogs to the progenitors of SN 2008S and NGC 300 are intrinsically rare

(§4.1 & §4.5), this issue deserves discussion. The magnitude limit is straightforward:

100

it is meant to select objects that have bolometric luminosities indicative of massive

stars (∼> 8 M⊙). In §4.4.3 below, we show that M4.5 < −10 is conservative; only half

of the 18 sources selected have bolometric luminosities large enough to be massive

stars. Note that had we required M4.5 to be brighter than or equal to the NGC 300

progenitor, we would exclude 8 of our 18 sources (see Fig. 4.5). Our goal was to not

miss any deeply-embedded massive stars and the criterion M4.5 < −10 accomplishes

that goal.

The cut on color is more complicated. We were motivated by several factors.

First, we wanted to avoid the AGB sequence blueward of 1.5, where the density of

points increases dramatically and where the sample would consist largely of Carbon

stars (see Figs. 4.1 & 4.2). Second, this cut essentially eliminates contamination

from background active galaxies (Stern et al. 2005). Third, the [3.6] − [4.5] > 1.5

cut gives us a reasonable number of objects to assess individually — it is neither too

many, nor too few (again, had we taken [3.6] − [4.5] > 2.0 the sample would consist

of just one or two objects).

These considerations leave us open to the potential criticism that optically-

obscured massive stars may exist in the region M4.5 < −11.5 (again, avoiding the

AGB feature in the CMD) and with 0.5 < [3.6]− [4.5] < 1.5. To address this, we have

examined the 45 sources that occupy this region of the CMD. Four are identified

with the LBV candidate catalog from Massey et al. (2007), which we discuss more

below. Sixteen of the remaining 41 sources have bright optical counterparts from

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the catalog of Massey et al. (2006). Of the 25 sources that do not appear in Massey

et al. (2006), 16 are optically-detected, but at flux levels where the catalog is highly

incomplete. That leaves 9 optically-obscured sources that have MIR luminosities

indicative of massive stars. Like the very brightest of the 18 sources within the

dotted lines in Figure 4.2, about half of these have very large MIR luminosities,

considerably larger than the progenitors of SN 2008S and NGC 300. We conclude

that in this region there is just a handful of sources that might be true analogs to

these progenitors — and we note that these have colors ∼ 0.5 − 1.0 magnitudes

bluer than the lower limit on the progenitor of SN 2008S. Indeed, self-obscuration

to the extent of the NGC 300 and SN 2008S progenitors is exceedingly rare for the

most luminous stars, as evidenced by the lack of objects in the upper right corner of

Figures 4.1 and 4.2.

Finally, we note that we have searched for 4.5µm sources without 3.6µm

detections that would lie within the dotted lines in Figure 4.2 and we find just one

source. Close inspection of the images reveals a marginal 3.6µm detection and

[3.6] − [4.5] ≈ 1.5.

Despite this long discussion of color and magnitude selection, the primary point

of Figures 4.1 and 4.2 still stands: there are remarkably few massive stars in M33

that have the color and luminosity of the progenitors of SN 2008S and NGC 300. As

we discuss in §4.4.3, all are consistent with relatively low-mass massive stars.

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The second point to note from Figures 4.1 and 4.2 is that the EAGB stars we

have selected are not optically-luminous LBVs, η-Carina analogs, or cool hyper-

giants. Indeed, all of the LBV candidates (open triangles) have [3.6] − [4.5] ∼< 0.8,

and about half have colors ∼< 0.3. The cool hyper-giants and η-Carina-like objects

are also bluer than the EAGB population, and considerably brighter than the

progenitors of SN 2008S and NGC 300. Indeed, the latter are most naturally

associated in this diagram with the luminous red extremum of the AGB population,

hence our use of “extreme-AGB” (EAGB) stars.

The color-color diagram for all the sources detected in the four IRAC bands

is shown in Figure 4.4. The symbols are the same as in Figure 4.2. The small

points with extremely red [5.8] − [8.0] colors are relatively dim, with M4.5 > −10

and are likely young stellar objects (YSOs; e.g., Bolatto et al. 2007). The strong

deviation of the SN 2008S and NGC 300 progenitors from the blackbody curve

(solid line) reinforces the fact that the SEDs of these sources are not well-fit by a

simple blackbody (Prieto et al. 2008c). Despite this, the SN 2008S and NGC 300

progenitors, EAGBs, LBV candidates, and cool-hypergiants do not stand out as

separate populations in [5.8] − [8.0] color. This is important because it means that

[5.8]− [8.0] color alone cannot be used as a metric for inclusion or exclusion from the

class of SN 2008S/NGC300-like progenitors.

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4.4.3. Spectral Energy Distributions

The limits on the optical emission from the progenitors of NGC 300 and SN

2008S are tight, and effectively rule out an optically unobscured massive star (Berger

& Soderberg 2008; Prieto et al. 2008c; Fig. 4.5). For the purposes of finding analogs

to these sources, it is critical to derive the optical luminosities and/or upper limits

for the 18 EAGBs identified in §4.4.2. Here, we present these results and we compare

the derived optical-to-MIR SEDs of the EAGBs to the optically-bright (narrow band

Hα-selected) LBV candidates from Massey et al. (2007).

As part of the Survey of Local Group Galaxies Currently Forming Stars

(SLGG), Massey et al. (2006) presented a catalog of ∼ 150, 000 point sources in M33

with well-calibrated UBV RI photometry obtained from observations at the KPNO

4m telescope with the Mosaic imager. We use the published photometric catalog

and images from Massey et al. (2006) to complement the Spitzer MIR photometry

described in §4.4.2.

We first cross-matched the positions of the EAGB stars with the Massey et

al. (2006) photometric catalog. Importantly, we do not find any optical counterpart

to the EAGB sources using a matching radius of 0.′′5. Since the completeness of the

catalog starts to decline rapidly at V ≃ 22 mag, we analyze the images independently

to look for faint optical counterparts to the EAGB stars. We use SEXtractor (Bertin

& Arnouts 1996) with a low detection threshold (2σ above the local background)

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to detect and measure aperture photometry (using a small aperture of 3 pixels

radius) of all the sources detected in the KPNO/Mosaic UBV RI images of M33

from Massey et al. (2006). We calibrate the photometry relative to the magnitudes

in the catalog of Massey et al. (2006). Using a radius of 0.′′5 to cross-match the MIR

positions of the EAGB stars with our multi-band catalogs of faint optical sources,

we only detect two EAGB stars in the BV R bands (again, 2σ). The remaining

17 sources do not have optical counterparts. We estimate 3σ upper limits on the

UBV RI magnitudes using the local background RMS at the positions of the EAGB

stars. The median 3σ upper limits are: U = 24.0, B = 24.0, V = 23.5, R = 23.0,

I = 22.5. The data for each of the 18 EAGB stars is listed in Table 4.2.

In order to convert MIR and optical magnitudes to fluxes, we used zeropoints

in Reach et al. (2005) and Cohen et al. (2003) for the EAGBs and LBV candidates,

respectively. The luminosities of all sources were calculated assuming a constant

reddening of E(B−V ) = 0.15 mag and a distance of µ = 24.92 (Bonanos et al. 2006).

The reddening correction is motivated by the uncorrected B − V color-magnitude

diagram of Massey et al. (2006), which shows that the bluest sources only reach

B − V ≈ −0.2, instead of ≈ −0.33, as would be expected from an un-reddened

massive star. In addition, Bonanos et al. (2006) also quote an average reddening

correction of 0.1 mag to massive stars for M33. A larger adopted reddening

correction increases our upper-limits for the optical fluxes of the EAGB stars.

105

The primary result of this procedure is Figure 4.5, which shows the SEDs of

the EAGB stars (left panel) and LBV candidates (right panel). The dotted lines in

the left panel show the range of upper limits (and one BV R detection, as described

above) for the 18 EAGB stars at UBV RI (see Table 4.2). The filled triangles show

the optical and 3.6µm upper limits for the progenitor of SN 2008S. The solid squares

show the MIR detections and optical upper limits for the NGC 300 progenitor.

These sources should be contrasted with the optically-luminous LBV candidates

from Massey et al. 2007 (right panel). Note that 3 of the LBVs do not have 5.6 and

8.0µm detections. The fact that one of the non-detections would appear to have a

5.6 and 8.0µm flux larger than some of the other detections is a consequence of the

locally higher MIR diffuse flux near that particular object.

There are a number of points to take away from the two panels of Figure 4.5.

First, 9 of the 18 EAGB stars we identified in Figure 4.1 do not have bolometric

luminosities indicative of massive stars; they have Lbol ∼< 2× 104 L⊙. Thus, these are

not likely to be true analogs to the SN 2008S and NGC 300 progenitors. Second, all

of the sources are highly optically-obscured, with λLλ[V ]/λLλ[4.5 µm] ∼ 10−2. Third,

these sources are qualitatively different from the more bolometrically luminous LBV

candidates (right panel). The LBVs are interesting in their own right, dividing

approximately into two classes: (1) relatively optically-dim with a strong MIR

excess and (2) optically-bright with little MIR excess, if at all. This division is also

106

evidenced by their positions in the CMD (Fig. 4.1), which indicates a bimodality in

MIR color. For a possible analog to LBVs with a MIR excess, see Smith (2007).

4.4.4. Variability

Because four epochs of archival Spitzer data were available for NGC 6946 in

the three years prior to the discovery of SN 2008S, Prieto et al. (2008c) investigated

potential variability of the progenitor. They found that there was remarkably little

and showed that this fact could be used to constrain the motion of the obscuring

medium, under the assumption of a geometrically-thin, but optically-thick shell (see

§4.2; Prieto et al. 2008c).

Motivated by this result, and by the fact that six epochs of archival data over

two years exist for M33, we investigated the MIR variability of all the sources

detected at 3.6 and 4.5µm. To generate lightcurves, we used the difference imaging

analysis package ISIS, based on the techniques of Alard & Lupton (1998), Alard

(2000).12 For a discussion, see Hartman et al. (2004).

There is a striking difference between the 4.5µm light curves of the 18 EAGB

stars and the LBV candidates from Massey et al. (2007). Most of the EAGB stars

are highly variable, both in magnitude (0.1 ∼< RMS(mag) ∼< 0.8) and in [3.6] − [4.5]

color. In contrast, the majority of the LBVs are not variable (RMS ∼< 0.1 mag); only

two sources show clear variability (both are blue in [3.6] − [4.5]). To illustrate these

12 See http://www2.iap.fr/users/alard/package.html.

107

differences, we show in Figure 4.6 the light curve at 4.5 µm and color variations

of the reddest source in our EAGB sample S1 (left panel) and an LBV candidate

(right panel). For completeness, we present all of the 4.5 µm light curves and color

variations of the 18 EAGB stars in Appendix A (Figures A.1 & A.2), and of the

LBVs in Appendix B (Figures B.1 & B.2).

Figure 4.7 summarizes these findings. It shows the measured RMS at 4.5 µm

as a function of [3.6] − [4.5] color for all the bright sources with M4.5 < −10. The

symbols are the same as in Figures 4.2 and 4.4. There is a clear correlation evident

between the RMS (or amplitude) and color for the AGB stars (see also McQuinn et

al. 2007).

For comparison, the RMS variation of the progenitor of SN 2008S derived from

its 3-year light curve is also shown. In the case of the NGC 300 progenitor we can

only put a lower limit on its RMS variation, because only two epochs of archival

Spitzer imaging exist. Nevertheless, it is striking that both the SN 2008S and

the NGC 300 progenitors are consistent with very little variation in the few years

preceding their explosions. In particular, SN 2008S is inconsistent with the clear

trend among the AGB stars to become more variable as they become redder. Only

a handful of the EAGB stars vary so little, which suggests that a lack of variability

among an otherwise variable EAGB star population may be used as a selection

criterion for analogs to the SN 2008S and NGC 300 progenitors. As an example,

requiring the RMS to be ∼< 0.3 magnitudes, we find 5 sources. They are S12, S7,

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S9, S17, and S8, in order of increasing RMS (see Table 4.2). Sources S8 and S9 are

among the lower luminosity sources in the left panel of Figure 4.5, and are therefore

not likely true SN 2008S analogs. In contrast, S7 is the brightest of our EAGB stars.

Finally, the least variable source (S12) has MV ≈ −11.6, which is quite close to the

SN 2008S progenitor, even though it is ∼ 0.5 magnitudes bluer.

Although we have identified a few rare and interesting sources, the sparsity of

data in Figures 4.2 and 4.7 with [3.6] − [4.5] > 1.5 makes it difficult to construct a

strict quantitative joint criterion in the space of luminosity, color, and variability for

inclusion in the class of SN 2008S-like progenitors. Based on the discussion above,

as in §4.4.2 and §4.4.3, we expect a total of ∼ 1 − 10 in M33. A more complete

multi-epoch survey of EAGB stars in the local universe may fill in the region in

Figure 4.7 between the AGB locus and the SN 2008S and NGC 300 progenitors.

Most importantly, it might make clear a quantitative criterion for “SN 2008S-like”

in the RMS-color plane.

4.4.5. Other Galaxies

Blum et al. (2006) and Bolatto et al. (2007) present Spitzer point source

catalogs for the LMC and SMC, respectively. We have searched these catalogs for

sources that satisfy the selection criteria M4.5 < −10 and [3.6] − [4.5] > 1.5 used

to identify the 18 EAGB stars discussed throughout this section. In the catalog of

Blum et al. (2006), we find 9 sources. Three are coincident with 2MASS sources and

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5 appear in the IRAS catalog. Although more careful follow-up is clearly required, a

subset of these 9 sources may be EAGB stars. In the catalog of Bolatto et al. (2007)

for the SMC, we find a single source that satisfies M4.5 < −10 and [3.6]− [4.5] > 1.5.

Finally, we have also completed a cursory search for EAGB stars in archival imaging

of NGC 300 (PI R. Kennicutt; ID 40204), and we find just 4 potential sources.

In sum, even the relatively conservative criteria M4.5 < −10 and [3.6]−[4.5] > 1.5

pick out remarkably few stars in any galaxy — M33 is not peculiar in this regard.

Given the fact that we expect only a fraction of the 18 sources in M33 to be bona

fide analogs to the progenitors of SN 2008S and the transient in NGC 300 (based

on luminosity, color, and variability; see §4.4.3, §4.4.4), a more careful look at the

sources of interest in the LMC (≈ 9), SMC (≈ 1), and NGC 300 (≈ 4) is likely

to further decrease the total number of sources of interest in these systems. As

emphasized in §4.1 and in §4.5 below, the scarcity of SN 2008S-like progenitors with

respect to the total massive star population is remarkable in light of the fact that SN

2008S-like transients are likely to be relatively common with respect to the overall

supernova rate.

4.5. Discussion

We have shown that in the primary metrics of color, luminosity, and variability,

stars analogous to the progenitors of SN 2008S and NGC 300 are exceedingly rare

in star-forming galaxies. They have luminosities characteristic of low-mass massive

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stars, are deeply dust-obscured with extremely red MIR colors, and show little MIR

variability (see Figs. 4.1, 4.2, 4.5, and 4.7). In luminosity and color they are distinct

from the population of optically-luminous LBV candidates selected from Massey et

al. (2007). Although many of the reddest objects selected as EAGB stars in Figure

4.2 are highly variable, the few (∼ 1− 5) least-variable sources most closely resemble

the SN 2008S and NGC 300 progenitors. In this way (but in only this way), they are

similar to the LBV candidates.

In this section we discuss the implications of our finding that stars with

characteristics analogous to the progenitors of SN 2008S and NGC 300 are rare. In

§4.5.1 we estimate the overall fraction of massive stars that are deeply dust-embedded

and the lifetime of stars in that state. In §4.5.2, we connect with the evolution of

massive stars, including the possibility that SN 2008S-like transients are the result

of electron-capture supernovae or massive white dwarf birth. Section 4.5.3 discusses

how many EAGB stars can be found within the local universe (D ∼< 10 Mpc) using

Spitzer.

4.5.1. Numbers & Rates

The total number of analogs to the progenitors of SN 2008S and NGC 300

in M33 is uncertain. This uncertainty comes primarily from the fact that we are

unable to identify an absolute quantitative criterion for inclusion into this progenitor

class based on our three primary metrics: color, luminosity, and variability. To

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be conservative, we have identified 18 sources in the region M4.5 < −10 and

[3.6] − [4.5] > 1.5 of the CMD that satisfy the minimal criteria of being bright and

extremely red (larger open circles in Fig. 4.2). However, a very strict cut in color

and magnitude (i.e., all sources redder than the lower limit to SN 2008S and brighter

than NGC 300) yields just two sources. Among the 18 selected sources, we have

shown that roughly half do not have bolometric luminosities indicative of massive

stars. That is, they do not have luminosities as large as one would expect for stars

who are traditionally thought to end their lives as supernovae (Lbol ∼> 4 × 104 L⊙).

It is important to note, however, that at fixed final luminosity of a massive star its

initial ZAMS stellar mass may be multi-valued, implying progenitors with either

∼ 5 − 7, ∼ 8 − 9, and ∼ 11 − 14 M⊙ for L ≈ 5 − 10 × 104 L⊙. (see Fig. 2 of Smartt

et al. 2009; §4.5.2; see footnote 4). Of course, the explosions SN 2008S and NGC

300 may not be true supernovae, but rather a new class of bright eruptions from

obscured massive stars (see §4.5.2). Nevertheless, comparing these 18 sources to

the progenitors of SN 2008S and NGC 300 in Figure 4.5 we would argue that only

roughly half belong to this progenitor class based on SED alone. Finally, Figure

4.7 shows that only ≈ 1 − 5 of the 18 sources vary as little as the progenitors of

SN 2008S and (potentially) NGC 300. Importantly, 16 of the 18 sources satisfy the

criterion of being highly optically-obscured (see Table 4.2).

In summary, very few massive stars have the color, luminosity, and variability

of the SN 2008S and NGC 300 progenitors. Our best guess is that the number of

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true analogs may be as few as zero and as large as ∼ 10−20. We denote this number

in M33 — the number of true analogs — as NEAGB. A larger sample of stars, culled

from a larger multi-epoch study of local star-forming galaxies (see §4.5.3) is clearly

needed to fill in the parameter space in the extreme red and bright side of the CMD.

This is the most robust way to understand NEAGB and its uncertainty.

In order to evaluate the fraction of stars in M33 that might be analogs to the

progenitors of SN 2008S and NGC 300, and so constrain the rate of production

of such objects, we must first estimate the total number of massive stars in M33

(N⋆; i.e., with MZAMS ∼> 8 M⊙). This number can be estimated in several ways: (1)

extinction-corrected Hα luminosity (e.g., Hoopes et al. 2001; Hoopes & Walterbos

2000; Greenawalt 1998), (2) dust-reddening corrected UV continuum luminosity (for

GALEX observations, see Thilker et al. 2005), (3) total number of main-sequence

optical point sources detected with MV ∼< −2 (appropriate for stars with ZAMS

masses above ≈ 9 − 10 M⊙; Lejeune & Schaerer 2001), or (4) total number of red

supergiants (RSGs) (MV ∼< −3.5) times the ratio of the lifetime of a massive star

to the time spent as an RSG (t⋆/tRSG ≈ 10; e.g., Schaller et al. 1992). Using the

latter method, and selecting RSGs with V − R > 0.5 and MV < −3.5 from the

catalog of Massey et al. (2006), we find ≈ 5400 sources, implying N⋆ ≈ 5.4 × 104.

Taking a more conservative color cut of V − R > 0.7 and MV < −3.5, we find that

N⋆ ≈ 3.5× 104. Similar estimates in the range of N⋆ ≈ 3− 6× 104 are obtained using

method (3) with the Massey et al. (2006) catalog, although this estimate suffers

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significantly from incompleteness. We take N⋆ = 5 × 104 as a fiducial number and

include it in our scalings below. Note that estimates of the total star formation rate

in M33 range from ∼ 0.3 to ∼ 0.7 M⊙ yr−1, consistent with the UV, Hα, and FIR

luminosities (e.g., Gardan et al. 2007 and references therein), implying a supernova

rate for the galaxy of ∼ 0.005 yr−1 (e.g., Gordon et al. 1998).

Taking NEAGB ∼ 5 and N⋆ ≈ 5 × 104, we find that a fraction

fEAGB =NEAGB

N⋆

∼ 1 × 10−4[

NEAGB

5

]

[

5 × 104

N⋆

]

(4.1)

of the massive stars in M33 may be analogs to the progenitors of SN 2008S and NGC

300.

As noted in §4.1 (point 2) and §4.3, only a fraction of all massive stars go

through this highly dust-enshrouded phase, and produce transients like SN 2008S

and NGC 300. Since, by assumption, roughly all of the massive stars in any galaxy

become normal core-collapse supernovae (but, see Kochanek et al. 2008), the rate of

SN 2008S-like explosions can be characterized by their fractional rate with respect to

the overall supernova rate. This fraction is determined by dividing the observed rate

of SN 2008S-like transients by the total number of supernovae within some volume,

times an incompleteness correction that accounts for the fact that SN 2008S-like

transients are intrinsically less optically luminous. Based on the numbers presented

in §4.3, we estimate that fSN ≈ 0.2, although higher and lower values are not

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excluded. For example, it is possible that SN 2008S-like transients are intrinsically

rare and the fact that NGC 300 and SN 2008S occurred in the same year was

simply chance. Although we cannot exclude this possibility, we note that such an

explanation appears improbable in the face of what is known about the rarity of

their progenitors (fEAGB; eq. 4.1). Conversely, it is possible that the incompleteness

correction exceeds the factor of ∼ 2 advocated in §4.3 within 10 Mpc and that such

transients are indeed common with respect to supernovae. However, it then becomes

increasingly difficult to explain why no more SN 2008S-like transients were observed

within 10 Mpc in the last 10 years. There is no way to circumvent these uncertainties

without a more complete census of progenitors and outbursts.

As stated in §4.1, the simplest explanation for the fact that SN 2008S- and

NGC 300-like transients are simultaneously common with respect to supernovae

(fSN ∼ 0.2) and that their progenitors are very rare by number at any moment,

in any star-forming galaxy (fEAGB ∼ 10−4) with respect to massive stars, is that

a significant fraction of all massive stars (∼ 0.2) go through a brief evolutionary

epoch in which they are highly dust-obscured, just before explosion. Taking the

average lifetime of massive stars with ZAMS masses in the range of 9 − 10 M⊙ to

be t⋆ ≈ 3 × 107 yr (e.g., Schaller et al. 1992), we find that the duration of this

dust-obscured phase is

tEAGB ∼ 1 × 104

[

t⋆107.5 yr

]

[

NEAGB

5

]

[

5 × 104

N⋆

] [

0.20

fSN

]

yr. (4.2)

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We consider the uncertainty in fSN to be at the factor of two level and the uncertainty

in N⋆ to be at the level of a factor of 1.5. However, as we have stressed, NEAGB may

be as much as a factor of 5 or more lower (tEAGB ∼< 103 yr), or a factor of ∼ 2 − 4

higher (tEAGB ∼ 6 × 104 yr). To improve these numbers significantly, a careful

monitoring program for optical transients like SN 2008S within the local universe

(D ∼< 10 Mpc), coupled with a survey of all local galaxies for bright MIR point

sources with (warm) Spitzer (see §4.5.3) should be undertaken. The combination

of watching for more transients of this type and associating them with individual

progenitors whose luminosities and variability have been cataloged will significantly

decrease the uncertainty in both fSN and NEAGB, and significantly increase our

understanding of the causal mapping between progenitors and their outbursts.

4.5.2. Connection to The Evolution of Massive Stars

The relation between final luminosity and initial stellar mass may be

triply-valued (Smartt et al. 2009) at ∼ 5 − 7, ∼ 8 − 9, and ∼ 11 − 14 M⊙ for

L ≈ 5 − 10 × 104 L⊙. This relation is of course uncertain, particularly in the

mass range singled out by the bolometric luminosity of the 2008S and NGC 300

progenitors near ∼ 10 M⊙. It is likely further complicated by binarity, and by the

mass-loss history, metallicity, and rotation of massive stars. Because the absolute

rate of these outbursts as well as whether or not they should be associated with

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the death of the progenitor are still uncertain, we consider a number of potential

scenarios below. We list a subset of the possibilties in order of increasing progenitor

mass.

Massive White Dwarf Birth: M ≈ 6 − 8 M⊙

We have referred to the progenitors of SN 2008S and NGC 300 throughout

this work as extreme (“E”) AGB stars because they lie at the red extremum of

the AGB sequence in the MIR color-magnitude diagram (Figs. 4.1 & 4.2). Taken

literally, these stars may indeed be the progenitors of the most massive O-Ne-Mg

white dwarfs, undergoing explosive core-envelope separation as they transition to

proto-planetary nebulae (e.g., Riera et al. 1995; Garcıa-Hernandez et al. 2007).

Perhaps the 2008S and NGC 300 progenitors were then akin to the most massive

highly-evolved carbon- or oxygen-rich AGB stars (Kwok 1993).

Based on analogy with local proto-planetary nebulae, we would expect bi-polar

explosion morphology and eventually the emergence of a hot ionizing continuum

source as the newly-born white dwarf begins its cooling phase (perhaps similar

to Hen 3-1475/IRAS 17423-1755; Riera et al. 2003). The initial luminosity of

the central source would be of order ∼ 5 × 104 L⊙ for a white dwarf near the

Chandrasekhar mass and it should cool on a timescale comparable to ∼ 105 yr.

Thus, the bolometric luminosity of the transient should eventually decrease back to

approximately pre-outburst levels. The primary distinguishing characteristic of this

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particular scenario is the (eventual) emergent hot continuum source and emission

lines, bi-polar morphology, and the fact that the bolometric luminosity should not

decrease to pre-outburst levels in the next decades (e.g., Kwok 1993).

Electon-Capture Supernova: M ≈ 9 M⊙

The timescale estimated in equation (4.2) is of the right order of magnitude to

be associated with the onset of carbon burning in relatively low-mass massive stars.

This is traditionally a very difficult phase to model (see the summary in Woosley et

al. 2002; Siess 2006; Poelarends et al. 2008).

One of the most intriguing explanations for the physics of SN 2008S-like

transients is that they result from electron-capture SNe (ecSNe) of O-Ne-Mg cores

of relatively low-mass massive stars (Miyaji et al. 1980). While speculative, this

explanation accounts for many of the observed characteristics of both the transients

and their progenitors. In particular, it accounts for the fact that the progenitors of

NGC 300 and SN 2008S were relatively low luminosity and deeply embedded. Here

we follow the scenario detailed by Poelarends et al. (2008) (see also Nomoto 1984a,

1987; Ritossa et al. 1996, Seiss 2006, 2007; as well as Chugai 1997b; Wheeler et

al. 1998; Woosley et al. 2002; Eldridge et al. 2007; Wanajo et al. 2009).

We know from the properties of the observed progenitors and our analysis of the

luminous stars in M33 that the progenitors are extreme AGB stars. In these systems,

the combination of thermal pulses due to He shell burning and dredge-up produces a

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massive, dusty wind (for lower luminosity and less enshrouded analogs, see the work

on carbon stars in the Magellanic clouds by Groenewegen et al. 2007, as well as van

Loon et al. 2005, 2006). In the Poelarends et al. (2008) models, the mass loss peaks

at nearly 10−4 M⊙ yr−1 for stars of mass M ≃ 9M⊙ and luminosity L ≃ 105L⊙, and

then drops precipitously for slightly more massive stars which can support core Neon

burning, and which eventually become normal iron-core core-collapse supernovae

(ccSNe) (see Fig. 13 from Poelarends et al. 2008). The thermal pulses driving the

mass loss occur at very high rates for the EAGB stars (on timescales of years),

suggesting that mass loss may appear as a steady wind, as seems to be required for

the small variability in the lightcurves of the SN 2008S and NGC 300 progenitors

(see Fig. 4.7; discussion in Prieto et al. 2008c), rather than as impulsive ejections of

optically-thick shells expected for normal carbon stars. The low degree of variability

seen in SN 2008S (particularly when contrasted with the EAGB stars in Fig. 4.7)

might also be explained by the onset of core carbon burning as the final phases of

stellar evolution commence. Note that the work of Nomoto (1984a, 1987) implies

that the EAGB envelope and mass-loss would be carbon-enhanced.

For a narrow mass range near 9 M⊙, the balance between mass loss and the

growth of the core allows the core to become unstable to collapse without igniting core

Neon burning, leading to an ecSNe. This occurs only for a small model-dependent

mass range near 9 M⊙ (see also, e.g., Nomoto 1984a; Seiss 2006, 2007). Poelarends

et al. (2008) estimated a mass range of Mmin ≃ 9 M⊙ to Mmax ≃ 9.25 M⊙. For

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a standard Salpeter IMF, and assuming that all stars from Mmax to 20 M⊙ form

ccSNe, then the ecSN fraction is just ≈ 6% (9.0 ≤ M ≤ 9.25 M⊙). This is below

our (albeit, uncertain) fiducial estimate for the rate of SN 2008S-like transients

relative to the normal ccSN rate, fSN ∼ 0.2. However, other studies have found

somewhat broader mass ranges for ecSNe. For example, Seiss (2007) find that

Mmax −Mmin ≈ 1− 1.5 M⊙, which implies a fractional rate for ecSNe more in accord

with our nominal estimate for fSN.

While the observed rate of SN 2008S-like transients is uncertain, we have argued

that they represent a modest fraction of the normal ccSN rate, consistent with a

limited progenitor mass range. The relatively low luminosity of their progenitors

implies that they are low-mass massive stars, potentially near the boundary between

electron-capture and normal core-collapse supernovae. In addition, Kitaura et

al. (2006) argue that ecSNe should be sub-luminous compared to normal ccSNe,

because of their low Ni yields, potentially explaining the low luminosity of SN

2008S-like transients. Finally, for the fiducial ecSN model of Poelarends et al. (2008)

(see also Nomoto 1984a), the AGB phase lasts for ≈ 4 × 104 yr, which although

short with respect to the lifetime of the star itself, is of order tEAGB in equation (4.2)

based on the number of analogs to the SN 2008S and NGC 300 progenitors in M33.

Of course, these timescales need not be identical, since the fiducial ecSN progenitors

of Poelarends et al. may evolve significantly in color as a result of mass-loss during

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their super-(“extreme”) AGB phase, becoming increasingly like the SN 2008S and

NGC 300 progenitors as they approach the end of their lives.

Fortunately, this speculative explanation has at least one simple and testable

prediction: there should be no surviving progenitor once the transient fades. This

may be testable in the optical, since most of the dust enshrouding the progenitor

was likely destroyed by the explosion, but observations in the MIR will be required

to be certain that a new shroud has not formed. A second test is to find strong

evidence in the late-time lightcurve for synthesized 56Ni. This may be difficult

both because ecSN may produce little 56Ni (Kitaura et al. 2006), and because dust,

whether that remaining from the EAGB phase or dust formed in the ejecta (e.g.,

Prieto et al. 2008e for M85), may make it difficult to correctly measure the late-time

decay rate. Spitzer may again be key in constraining the nature of these events

because of this obscuration. A third test is to ensure that Spitzer has surveyed

all the nearby galaxies so that future examples of these transients can be causally

connected to their deeply obscured progenitors. Finally, as with ccSNe, very nearby

ecSNe (D ∼< 300 kpc) should produce neutrino signatures characteristic of neutron

star formation, which detectors such as SuperKamiokande and its successors would

observe (see, e.g., Thompson et al. 2003; Kistler et al. 2008).

The recent paper by Botticella et al. (2009) in part corroborates the

interpretation discussed in Prieto et al. (2008a) and proposed here that SN 2008S

may be an electron-caputre supernova. They present the late-time quasi-bolometric

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lightcurve of SN 2008S, which shows evidence for a power-law time dependence with

a slope indicative of being powered by the radioactive decay 56Co. Although this

need not uniquely signal ecSN as the physical mechanism (see §4.5.2), it provides

some evidence for core-collapse, and something perhaps akin to normal neutron star

formation (Kitaura et al. 2006). A lightcurve with similar cadence and photometric

coverage has recently been published in Bond et al. (2009) for NGC 300.

Intrinsically Low-Luminosity Iron Core-Collapse Supernova:

M ∼ 10 − 12 M⊙

Heger et al. (1997) discuss a mechanism for generating a potentially obscuring

“superwind” via pulsational mass-loss in red supergiants between 10 and 20 M⊙

during the last 104 yr before explosion. The prediction of enhanced AGB-like

obscuration (they compare directly with luminous OH/IR stars), the evolutionary

timescale (compare their 104 yr with our eq. 4.2), and the secular increase in the

fundamental mode pulsation period as the star approaches death (see their Fig. 7)

are all in good agreement with the requirements on the 2008S and NGC 300

progenitors we discuss in this paper.

The physical mechanism of iron core-collapse supernovae is unknown (Rampp

& Janka 2000; Liebendorfer et al. 2001; Thompson et al. 2003; Buras et al. 2003;

Burrows et al. 2006). Recent observations hint that low-luminosity Type IIP

supernovae may be more common than previously thought (e.g., Chugai &

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Utrobin 2000; Pastorello et al. 2004, 2007), particularly when one accounts for the

incompleteness corrections discussed in §4.3. Because the mechanism of supernovae

has yet to be conclusively identified, it is difficult to interpret the diversity in

inferred 56Ni yield physically. In fact, that diversity may be larger than previously

thought, and we are only now appreciating the existence of a very low-luminosity

tail to the Type IIP luminosity function. If so, it is natural to imagine that these

low-luminosity core-collapse events might have analogs that occur in the very dusty

circumstellar medium of their massive stellar progenitors, as in Heger et al. (1997),

and thus may give rise to events like SN 2008S and NGC 300.

This scenario yields many of the predictions of the ecSN scenario discussed in

§4.5.2. Indeed, even with a complete sampling of “ec-” and “cc-” supernovae, it

may be difficult to disentangle the two populations since many of the predictions

— radioactive decay powered lightcurves, potentially embedded progenitor, no

“postgenitor” — are the same in both.

Massive Star Outburst: M ≈ 10 − 15 M⊙

On the basis of the relatively low luminosity of their progenitors, we view the

ecSN and massive white dwarf birth scenarios discussed above as the most probable

explanation SN 2008S and NGC 300. Nevertheless, there is of course the possibility

that they are instead a new class of outbursts from relatively low-mass massive

stars, potentially analogous to the pulsational instabilities discussed in Poelarends

123

et al. (2008) or Heger et al. (1997). The majority of the true “LBV” eruptions

with documented progenitors (e.g., 1997bs, 2002kg) came from optically bright

massive stars significantly more bolometrically luminous than the progenitors of

SN 2008S and NGC 300. As we have shown in Figures 4.2 and 4.5, the EAGB

population is separate from the sources traditionally classified as LBVs: they are

less bolometrically luminous and much more dust-obscured. These facts suggest

that if these transients were the outbursts of massive stars then they are distinct

from from the classical supernova impostors. If these events are not supernovae,

but merely outbursts, then their existence is likely connected to the physics of the

transition between stars that become ecSNe and/or ccSNe, and those that do not.

The degree of dust-obscuration at outburst is a crucial clue to their evolution. A

simple prediction of this possibility is that the progenitors of SN 2008S and NGC 300

should eventually be re-discovered in the optical and/or infrared after the outburst

emission has faded. For further arguments on the nature of SN 2008S and NGC 300

related to this discussion, see Bond et al. (2008), Smith et al. (2009b), and Berger et

al. (2008).

4.5.3. A More Complete Census

Equation (4.1) implies that a fraction ∼ 1× 10−4 of the massive star population

in any given galaxy appears to be in the evolutionary state that led to the explosions

observed as SN 2008S and NGC 300. The simplest explanation, adopted throughout

124

this work, is that the deeply dust-enshrouded phase marks the last tEAGB ∼< 104 yr

(eq. 4.2) in the life of a fraction fSN ≈ 0.2 of the massive star population.

Compilations of star formation and supernova rates in the local universe (e.g.,

Ando et al. 2005) suggest that the latter is ≈ 2 yr−1 within 10 Mpc (see §4.1, point

2), implying that there are ∼ 5 × 106 massive stars and ∼< 103 EAGB stars within

this volume. If the lifetime in the pre-explosion, highly dust-obscured phase is tEAGB

(eq. 4.2), we would expect to see one SN 2008S-like transient every few years, in

accord with our estimate for fSN.

A multi-epoch survey of all the local star-forming galaxies within 10 Mpc

with Spitzer would allow for a comprehensive census of EAGB stars. It would

significantly increase our knowledge of the variability properties and SED evolution

of these objects, and it might allow us to define more strict criteria for inclusion in

the class of SN 2008S/NGC 300-like progenitors. It would therefore decrease the

considerable uncertainty in NEAGB in equations (4.1) and (4.2). Coupled with the

supernova surveys in the local volume, such a study would improve our knowledge of

the fraction fSN of stars that eventually go through the deeply embedded phase just

before explosion. Of course, the most intriguing possibility is that the number of

true analogs to the progenitors of SN 2008S and NGC 300 is in fact NEAGB ∼ 0 − 1

in M33 and that tEAGB ∼< few − 103 yr.13 If so, the final catalog of EAGB stars that

13The lower limit here comes from the fact that pre-explosion imaging of the SN 2008S and NGC

300 progenitors establishes a few year baseline.

125

would be produced by a Spitzer survey would have just ∼ 50 − 100 members. These

could be followed up repeatedly, since, given these numbers one would expect to wait

just ∼ 10 years before one of these individual sources exploded. This would give a

direct observational link in the causal mapping between a sub-population of massive

stellar progenitors and their explosions, connecting them with a short timescale.

Indeed, the ability to identify an individual star as marked for imminent death (or

eruption) would be an astonishing consequence of this work.

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Fig. 4.1.— Mid-infrared color-magnitude diagram for M33. Absolute magnitudeat 4.5µm is plotted versus [3.6] − [4.5] color for all detected sources (3σ limitsdenoted by dashed lines; see §4.4.1; Table 4.1). For comparison, the positions of theprogenitors of NGC 300 (filled square) and SN 2008S (filled triangle; lower limit) arealso plotted (Prieto et al. 2008c; Prieto 2008d). Stars analogous to these progenitorsare exceedingly rare. Compare with Figure 4.2. The main sequence, AGB, and EAGBstars are clearly visible. Note the “spur” in the data extending to fainter M4.5 andredder color at [3.6] − [4.5] ≈ 1 and M4.5 ≈ −10.5, originating at the red extremumof the AGB population. To our knowledge, this is the first time that such a featurehas been identified in a MIR CMD.

127

Fig. 4.2.— Same as Figure 4.1, but focused on the red and bright region of interest.Here, the [3.6] − [4.5] > 1.5 and M4.5 < −10 selection for extremely red and brightobjects is shown explicitly by the dotted lines, as are the 18 EAGB sources in M33that meet these criteria (large open circles; see §4.4.2). We denote them “S1-S18,”ordered by color. The reddest source, S1, is shown in Figure 4.3. Most of these sourcesare highly variable (see Figs. 4.7, A.1, & A.2). SEDs are shown in Figure 4.5 (seeTable 4.2). The EAGB population is clearly distinct from the optically-luminous LBVcandidates from the catalog of Massey et al. 2007, shown by the open triangles (see§4.4; Table 4.3). Most of the LBV candidates are considerably more bolometricallyluminous and much less variable at 4.5µm than the EAGB sample (see Figs. 4.5, 4.7,B.1, & B.2). A number of cool hyper-giants such as VY CMa, NML Cyg, and M33Var A, as well as η-Carina, are shown for comparison. The brightest LBV candidateis M33 Var C. The LBV candidates appear to be bimodal in MIR color (right panelof Fig. 4.5).

128

3.6 4.5

V I

Fig. 4.3.— Image showing the point source S1 at 3.6µm, 4.5µm, as well as V and Iband (Massey et al. 2006). This is the reddest of the 18 sources in [3.6] − [4.5] colorselected in Figure 4.1. It is optically obscured (see Fig. 4.5). S1 is also highly variablein color and 4.5µm magnitude (see Figs. 4.6 & 4.7).

129

Fig. 4.4.— Color-color diagram showing [5.8] − [8.0] versus [3.6] − [4.5] colors for all(≈ 1800) the sources detected in all four IRAC bands. Symbols are the same as inFigure 4.1. The solid line and open squares show the expectation for a blackbodyof temperature TBB = 5000, 1000, 400, and 280 K. The small filled points with[5.8] − [8.0] ≈ 1.6 − 1.8 and [3.6] − [4.5] ∼> 2 are not sufficiently bright at 4.5µm(M4.5 < −10) to be included in the sample defined by the dotted lines in Figure 4.2.

130

Fig. 4.5.— Left Panel: Spectral energy distributions of the 18 EAGB stars in M33with [3.6] − [4.5] > 1.5 and M4.5 < −10 (see Fig. 4.2; Table 4.2). The lower dottedline shows the best lower limits obtained in the optical, whereas the upper dotted lineshows the worst lower limits at U and I and the two BV R detections (see discussion§4.4.3; Table 4.2). The SEDs of the SN 2008S (filled triangles) and NGC 300 (filledsquares) progenitors are also shown. For all the sources in M33 we assume a totalextinction of E(B − V ) = 0.1 mag. Right Panel: Spectral energy distributions of the16 LBVs detected in MIR from the Hα selected catalog of Massey et al. 2007 (seeFig. 4.2; Table 4.3). The relative increase in λLλ in the R-band (0.6 µm) is due tothe presence of strong Hα emission.

131

Fig. 4.6.— Left panel: Lightcurve for S1 (see Fig. 4.3), the reddest of the 18 EAGBstars selected in Figure 4.2. Note the high degree of variability, which for this source, isinconsistent with the progenitor of SN 2008S and (potentially) NGC 300 (see Fig. 4.7).Lightcurves for all of the EAGB stars are shown in Figures A.1 and A.2. Althoughnot strong in this example, most of the sources exhibit correlated color-magnitudevariations. Right panel: 4.5µm lightcurve for one of the LBV candidates from Masseyet al. 2007 (see Figs. B.1 & B.2 for the complete set). The large majority of the 16LBV candidates are not highly variable at 4.5µm, although there are two exceptions(Fig. 4.7).

132

Fig. 4.7.— RMS variation in 4.5µm magnitude as a function of [3.6] − [4.5] color forall of the sources detected in 3.6µm and 4.5µm with M4.5 < −10 (points) and the 16LBV candidates (open triangles). As in Figure 4.2, the 18 EAGB stars are denotedwith larger open circles. As noted in Figures A.1 & A.2, as well as B.1 & B.2, theextreme-AGB stars are highly variable, whereas all but two of the LBV candidatesare not. Variability of the progenitors of SN 2008S and NGC 300 is also shown. Forthe former, the data are taken from Prieto et al. 2008c (their Fig. 2). For NGC 300only two epochs are available and hence the value of the RMS (≈ 0.05) is a lowerlimit on the variability of the progenitor.

133

RA Dec [3.6] σ3.6 [4.5] σ4.5

(deg) (deg) (mag) (mag)

23.03815 30.82294 18.70 0.21 18.19 0.12

23.04489 30.82397 17.67 0.07 18.24 0.10

23.04532 30.81396 18.15 0.14 18.59 0.15

23.04906 30.81863 15.69 0.05 16.19 0.07

23.04955 30.82077 17.52 0.08 17.82 0.13

23.05103 30.82549 16.79 0.06 17.13 0.07

23.05131 30.81658 17.96 0.09 18.33 0.12

23.05277 30.79529 16.50 0.07 17.13 0.07

23.05287 30.80321 17.45 0.06 17.94 0.09

23.05303 30.82159 16.96 0.05 17.38 0.07

23.05314 30.78955 18.73 0.11 18.41 0.19

23.05333 30.81954 18.14 0.07 18.13 0.11

23.05470 30.79727 17.57 0.06 18.04 0.12

23.05473 30.81597 17.84 0.07 18.29 0.10

23.05481 30.80367 17.86 0.06 18.14 0.10

. . . . . . . . . . . . . . . . . .

Table 4.1. MIR Catalog for 53,194 Point Sources in M33

134

Name RA Dec Ua Ba V a Ra Ia [3.6] [4.5] [5.8] [8.0]

(deg) (deg) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag)

S1 23.45485 30.85704 23.47 23.76 23.37 22.90 22.47 16.35 14.21 12.55 11.37

S2 23.44397 30.79731 23.96 24.15 23.55 23.15 22.47 16.84 14.85 13.52 12.42

S3 23.56813 30.87755 23.91 23.97 23.48 23.10 22.50 16.56 14.73 13.42 12.55

S4 23.43452 30.57106 23.87 23.74 23.21 22.71 22.02 15.87 14.15 13.23 12.07

S5 23.40436 30.51738 22.99 23.10 22.97 22.71 22.10 15.08 13.38 12.12 11.32

S6 23.49194 30.82791 24.10 23.84 23.42 22.89 22.09 16.51 14.84 13.46 12.32

S7 23.55640 30.55211 24.52 24.50 23.86 23.31 22.75 14.29 12.63 11.18 10.17

S8 23.46408 30.64138 22.79b 23.11b 22.32b 21.75b 20.96b 16.17 14.52 14.14 12.87

S9 23.53817 30.73269 21.75b 21.93b 22.14b 22.10b 21.89b 15.67 14.08 12.69 11.51

S10 23.29877 30.59901 24.31 23.97 23.67 23.30 22.75 16.35 14.78 13.55 12.37

S11 23.29777 30.50744 23.82 23.64 23.53 23.05 22.83 15.30 13.72 12.47 11.26

S12 23.55236 30.90564 24.20 23.96 23.50 23.21 22.62 14.90 13.33 12.01 11.08

S13 23.37907 30.70096 24.05 24.14 23.60 23.03 22.41 16.38 14.82 13.44 12.36

S14 23.26238 30.34469 23.96 23.95 23.48 23.03 22.47 16.33 14.79 13.61 12.43

S15 23.39709 30.67737 24.13 24.02 23.29 22.80 22.23 15.25 13.71 12.22 10.97

S16 23.43722 30.64242 23.75 23.39 22.79 21.99 21.39 16.30 14.76 14.15 12.76

S17 23.47176 30.67430 23.81 23.67 23.03 22.51 21.60 15.27 13.74 13.15 11.56

S18 23.34248 30.64602 24.21 23.95 23.65 23.22 22.63 14.86 13.32 12.03 10.88

aExcept where otherwise noted, all UBV RI data in this table are upper limits.

bSource detections. Magnitudes from Massey et al. (2006).

Table 4.2. Photometry for the 18 EAGBs in M33

135

Identifier U B V R I [3.6] [4.5] [5.8] [8.0]

(from Massey et al. 2007) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag) (mag)

J013248.26+303950.4 16.13 17.32 17.25 17.01 16.89 15.70 15.36 15.02 14.57

J013324.62+302328.4 18.66 19.49 19.58 19.27 19.53 14.26 13.55 13.04 12.27

J013333.22+303343.4 18.30 19.30 19.40 18.44 18.89 13.09 12.50 12.20 11.65

J013335.14+303600.4 15.54 16.53 16.43 16.30 16.14 14.45 14.35 14.48 14.15

J013341.28+302237.2 15.18 16.24 16.29 16.28 16.32 15.08 14.96 15.07 14.45

J013350.12+304126.6 15.76 16.85 16.82 16.43 16.30 12.22 11.65 11.06 10.64

J013406.63+304147.8 15.12 16.26 16.08 15.86 15.76 14.41 13.95 13.83 · · · a

J013410.93+303437.6 15.13 16.13 16.03 15.87 15.70 14.84 14.65 14.69 13.97

J013416.10+303344.9 16.32 17.17 17.12 16.96 16.85 15.69 15.48 15.01 · · · a

J013422.91+304411.0 16.36 17.28 17.21 17.14 17.07 16.24 16.35 · · · a · · · a

J013424.78+303306.6 16.21 16.97 16.84 16.72 16.54 15.24 15.00 14.74 13.97

J013426.11+303424.7 18.38 19.23 18.97 18.59 18.27 14.57 13.79 13.09 12.24

J013429.64+303732.1 16.24 17.12 17.11 17.05 16.93 16.55 16.64 · · · a · · · a

J013442.14+303216.0 18.35 18.20 17.34 16.89 16.44 13.56 13.02 12.39 11.62

J013459.47+303701.9 17.35 18.59 18.37 17.88 17.69 14.42 13.73 13.16 12.31

J013500.30+304150.9 18.37 19.22 19.30 18.60 19.11 13.93 13.24 12.54 11.93

aNon-detection.

Table 4.3. Photometry for the 16 LBV Candidates in M33 from Massey et al. (2007)

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Chapter 5

A Spitzer/IRS Spectrum of the 2008 Luminous

Transient in NGC 300: Connection to

Proto-Planetary Nebula

5.1. Introduction

An intriguing luminous optical transient was discovered in the nearby galaxy

NGC 300 (hereafter NGC 300-OT) by the amateur astronomer B. Monard on May

14.1, 2008 (Monard 2008). The transient was faint compared to normal core-collapse

supernovae, with an absolute magnitude at discovery of MV ≃ −13 (Bond et

al. 2008). The optical spectrum obtained by Bond et al. (2008) close to discovery was

dominated by relatively narrow Hydrogen Balmer and Ca II lines (infrared triplet

and forbidden doublet) in emission, as well as strong Ca II H&K in absorption.

Shortly after discovery, Berger & Soderberg (2008) reported strong upper limits

on the optical luminosity of the progenitor star obtained from deep archival HST

data, which led them to suggest that the progenitor was a low-mass main sequence

star and the transient was a stellar merger, similar to the red Galactic nova V838

Monocerotis (e.g., Bond et al. 2003).

137

However, Prieto (2008d) reported the discovery of a luminous mid-infrared

(mid-IR) progenitor to the transient in archival Spitzer images. The progenitor was

a luminous dust-enshrouded star, whose spectral energy distribution was consistent

with a black-body of R ≃ 300 AU radiating at T ≃ 300 K, with Lbol ≃ 6 × 104 L⊙.

This discovery showed that NGC 300-OT was connected to an energetic explosion

in a relatively low mass ∼ 10 M⊙ star. The relatively low luminosity of the

transient compared to normal core-collapse supernovae, spectral properties, and

dust-enshrouded nature of the progenitor star, made NGC 300-OT a “twin” of

SN 2008S (Prieto et al. 2008c; Prieto 2008d), which was discovered earlier in 2008

in the galaxy NGC 6946 (Arbour & Boles 2008; Stanishev et al. 2008; Chandra &

Soderberg 2008; Steele et al. 2008; Yee et al. 2008; Wesson et al. 2008).

There have been a number of studies of NGC 300-OT and SN 2008S, and

the nature of these transients is still under debate (Prieto et al. 2008c; Thompson

et al. 2008; Smith et al. 2009b; Bond et al. 2009; Berger et al. 2009; Boticella et

al. 2009; Wesson et al. 2009). Thompson et al. (2008) present and discuss various

possible physical mechanisms that can explain these transients and the likely

range of main-sequence masses of the dusty progenitor stars that are consistent

with the observations of NGC 300-OT and SN 2008S: (1) massive white-dwarf

birth (MZAMS ≈ 6 − 8 M⊙); (2) electron-capture supernova (MZAMS ≈ 9 M⊙); (3)

intrinsically low-luminosity iron core-collapse supernova (MZAMS ≈ 10 − 12 M⊙);

and (4) massive star outburst (MZAMS ≈ 10 − 15 M⊙). Any of these potential

138

scenarios suggests that these transients are very important for our understanding

of the evolution of stars at the dividing line between “high” and “low” mass (i.e.,

8 − 10 M⊙).

Here we report on a low-resolution mid-IR spectrum of NGC 300-OT obtained

with Spitzer on August 14, 2008, 93 days after the discovery and 113 days after the

first detection (Monard 2008). The transient is luminous in the mid-IR spectral

range and shows broad emission features that we interpret as signs of carbon-rich

dust, similar to the spectra of carbon-rich proto-planetary nebulae in the Galaxy.

The paper is organized as follows. In §5.2, we discuss the observations and data

reduction. In §5.3, we present the analysis of the spectrum and spectral energy

distribution of the transient. In §5.4, we discuss the implications of our findings.

Hereafter we adopt a distance of 1.88 Mpc to NGC 300 (Gieren et al. 2005).

5.2. Spitzer Observations

We observed NGC 300-OT with the Short-Low (SL; 5.2 − 14 µm,

R = ∆λ/λ = 60 − 120) module of the Infrared Spectrograph (IRS; Houck et

al. 2004) on August 14.4, 2008 (UT). The observations were obtained in staring

mode as part of a Spitzer Director’s Discretionary Time (DDT) proposal request

(PID 487; AOR key 28139008). The ramp time was set to 60 sec in both SL orders

(SL1: 7.4 − 14 µm, SL2: 5.2 − 7.5 µm) and 10 cycles were obtained, for a total

139

exposure time of 1200 sec on source, including the two nod positions (total of 20

images).

To reduce the data, we started from the basic calibrated data (BCD) from the

Spitzer Science Center pipeline (S18.1.0). We constructed a high S/N background

image for each order and nod by median combining all the other images. We

subtracted the background from the individual 2D images. Rogue pixels in the

background-subtracted images were cleaned with IRSCLEAN (v1.9). We used the

routines profile, ridge, extract and tune in the Spitzer IRS Custom Extractor

(SPICE) software package in order to extract flux-calibrated 1D spectra. The

spectra of each nod were median combined and the two orders were merged together

after applying a small multiplicative correction factor of 3.5% to the SL1 spectra.

Figure 5.1 shows the final combined spectrum with ±1σ error bars on the fluxes

estimated from the RMS in each pixel. The mean signal-to-noise ratio of the final

spectrum is ≃ 35 per pixel.

5.3. Analysis

5.3.1. Spectral Features

The Spitzer spectrum of NGC 300-OT presented in Figure 5.1 shows two

prominent broad emission features at ≈ 8.3 µm and ≈ 12.2 µm. There is also

140

a relatively narrow, but resolved, fainter feature at ≈ 9.7 µm. The significance

of this faint feature is quite uncertain and depends sensitively on the spectra

used to obtain the final combined spectrum. The main properties of the emission

features (central wavelength, FWHM, and integrated fluxes) present in the mid-IR

spectrum are shown in Table 5.1. They were obtained after fitting Gaussians to the

continuum-subtracted spectrum. The continuum was modeled using a high-order

(6th) polynomial fit over the wavelength regions: λ ≤ 7.5 µm, 9.2 − 9.3 µm,

10.25 − 10.6 µm, and λ ≥ 13.4 µm. We obtain consistent results if we use a spline

function to model the continuum.

In Figure 5.2 we compare the Spitzer spectrum of NGC 300-OT with mid-IR

spectra of type IIP supernovae. The spectra of SN 2004et (Kotak et al. 2009) and

SN 2005af (Kotak et al. 2008) were obtained from the Spitzer archive (PID 237,

20256). Unlike the late-time mid-IR spectra of normal type IIP supernovae (e.g.,

SN 2004dj, Kotak et al. 2007; SN 2005af, Kotak et al. 2008; SN 2004et, Kotak

et al. 2009) and SN 1987A (e.g., Roche et al. 1993; Wooden et al. 1993) that

are dominated by narrow fine-structure lines of stable Ni, Ar, Ne, Co, and some

molecular SiO in emission at ∼ 8 − 9 µm, the mid-IR spectrum of NGC 300-OT

presents broad features that are most likely dominated by emission from dust grains

in the circumstellar environment. The non-detection of fine-structure lines of Fe-peak

elements in emission suggests that the main source of dust heating at this epoch is

141

not the decay of radioactive 56Ni. This is consistent with the low 56Ni production

estimated by Botticella et al. (2009) in the case of a supernova explosion.

In Figure 5.3 we compare the Spitzer spectrum of NGC 300-OT with mid-IR

spectra of three evolved massive stars that have circumstellar dust. The spectrum

of the yellow-hypergiant IRC+10420 (e.g., Humphreys et al. 2002) was obtained

from the ISO catalog of SWS spectra (Sloan et al. 2003). The spectra of the

yellow-hypergiant M33 Var A (Humphreys et al. 2006) and the LMC B[e] supergiant

R66 (Kastner et al. 2006) are from the Spitzer archive (PID 5, 3426). The spectra of

IRC+10420, M33 Var A and R66 are dominated by the amorphous silicate emission

feature at 9.7 µm, characteristic of oxygen-rich dust. The spectrum of R66 also

contains emission features of polycyclic aromatic hydrocarbons at 6.2, 7.7, 8.6 and

11.3 µm (PAHs), indicating the presence of carbon-rich dust as well. It is clear from

Figure 5.3 that the mid-IR spectrum of NGC 300-OT does not resemble the spectra

of these evolved massive stars with circumstellar dust, even though the optical

spectrum of the transient is strikingly similar to IRC+10420 (Bond et al. 2009; see

Smith et al. 2009b for the case of SN 2008S).

In Figure 5.4 we compare the spectrum of NGC 300-OT with mid-IR spectra of

Galactic proto-planetary nebulae (pPNe). Two of the pPNe have carbon-rich dust

(IRAS 20000+3239 and IRAS 13416-6243) and the others have oxygen-rich dust

(IRAS 15452-5459 and IRAS 17150-3224). The spectra are all from the ISO/SWS

catalog (Sloan et al. 2003). The spectra of oxygen-rich pPNe contain SiO absorption

142

at 7.9 µm and a strong silicate absorption feature at 9.7 µm, which do not appear

to be present in the spectrum of NGC 300-OT. Also the central wavelengths and

FWHM of the two “bumps” at ∼ 8.5 µm and ∼> 12 µm are inconsistent with the

features in the spectrum of NGC 300-OT.

The spectrum of NGC 300-OT is most similar to the spectra of carbon-rich

pPNe in Figure 5.4. They contain broad emission features at ∼ 8 µm and

∼ 12 µm, which have been associated with C-C and C-H bending and stretching

modes identified as the carriers of PAHs (e.g., Duley & Williams 1981; Peeters et

al. 2002). Note, however, that the spectrum of NGC 300-OT does not contain

the 6.2 µm PAH feature that is clearly present in the two carbon-rich pPNe. We

can put a 3σ upper limit on the integrated flux of the 6.2 µm PAH feature of

I6.2 < 2.1 (FWHM/0.2 µm) × 10−14 erg cm−2 s−1. This gives a 3σ limit on the

flux ratio of I6.2/I8.3 < 0.09 for an assumed FWHM = 0.2 µm. The spectrum of

IRAS 20000+3239 also has the 6.9 µm PAH feature, which is not present in the

spectrum of NGC 300-OT. We can put a 3σ upper limit on the integrated flux of the

6.9 µm PAH feature of I6.9 < 1.9 (FWHM/0.2 µm)× 10−14 erg cm−2 s−1, which gives

a limit on the flux ratio I6.9/I8.3 < 0.08 for an assumed FWHM = 0.2 µm.

143

5.3.2. Spectral Energy Distribution

We can construct the full spectral energy distribution (SED) of NGC 300-OT

at the epoch of the Spitzer spectrum using the optical and near-IR photometry

presented in Bond et al. (2009). Figure 5.5 shows the optical to mid-IR SED of the

transient 93 days after the discovery date. We have corrected all the fluxes for a

total extinction along the line-of-sight of E(B − V ) = 0.25 mag, which is the mean

of the extinction values reported in Bond et al. (2009). We assume RV = 3.1 and use

the Schlegel et al. (1998) reddening law, with Aλ ∝ λ−1.6 in the near-to-mid infrared

range. The filled circles are the optical (BV RI) and near-IR (JHK) fluxes from

Bond et al. (2009). The thick line is the Spitzer mid-IR spectrum. For comparison,

we show the SED of the luminous dust-enshrouded progenitor of NGC 300-OT (filled

squares) obtained from pre-explosion Spitzer IRAC (3.6 − 8 µm) and MIPS (24 µm)

photometry (see Table 5.2; these data were used in Thompson et al. 2008 and Bond

et al. 2009). At the epoch of the Spitzer observation the transient is ∼ 20 times

more luminous than the progenitor at 8 µm. The results of black-body fits to the

transient and progenitor SEDs are presented in Table 5.3.

The evolution of the light curves of NGC 300-OT in different filters presented

by Bond et al. (2009) shows that the transient becomes redder in time, with the color

evolving from V −K ≃ 3.1 mag at discovery to V − K ≃ 5.2 mag at the time of the

Spitzer spectrum. This fast evolution in the V −K color, while the B − V color only

144

changes from ≃ 0.8 mag to ≃ 1.1 mag in the same time period, suggests the presence

of warm circumstellar dust formed in the explosion or heated pre-existing dust.

Botticella et al. (2009) analyzed the SED of SN 2008S and showed that

the evolution in optical+near-IR fluxes could be explained with a single “hot”

black-body until ∼ 120 days after explosion, but they needed a second “warm”

black-body component at later times. They concluded that the near-IR flux excess

of SN 2008S at ∼> 120 days after explosion was possibly due to newly-formed dust in

the ejecta or shock-heated dust in the circumstellar environment.

As in the case of SN 2008S, we find that the BV RIJHK fluxes of NGC 300-OT

at 93 days after discovery (113 days after the first detection) can be well-fit by

the sum of two black-bodies (χ2ν = 1.8), a hot component with T1 = 3893 K,

R1 = 10.7 AU and a warm component with T2 = 1511 K, R2 = 67.3 AU (see the

dotted line in Figure 5.5). The total luminosity of these components is 2.1× 106 L⊙.

However, the mid-IR SED of NGC 300-OT traced by the Spitzer spectrum is a

factor of ∼ 2 − 9 brighter between 5 − 14 µm than the extrapolated sum of the two

black-bodies. We add a third black-body component with lower temperature, while

keeping the fit to the optical and near-IR fluxes fixed, to account for the mid-IR

excess. We find a good fit (χ2ν = 1.5) to the continuum of the Spitzer spectrum

(defined in §5.3.1) with T3 = 485 K and R3 = 515 AU. The luminosity of this

component is 6.0× 105 L⊙, which is 29% of the optical+near-IR luminosity and 22%

of the total integrated luminosity of the transient.

145

The two black-body components that can reproduce the near-IR and mid-IR

excesses in the SED of NGC 300-OT are likely due to emission from circumstellar

dust. The hotter component (T2 = 1511 K) can be reasonably explained with newly

formed dust in the ejecta, as proposed by Botticella et al. (2009) for SN 2008S.

The velocity inferred from the black-body radius R2 = 67 AU and the time after

the first detection of the transient1 is ∼ 1000 km s−1. This velocity is in the range

of velocities measured from emission lines in the optical spectra of NGC 300-OT of

∼ 100 − 1000 km s−1 (Bond et al. 2009; Berger et al. 2009).

Another way to explain this warm dust component would be emission from

pre-existing progenitor dust, although this seems less likely because of dust

destruction from the initial outburst light (e.g., Dwek 1983). Assuming a luminosity

at maximum light of Lmax ≈ 5 × 106 L⊙ (luminosity at discovery) and a dust

sublimation temperature of Tsub ≃ 1500 K, we obtain a radius of the dust-free cavity

of ∼ 80 − 250 AU depending on the assumed dust emisivity law, Qλ ∝ λ0 − λ−1.

This radius is a factor of ∼ 1.2 − 3.7 times larger than the black-body scale R2 of

the ∼ 1500 K temperature dust, which suggest the near-IR emitting dust may have

been formed in the ejecta.

1We assume this is a good estimate of the time after explosion. This seems like a reasonable

assumption given that the light curve of the transient is rising fast at that time (see Figure 1 in

Bond et al. 2009).

146

We can estimate the total dust mass needed to account for the luminosity of the

warm black-body component using equation (3) in Smith et al. (2005) and assuming

an average carbon grain density of ρdust ≃ 2.24 g cm−3,

Md ≈ 2.5 × 10−6(

Td

1000 K

)−6(

Ld

106L⊙

)

M⊙ (5.1)

where Td is the dust temperature in Kelvin, and Ld is the dust luminosity in L⊙.

This equation assumes that the dust grains have a constant radius a ∼< 0.2 µm,

and that the dust emissivity is approximated by the function presented in Gilman

(1974). Using T2 and L2 in Table 5.3, we obtain Md ≈ 2 × 10−7 M⊙ for the mass

of newly formed dust at 93 days after discovery. We can compare this dust mass

to that in SN 2008S. Under the same assumptions about dust properties considered

in equation (5.1), this dust mass is ∼ 3 times larger than was needed to explain

the near-IR excess in SN 2008S ∼ 120 days after explosion, but comparable to the

mass needed ∼ 230 days after explosion (using the data in Table 8 of Botticella et

al. 2009).

The mid-IR excess revealed by the Spitzer spectrum of NGC 300-OT cannot be

explained by newly formed dust. The constant expansion velocity needed to reach

a black-body radius of R3 = 515 AU at this epoch is v ∼ 8000 km s−1, far larger

than the velocities observed in the optical spectra of the transient. This component

must be emission from pre-existing dust from the progenitor. Dust grains that were

not destroyed by the initial outburst will absorb the outburst light, warm up, and

147

re-radiate at mid-IR wavelengths. Using equation (5.1) we estimate that a dust

mass of Md ≈ 10−4 M⊙ and dust optical depth of τV ≈ 0.4 is needed to account

for the mid-IR excess. For a gas-to-dust mass ratio of 200 that is consistent with

observations of evolved stars with carbon-rich dust in the LMC (e.g., Matsuura et

al. 2009), we find a gas mass of ≈ 0.02 M⊙. A similar mass of dust of ∼ 10−4 M⊙

is needed to explain the SED of the progenitor of NGC 300-OT, suggesting that

a substantial fraction of the dust in the progenitor wind survives the explosion.

Interestingly, a similar mass of dust is also needed to explain the mid-IR excess

observed in SN 2008S ∼ 17 days after explosion (Botticella et al. 2009).

We performed radiation transport calculations using DUSTY (Ivezic &

Elitzur 1997) to check if the SED of NGC 300-OT can also be explained by radiation

through a spherical shell of dust. We found a reasonable fit to the optical to mid-IR

SED using a central black-body with T= 4000 − 5000 K illuminating a dust shell

with density profile ρ(r) ∝ r−1, inner and outer radius of the shell Rin ≃ 100 AU and

Rout ≃ 10000 AU, temperature of the dust at the inner radius T = 1500 K, and dust

optical depth τV ≃ 1.2. We note that this solution is not unique since we also obtain

a reasonable fit to the SED with ρ(r) = constant, Rin ≃ 100 AU, and Rout/Rin ∼ 10,

with all the other parameters being equal. Wesson et al. (2009) analyzed the optical

to mid-IR SED of the progenitor of SN 2008S and the transient ∼ 17 days and

∼ 180 days after explosion using a radiative transfer code. They find that the SED

of SN 2008S can be explained with a central source illuminating a spherical dust

148

shell with density profile ρ ∝ r−2, inner radius Rin ≃ 1250 AU, and τV ≃ 0.8, where

∼> 98% of the pre-existing progenitor dust (Mdust ≃ 1.2 × 10−5 − 3.5 × 10−3 M⊙,

depending on Rout) survives the explosion. This is qualitatively similar to our results

for NGC 300-OT.

5.4. Discussion & Conclusions

We have presented a low-resolution mid-IR spectrum of NGC 300-OT obtained

with Spitzer on August 14.4, 2008, 93 days after the discovery of the transient. We

now present our discussion and interpretation of the results.

5.4.1. Mid-IR Spectrum and SED of NGC 300-OT

The mid-IR spectrum of NGC 300-OT shows broad emission features at 8.3 µm

and 12.2 µm that are similar to the broad features seen in the spectra of carbon-rich

pPNe in the Galaxy (e.g., Kwok 1993; Kwok et al. 2001), called “Class C” PAH

sources by Peeters et al. (2002). Joblin et al. (2008) derive profiles for the broad

8 µm and 12 µm features from the spectrum of the archetypal “Class C” PAH source,

the pPN IRAS 13416-6243 (see Figure 5.4). These broad features are attributed to

hydrocarbons with a predominantly aliphatic nature, which undergo photochemical

processing in proto-planetary nebula to transform into the more aromatic material

observed in carbon-rich planetary nebulae (e.g., Kwok et al. 2001; Sloan et al. 2007).

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Joblin et al. (2008) show that these broad features are also observed in young

planetary nebulae, and are distinct from the spectral features of neutral PAHs,

ionized PAHs, and very small grains.

It is interesting to note that the position of the center of the PAH complex

at 7 − 8 µm, observed in many astrophysical environments including the ISM

and evolved stars (e.g., Tielens 2008 and references therein), has been shown to

correlate with the effective temperature of the host star in Herbig A/Be stars,

planetary-nebulae, and pPNe (e.g., Sloan et al. 2007; Keller et al. 2008; Boersma

et al. 2008). The correlation goes in the sense that stars with lower effective

temperatures (i.e., weaker UV-optical radiation field) show the central peak of this

complex at longer wavelengths. The central wavelength at 8.3 µm detected in the

Spitzer spectrum of NGC 300-OT would imply an effective temperature of ∼ 4000 K

(see Fig. 8 in Keller et al. 2008), which is consistent with the temperature of the

hot black-body (T ≃ 3900 K) derived from the optical SED of the transient. This

provides indirect and independent support for our interpretation of the mid-IR

spectrum, and also evidence that UV processing has not yet converted predominantly

aliphatic hydrocarbons into PAHs in NGC 300-OT.

A noteworthy difference between the Spitzer spectrum of NGC 300-OT and the

spectra of carbon-rich pPNe is the non-detection of the 6.2 µm PAH emission feature

to fairly deep limits. In the carbon-rich pPN IRAS 13416-6243, for example, the

ratio of the integrated fluxes of the 6.2 µm and 8 µm features is I6.2/I8 ≃ 0.13, which

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is a factor of 1.3 higher than the 3σ limit for NGC 300-OT. The 6.2 µm emission

feature is thought to be produced by the C-C stretching mode in ionized PAHs. The

astronomical and laboratory spectra of PAHs and PAH-like molecules show such a

wide variety that the absence of the 6.2 µm feature may be explained by differences

in shape, ionization state, impurities, and size of the molecules (e.g., Pathak &

Rastogi 2008; Bauschlicher et al. 2009). On the other hand, the non-detection of the

6.9 µm PAH emission feature in the spectrum of NGC 300-OT seems to be consistent

with the integrated flux of this feature measured in some Galactic carbon-rich pPNe

like IRAS 20000+3239. This emission feature is thought to be associated with

aliphatic material (C-H bending mode) and is detected only in a fraction of pPNe

(e.g., Kwok et al. 1999).

The mid-IR excess traced by the Spitzer spectrum can be well explained by the

presence of warm circumstellar dust (T ∼ 500 K) with mass Md ∼ 10−4 M⊙. This

dust was most likely part of the dusty progenitor wind, pre-existing the luminous

explosion that produced the optical transient. The SED of NGC 300-OT at the

epoch of the Spitzer spectrum also shows a near-IR excess, which can be explained

with a small mass (∼ 2× 10−7 M⊙) of warm circumstellar dust (T ∼ 1500 K) formed

in the ejecta.

Alternatively, the SED of NGC 300-OT can be reasonably well explained

with a T = 4000 − 5000 K black-body illuminating a spherical shell of pre-existing

progenitor dust that extends from ∼ 100 AU to ∼ 10000 AU, where the inner radius

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of the dust shell marks the destruction of dust by the initial outburst light. The

presence of a substantial mass of pre-existing dust from the progenitor wind in the

overall SED of NGC 300-OT was also characteristic of SN 2008S (Wesson et al. 2009)

and indicates that most of the dust survived the explosion.

5.4.2. NGC 300-OT and SN 2008S: Connection to

Proto-Planetary Nebulae

The similarity of the mid-IR spectrum of NGC 300-OT with carbon-rich

pPNe is very striking and may shed new light on the nature of this transient,

SN 2008S, and the other optical transients (e.g., SN 1999bw, M85-OT) that show

similar characteristics and appear to be part of the same class (Prieto et al. 2008e;

Thompson et al. 2008). The optical spectra of NGC 300-OT and SN 2008S transients

were compared with the spectrum of the massive Galactic yellow-hypergiant star

IRC+10420 (Smith et al. 2009b; Bond et al. 2009; Berger et al. 2009), which shows

an F-type supergiant spectrum with Balmer lines in emission, as well as strong Ca II

triplet, [Ca II] doublet, and [Fe II] lines in emission. Given the similarity of the

mid-IR spectrum of NGC 300-OT with pPNe, we have searched in the literature

for their optical spectra. We found several examples of pPNe with optical spectra

that are remarkably similar to NGC 300-OT and SN 2008S. In the sample of echelle

long-slit spectra of evolved stars of Sanchez Contreras et al. (2008), the Galactic

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pPNe IRAS 17516-2525 (O-B spectral type), M1-92 (B2-F5), Hen 3-1475 (Be),

IRAS 22036+5306 (F4-7), and IRAS 08005-2356 (F4 Ie) show Balmer lines and also

strong Ca II triplet and [Ca II] doublet in emission. Other examples can be found

in the atlas of optical spectra of post-AGB stars presented in Suarez et al. (2006).

The presence of forbidden Ca II in emission in the spectra of NGC 300-OT and

SN 2008S, pPNe, and IRC+10420, which is rarely present in stellar spectra, means

that calcium is not depleted onto dust grains, most likely due to the destruction of

grains by relatively fast shocks (e.g., Hartigan et al. 1987).

Another similarity between NGC 300-OT and pPNe is revealed by the

kinematics and the detection of double-peaked Balmer and Ca II triplet lines in the

spectrum of NGC 300-OT. Bond et al. (2009) interpreted these double features as the

presence of a bipolar outflow with an expansion velocity of ≈ 75 km s−1, and possibly

faster components moving at ∼ 200 km s−1. Berger et al. (2009) discussed evidence

for even faster velocity components (including inflow) going up to ∼ 1000 km s−1.

Aspherical winds or outflow moving at velocities of a few × 100 km s−1 up to

∼ 1000 km s−1 are observed in pPNe (e.g., Balick & Frank 2002). Multiple studies

using high-resolution imaging of Galactic pPNe with HST have shown a variety of

complex morphologies, with bipolar, multipolar and point-symmetric structures (e.g.,

Sahai et al. 1999). In particular the pPN Hen 3-1475 (also classified as a young PN in

some studies), which is in the spectroscopic sample of Sanchez Contreras et al. (2008)

and whose spectrum shares many features with the spectra of NGC 300-OT, has

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a bipolar morphology and velocity components up to ∼ 1200 km s−1 (e.g., Riera

et al. 1995, 2003). Given the inferred luminosity of ∼ 104 L⊙, chemistry, and

kinematics, Riera et al. (1995, 2003) proposed that Hen 3-1475 was a relatively

high-mass star (∼ 3−5 M⊙) in the post-AGB phase of evolution. Another interesting

example is the Red Rectangle, an extensively studied intermediate-mass pPNe with

carbon-rich dust chemistry and a fast (∼ 560 km s−1) bipolar outflow traced by Hα

in emission (e.g., Witt et al. 2009).

Several studies have proposed that SN 2008S and NGC 300-OT were the result

of an energetic eruption in a dust-enshrouded 10 − 20 M⊙ star, where the star

survives the eruption. Smith et al. (2009b) discussed a super-Eddington wind as

the physical mechanism that produced SN 2008S, similar to the super-outbursts

of massive LBVs (e.g., van Marle et al. 2008). Berger et al. (2009) presented

possible observational evidence for this model from the complex kinematics that

they inferred from their high-resolution spectra of NGC 300-OT. Bond et al. (2009)

did not require that the progenitor of the transient was LBV-like, but rather an

OH/IR star (e.g., Wood et al. 1992) that was evolving to warmer temperatures (in a

blue-loop) at the time of the eruption. These studies relied heavily on comparing the

optical spectra of the transients with the spectrum of the massive yellow-hypergiant

IRC+10420 (e.g., Humphreys et al. 2002; Davies et al. 2007). However, as discussed

here, there are examples of pPNe in the Galaxy that share very similar optical

spectroscopic characteristics with NGC 300-OT and SN 2008S. In fact, the complex

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model of inflow-outflow put forth in Berger et al. (2009) to explain the spectra of

NGC 300-OT has been discussed in the context of fast winds of AGB and post-AGB

stars in binaries (Soker 2008).

Finally, in a mid-IR study of massive stars in the LMC, Bonanos et al. (2009)

argued that B[e] supergiants (e.g., R66 in Figure 5.3) may share a common

origin with NGC 300-OT and SN 2008S. Supergiant B[e] stars in the LMC are

very rare (only ∼ 10 discovered), have luminosities Lbol ∼> 104 L⊙, and dusty

circumstellar envelopes, properties that are broadly consistent with the properties

of the progenitors of NGC 300-OT and SN 2008S. However, the circumstellar dust

around B[e] supergiants in the LMC is significantly hotter (∼> 800 K) than the dust

around the progenitors, probably because they are oxygen-rich.

In summary, we have shown that NGC 300-OT and SN 2008S have several

properties (mid-IR spectrum, optical spectra, kinematics, and dusty circumstellar

medium) that are characteristic of pPNe in the Galaxy; they are not unique to

massive stars like IRC+10420.

5.4.3. The Progenitors of NGC 300-OT and SN 2008S:

Massive Carbon-rich AGB/post-AGB stars ?

The progenitors of NGC 300-OT and SN 2008S were luminous (∼ 4−6×104 L⊙)

dust-enshrouded stars with warm (T∼ 300− 450 K) circumstellar dust, found at the

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red extremum of the AGB sequence in a mid-IR color-magnitude diagram (Thompson

et al. 2008). They are part of the extreme-AGB (EAGB) sequence, which has been

identified as a continuation of the AGB to redder mid-IR colors in resolved stellar

population studies of nearby galaxies using Spitzer (e.g., LMC, Blum et al. 2006;

M33, Thompson et al. 2008). Their location in the mid-IR color-magnitude diagram

indicates extreme mass-loss and relatively cool circumstellar dust (e.g., Srinivasan

et al. 2009). Interestingly, Matsuura et al. (2009) find that most EAGB stars in

the LMC sample for which they have obtained mid-IR spectra have carbon-rich

dust, which is consistent with the evidence presented here for carbon-rich dust in

NGC 300-OT. Even though the number of carbon-rich AGBs in the LMC declines as

a function of luminosity with respect to oxygen-rich AGBs, interpreted as evidence

of Hot-Bottom-Burning which converts carbon into nitrogen and oxygen, there are

carbon stars with luminosities approaching those of the progenitors of NGC 300-OT

and SN 2008S (e.g., van Loon et al. 1997; Frost et al. 1998). One example is

IRAS 05278-6942, a carbon-rich AGB star in the LMC that has Lbol ∼ 4 × 104 L⊙

and M ∼ 3× 10−5 M⊙ yr−1 (Groenewegen et al. 2007). Indeed, Kastner et al. (2008)

in their study of the most luminous 8 µm sources in the LMC, point out that “more

high-Lbol carbon stars may lurk among the very red, unclassified objects” in their

sample.

The high-mass counterparts of AGB stars with MZAMS ≃ 8 − 10 M⊙, so-called

super-AGB stars, have been proposed as good candidates for the progenitors of

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NGC 300-OT and SN 2008S (see Thompson et al. 2008 and references therein;

Botticella et al. 2009). These stars end up with an O-Ne core and, depending on the

competing effects of core-growth after carbon ignition and strong mass-loss, they can

explode as electron-capture supernovae in a narrow and uncertain mass range around

∼ 9 M⊙ or end-up as O-Ne white dwarfs at lower masses (e.g., Nomoto 1984a;

Poelarends et al. 2008). The luminosities of SAGB stars in theoretical models

can reach ∼ 105 L⊙ at the end of their evolution (e.g., Siess 2007), comparable

to the luminosities of the progenitors of NGC 300-OT and SN 2008S. These

models also predict that the photospheric abundances of SAGB stars should be

oxygen-rich (C/O < 1) at the end of their evolution, through a combination of

Hot-Bottom-Burning and the occurrence of the third dredge-up. However, the

modeling of these processes in the AGB and SAGB evolution is very uncertain and

depends on several important factors like metallicity, the treatment of convection,

mass-loss, and the input opacities (e.g., Marigo 2008). In fact, there are theoretical

studies that have discussed the possibility of carbon-rich photospheres in massive

AGBs (e.g., Nomoto 1987; Marigo 2007).

An important difference between the progenitors of the luminous transients and

carbon-rich AGB and EAGB stars is that they did not show variability in the mid-IR

within 3 − 4 years of explosion (Prieto et al. 2008c; Thompson et al. 2008), whereas

most AGB and EAGB stars are highly variable (e.g., Gronewegen et al. 2007; Vijh

et al. 2009). Since variability in AGB stars is explained by pulsations that drive the

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mass-loss (thermal pulses), the lack of variability in the progenitors may indicate

that they were at the very tip of the AGB or SAGB phase before the explosion,

perhaps past a super-wind phase. If this is the case, the progenitor could be classified

as a pPN (i.e., it was in the post-AGB phase at the time of the explosion).

5.4.4. Progenitors and Transients: Concluding Remarks

The physical mechanism that produced the energetic explosions

(∼ 2 − 6 × 1047 erg in optical to near-IR light) of NGC 300-OT and SN 2008S

is still unknown. Although the observations presented here do not directly shed

light on the mechanism that produced the transients, we have shown that all the

observations of the transients and their progenitors presented thus far are consistent

with the explosion of a massive (MZAMS ∼ 6− 10 M⊙), carbon-rich AGB, super-AGB

or post-AGB star, either single or in a binary. An in-depth discussion of some of

the mechanisms that could explain the transients can be found in Thompson et

al. (2008). Here we briefly comment on the ones that involve a massive AGB or

post-AGB star: white dwarf formation and an electron-capture supernova.

In the case of an energetic eruption where the progenitor survives the explosion,

the transients could mark the birth of massive white dwarfs (Thompson et al. 2008).

Observations of mass-losing AGB stars show spherically symmetric envelopes, while

their descendants (proto-planetary and planetary nebulae) have highly asymmetric

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and complex morphologies and kinematics. This has been a long standing mystery

in stellar evolution for which several mechanisms have been proposed, with magnetic

fields, rotation and binaries suggested as primary suspects for breaking the symmetry

(e.g., Balick & Frank 2002; de Marco 2009; Sahai et al. 2009 and references therein).

In a recent study, Dennis et al. (2008) argue that pPNe outflows may be driven by

an explosive MHD launch mechanism similar to the ones discussed in the context

of supernovae and gamma-ray bursts. This model seems appealing when applied to

NGC 300-OT and SN 2008S – perhaps we are witnessing the launch of the jet in a

massive AGB which is shaping a pPN. In this scenario we expect that the pPN now

in formation will become a PN when the central white dwarf left behind ionizes the

surrounding gas. The timescale for this is very uncertain, but for a ∼ 8 M⊙ star it

can be of the order of ∼ 100 yr (e.g., Stanghellini & Renzini 2000). An interesting

prediction of an asymmetric outflow that can be tested with new observations is the

detection of strongly polaraized light from the transient.

An electron-capture supernova in a massive AGB star has been suggested as

a possible mechanism for NGC 300-OT and SN 2008S (e.g., Prieto et al. 2008c;

Thompson et al. 2008; Botticella et al. 2009). Two of the main predictions of this

scenario that can be tested with late-time observations are the disappearance of the

progenitor star years after the explosion and the detection of radioactive 56Ni decay

synthesized in the explosion. Botticella et al. (2009) presented detailed photometric

and spectroscopic observations of SN 2008S. Their main argument in favor of a

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supernova explosion as the origin of the transient was presented in the late-time

light curve. They found that the pseudo-bolometric light curve at t ∼> 140 days had

a decay slope consistent with radioactive decay of 56Co → 56Fe and inferred the

production of ∼ 10−3 M⊙ of 56Ni in the explosion. However, Smith et al. (2009b)

noted that the late time light curve of SN 2008S was slower than expected from 56Co

decay and argued against a supernova interpretation. While a late-time decay slope

slower than 56Co is a possibility in SN 2008S, it should be pointed out that slow

late-time light curve slopes (compared to 56Co decay) have also been observed in

some subluminous type IIP supernovae, including the very well-studied SN 2005cs

(Pastorello et al. 2009). Therefore, we think that a supernova explosion origin

cannot be excluded from this result alone. The late-time light curve of NGC 300-OT

should give very important clues about the possible supernova origin.

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Fig. 5.1.— Low-resolution Spitzer/IRS mid-IR spectrum of the NGC300-OT obtainedon 2008 August 14.4, 93 days after the transient was discovered. The vertical errorbars are the RMS of each pixel obtained after combining the 20 spectra.

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Fig. 5.2.— Comparison of the Spitzer/IRS spectrum of NGC300-OT (top) with mid-IR spectra of the type IIP supernovae SN 2004et and SN 2005af. In parenthesis arethe days with respect to the explosion date for each spectrum. We have subtracted alinear continuum fit in the wavelength region shown to each spectrum and scaled theflux arbitrarily.

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Fig. 5.3.— Comparison of the Spitzer/IRS spectrum of NGC300-OT (top) with mid-IR spectra of evolved massive stars that show circumstellar dust emission. Theseinclude the yellow-hypergiants IRC+10420 and M33 Var A, and the B[e] supergiantR66. We have subtracted a linear continuum fit in the wavelength region shown toeach spectrum and scaled the flux arbitrarily.

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Fig. 5.4.— Comparison of the Spitzer/IRS spectrum of NGC300-OT (top) withmid-IR spectra of Galactic proto-planetary nebulae. The pPNe have different dustchemistry: carbon-rich (IRAS 2000+3239 and IRAS 13416-6243) and oxygen-rich(IRAS 1545-549 and IRAS 17150-3224) dust. We have subtracted a linear continuumfit in the wavelength region shown to each spectrum and scaled the flux arbitrarily.

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Fig. 5.5.— SED of NGC 300-OT at 93 days after discovery. The filled circles are theoptical (BV RI) and near-IR (JHK) fluxes of NGC 300-OT from Bond et al. (2009).The Spitzer/IRS spectrum is shown as the thick black line. The dotted line is a fitto the optical+near-IR SED of the transient using the sum of two black-bodies withtemperatures T1 ≈ 3890 K and T2 ≈ 1500 K. The grey line is an extension of theprevious black-body fits to mid-IR by adding a cooler component with T3 ≈ 485 K.We show the SED of the progenitor of NGC 300-OT for comparison (filled squares).The dashed line is a single black-body fit with T ≈ 335 K to the progenitor SED.

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λc Intensity FWHM

(µm) (10−13 erg cm−2 s−1) (µm)

8.33 ± 0.01 2.27 ± 0.20 0.94 ± 0.03

9.71 ± 0.06 0.12 ± 0.11 0.49 ± 0.14

12.16 ± 0.07 0.39 ± 0.14 1.47 ± 0.16

Table 5.1. Features in the Spitzer/IRS Spectrum of NGC300-OT

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λ λFλ Telescope/Instrument

(10−14 erg cm−2 s−1)





24 µm 15.03 ± 1.37 Spitzer/MIPS

Table 5.2. Spectral Energy Distribution of the Progenitor of NGC 300-OT

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Parameter Value

NGC 300-OT BV RIJHK fluxes

T1 (K) 3893

R1 (AU) 10.7

L1 (L⊙) 1.1 × 106

T2 (K) 1511

R2 (AU) 67.3

L2 (L⊙) 9.8 × 105

NGC 300-OT adding fit to Spitzer data

T3 (K) 485

R3 (AU) 515

L3 (L⊙) 6.0 × 105

Progenitor SED

T (K) 335

R (AU) 332

L (L⊙) 5.7 × 104

Table 5.3. Black-body Fits to the Transient and Progenitor SEDs

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Chapter 6

A Study of the Type Ia/IIn Supernova 2005gj

from X-ray to the Infrared

6.1. Introduction

Thermonuclear supernova explosions (Type Ia supernovae, SN Ia hereafter) are

believed to be the detonation or deflagration of a white dwarf accreting matter from

a companion star (Arnett 1982). The mass of the white dwarf slowly increases until

it approaches the Chandrasekhar limit where the star becomes thermally unstable.

At this point fusion of Carbon and Oxygen begins near the center and quickly moves

through most of the star before degeneracy is lifted. The result is a spectacular and

powerful explosion that is visible across much of the Universe. Since SN Ia arise from

a narrow range of white dwarf masses, their peak luminosities are very consistent

and they make excellent distance indicators (e.g., Phillips 1993). SN Ia are powerful

probes of cosmology and have been instrumental in narrowing the uncertainty in

the Hubble parameter, discovery of the accelerating universe, and constraining dark

energy models (Hamuy et al. 1995, 1996b; Riess et al. 1998; Perlmutter et al. 1999;

Astier et al. 2006; Wood-Vasey et al. 2007).

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But the use of SN Ia as reliable distance indicators will always be questioned

until the progenitor and explosion physics are well-understood. What types of

binaries create SN Ia? How is matter transferred to the white dwarf without causing

thermonuclear runaways on the surface? Are there several types of progenitors?.

These big questions remain to be answered and detailed observations of hundreds of

events have yielded few clues.

In 2002, Hamuy et al. (2003) identified a new kind of supernova. The early

spectrum of SN 2002ic was a cross between a Type Ia event and a Type IIn (Deng

et al. 2004, have called this type of supernova a “IIa”), showing P-Cygni features

similar to SN Ia and resolved Balmer lines in emission. Type IIn supernovae are

core-collapse explosions going off in dense circumstellar environments (Schlegel 1990;

Chevalier & Fransson 1994). They are relatively common since the massive stars

that create core-collapse supernovae often have thick winds. If the interpretations of

the pre-explosion observations of SN 2005gl are correct, SN IIn could be associated

in some cases with luminous blue variables (Gal-Yam et al. 2007).

In the case of SN 2002ic, the presence of Balmer lines with profiles characteristic

of SN IIn, and the high luminosities and slow decline after maximum lead to the

conclusion that most of the energy came from the interaction of the ejecta with a

dense circumstellar medium (CSM). Other Type IIn events (SN 1997cy, Germany

et al. 2000; SN 1999E, Rigon et al. 2003) have been re-classified as SN 2002ic-like

that were caught late in their evolution (Hamuy et al. 2003; Wood-Vasey et al. 2004).

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SN 2002ic provided the first direct evidence that thermonuclear explosions can

also occur in a dense medium, but in this case the circumstellar medium is probably

generated by an Asymptotic Giant companion (Hamuy et al. 2003; Wang et al. 2004;

Han & Podsiadlowski 2006). However, there is still debate in the literature about the

origin of SN 2002ic. Livio & Riess (2003) proposed the merger of two white dwarfs

as a possible progenitor, with the explosion occurring in the common envelope phase.

Chugai et al. (2004) concluded that the properties of the circumstellar interaction

in 2002ic-like events can be broadly explained by the SN 1.5 scenario (Iben &

Renzini 1983): the thermonuclear explosion of the degenerate core of a massive AGB

star. Recently, Benetti et al. (2006b) questioned the earlier interpretation of the

observations and proposed that SN 2002ic can be equally well explained by the core

collapse of a stripped-envelope massive star in a dense medium.

SN 2005gj was discovered on 2005 September 28.6 (UT) by the SDSS-II

Supernova Collaboration (Frieman et al. 2008) in gri images obtained with the

SDSS-2.5m telescope at Apache Point Observatory (APO). The new supernova

(Barentine et al. 2005) was ∼1′′ from the center of its host galaxy at the

position α = 03h01m12.s0, δ = +0033′13.′′9 (J2000.0). It had SDSS magnitudes

(g, r, i) = (18.6, 18.6, 18.7) mag, obtained from PSF photometry after kernel-matching

and subtraction of a template image in each band. The SN was independently

discovered by the Nearby Supernova Factory on 2005 September 29 (Aldering et al.

2006).

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SN 2005gj was classified as a Type Ia candidate from the first three epochs of

the gri light curve using a light curve fitting program, and was sent to the queue

of the MDM-2.4m telescope for spectroscopic confirmation. The optical spectrum

obtained on 2005 October 1 (UT) showed a blue continuum with resolved Hydrogen

Balmer lines in emission, very similar to the spectrum of a young Type IIn supernova,

but also with an unusual continuum showing broad and weak absorption features.

Further spectroscopic follow-up showed a dramatic evolution. The continuum

became substantially redder and developed broad, P-Cygni features probably

associated with blended lines of Fe-peak mass elements, similar to a Type Ia SN

a few weeks after maximum. The spectrum obtained on 2005 Nov. 12 (UT) was

remarkably similar to that of the unusual Type Ia supernova SN 2002ic obtained on

2002 Dec. 27 (UT) (Prieto et al. 2005).

Aldering et al. (2006) presented optical photometry and spectroscopy of this

SN. Through detailed analysis they confirmed its photometric and spectroscopic

resemblance to SN 2002ic, confirming it was a new case of a Type Ia explosion

interacting with a dense circumstellar environment. From a spectrum obtained with

the slit oriented to overlap with its host galaxy, they calculated a redshift for the

host of z = 0.0616 ± 0.0002. SN 2005gj was not detected in the radio with the Very

Large Array (Soderberg & Frail 2005) or in the X-ray with Swift (Immler et al.

2005).

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Here we present extensive follow-up photometry and spectroscopy of the

Type Ia/Type IIn SN 2005gj during the first ∼150 days after discovery. These are

the most detailed observations ever obtained of a SN 2002ic-like event, and provide

insight into the early evolution, progenitor, and variety of these events. We also

present a sensitive, early X-ray observation with Chandra that gives an upper limit

on the X-ray luminosity of this peculiar object. We describe the optical and NIR

photometry of SN 2005gj in §6.2 and the optical spectroscopy in §6.3. We describe

the X-ray observation with Chandra in §6.4. An analysis of the photometric and

spectroscopic data are presented in §6.5. Finally, we discuss the significance of the

results in §6.6. We adopt a cosmology with H0 = 72 ± 8 km s−1 Mpc−1, ΩM = 0.3,

and ΩΛ = 0.7 throughout the paper (see Spergel et al. 2007), which yields a distance

modulus of µ = 37.15 mag to the host of SN 2005gj.

6.2. Photometry

6.2.1. SDSS and MDM

The Sloan Digital Sky Survey (SDSS) uses a wide-field, 2.5-meter telescope

(Gunn et al. 2006) and mosaic CCD camera (Gunn et al. 1998) at APO to survey

the sky. The SDSS-II Supernova Survey, part of a 3-year extension of the original

SDSS, uses the APO-2.5m telescope to detect and measure light-curves for a

large number of supernovae through repeat scans of the SDSS Southern equatorial

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stripe (about 2.5 deg wide by ∼120 deg long) over the course of three 3-month

campaigns (Sept-Nov. 2005-2007). SN 2005gj was discovered in the second month

of the first campaign (October 2005). Twenty epochs of ugriz photometry were

obtained between 2005 Sep 26-Nov 30 (U.T.). Details of the photometric system and

calibration are given in Fukugita et al. (1996), Hogg et al. (2001), Smith et al. (2002),

Ivezic et al. (2004), Tucker et al. (2006). Additional griz imaging of SN 2005gj was

obtained with the MDM Observatory 2.4m telescope using a facility CCD imager

(RETROCAM; see Morgan et al. 2005, for a complete description of the imager).

Photometry of SN 2005gj on the SDSS images was carried out using the scene

modeling code developed for SDSS-II as described in Holtzman et al. (2008). A

sequence of stars around the supernova were taken from the list of Ivezic et al.

(2007), who derived standard SDSS magnitudes from multiple observations taken

during the main SDSS survey under photometric conditions. Using these stars,

frame scalings and astrometric solutions were derived for each of the supernova

frames, as well as for the twenty five pre-supernova frames taken as part of either

the main SDSS survey or the SN survey. Finally, the entire stack of frames was

simultaneously fit for a single supernova position, a fixed galaxy background in each

filter (characterized by a grid of galaxy intensities), and the supernova brightness in

each frame.

The supernova photometry in the MDM frames was also determined using

the scene modeling code. Since the MDM observations had different response

174

functions from the standard SDSS bandpasses, the photometric frame solutions

included color terms from the SDSS standard magnitudes. To prevent uncertainties

in the frame parameters and color terms from possibly corrupting the galaxy model

(here affecting the SDSS photometry), the MDM data were not included in the

galaxy determination, but the galaxy model as determined from the SDSS was used

(with color terms) to subtract the galaxy in the MDM frames. The resulting SN

photometry from the MDM frames is reported on the native MDM system, since

the color terms derived from stars are likely not to apply to the spectrum of the

supernova.

Figure 6.1 shows a 3.5′ × 3.5′ field around SN 2005gj and 16 comparison stars

used for calibration of the SN magnitudes by SDSS and the Carnegie Supernova

Group (CSP; see §6.2.2 below for details). In Table 6.1 we present the SDSS ugriz

and CSP u′g′r′i′ photometry of the comparison stars in common. The final SDSS

and MDM griz photometry is given in Table 6.2.

6.2.2. CSP

Optical photometry from the CSP was obtained with the Swope-1m telescope at

LCO, using a SITe CCD and a set of u′g′r′i′ filters. A subraster of 1200× 1200 pixels

was read from the center of the CCD, which, at a scale of 0.435 ′′ pixel−1, yielded a

field of view of 8.7′×8.7′. Typical image quality ranged between 1′′ and 2′′ (FWHM).

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A photometric sequence of comparison stars in the SN field was calibrated with the

Swope telescope from observations of SDSS standard stars from Smith et al. (2002)

during four photometric nights. SN magnitudes in the u′g′r′i′ system were obtained

differentially relative to the comparison stars using PSF photometry. In order

to minimize the contamination from the host galaxy light in the SN magnitudes,

PSF-matched ugri template images from SDSS were subtracted from the SN images.

On every galaxy-subtracted image, a PSF was fitted to the SN and comparison

stars within a radius of 3′′. See Hamuy et al. (2006) for further details about the

measurement procedures.

The NIR photometry of SN 2005gj was obtained by the CSP using three

different instruments/telescopes. A total of 15 epochs in Y , J and H filters were

obtained using RetroCam, mounted on the Swope-1m telescope at LCO. A few

additional epochs in Y JHKs were obtained with the Wide Infrared Camera (WIRC;

Persson et al. 2002) mounted on the duPont-2.5m telescope, and the PANIC camera

(Martini et al. 2004) mounted on the Magellan-6.5m Baade telescope, both at LCO.

We refer to Hamuy et al. (2006) and Phillips et al. (2007) for details of the imagers

and procedures to extract the SN photometry. The host galaxy was not subtracted

from the NIR frames; therefore the SN photometry contains an unknown galaxy

contamination component.

The final CSP u′g′r′i′ photometry of SN 2005gj is given in Table 6.3 and

Y JHKs photometry in Table 6.4. A minimum uncertainty of 0.015 mag in the

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optical bandpasses and 0.02 mag in the NIR is assumed for a single measurement

based on the typical scatter in the transformation from instrumental to standard

magnitudes of bright stars (Hamuy et al. 2006).

6.3. Spectroscopy

The spectroscopic observations of SN 2005gj are summarized in Table 6.6.

They were obtained using five different telescopes/instruments at four observatories.

A total of twelve spectra were obtained between early-October 2005 and

late-January 2006 with the Boller & Chivens CCD Spectrograph (CCDS) mounted

at the MDM-2.4m telescope. This instrument uses a Loral 1200×800 pixel CCD

with 15 µm pixel−1 and a 150 l/mm grating (blazed at 4700 A). We used a 2′′ slit

which gives a dispersion of 3.1 A pixel−1 in the wavelength range ∼3800–7300 A.

Eight spectra were obtained at LCO with the Wide Field Reimaging Camera

(WFCCD) and the Modular Spectrograph (ModSpec) at the duPont-2.5m telescope,

and LDSS-3 at the Clay-6.5m telescope. In the case of WFCCD a 400 l/mm grism

and a 1.6′′ slit were used, reaching a dispersion of 3 A pixel−1 in the wavelength

range ∼3800–9200 A. For ModSpec we used a 300 l/mm grating and a 1′′ slit that

gave a dispersion of 2.45 A pixel−1 in the range ∼3800–7300 A. Further details of

WFCCD and ModSpec can be found in Hamuy et al. (2006). For LDSS-3, which

employs a STA0500A 4064× 4064 pixel CCD, we used the VPH blue and red grisms,

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the latter with an OG590 filter to block second order contamination, with a 0.75′′

slit reaching a dispersion of 0.70 A pixel−1 and 1.12 A pixel−1 in the blue and red

sides of the spectrum, respectively.

Additionally we obtained two early spectra with the Double Imaging

Spectrograph (DIS) mounted on the ARC-3.5m telescope at APO, and one spectrum

with the Intermediate dispersion Spectrograph and Imaging System (ISIS) at the

WHT-4.2m telescope at the Roque de Los Muchachos Observatory in La Palma,

Spain. The DIS spectrograph has blue and red detectors, each uses a Marconi

2048 × 1024 pixel CCD with 13.5 µm pixel−1. We used a 300 l/mm grating and a

1.5′′ slit which gives a dispersion of 2.4 A pixel−1. The ISIS spectrograph has a blue

and red arm, the blue arm using a EEV12 CCD and the red using a MARCONI2

detector. We used a 300 l/mm grating in both the blue and red arm and a 1′′ slit,

which gives a dispersion of 0.86 A pixel−1 in the blue and 1.47 A pixel−1 in the red.

Most of the spectra were obtained close to the parallactic angle to minimize

relative changes in the calibration of the blue and red parts of the spectrum due

to differential refraction through the atmosphere. The spectroscopic reductions

were performed using standard IRAF tasks and included: bias and overscan

subtraction, flat-fielding, combination of 2-4 individual 2D spectra to reach the best

signal-to-noise ratio in the final image, tracing and extraction of a 1D spectrum

from the combined 2D image, subtracting the background sky around the selected

aperture, wavelength calibration using an arc-lamp, and flux calibration. In order

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to flux calibrate the spectra we observed 1-2 spectrophotometric standard stars per

night. The spectra from LCO, WHT and MDM were corrected by atmospheric

telluric lines using the spectrum of a hot spectrophotometric standard star and the

spectra from APO were corrected using a model atmosphere. This correction is not

optimal for some of the spectra and there are evident residuals left in the corrected

spectra.

Figure 6.2 and 6.3 show a montage of the optical spectra of SN 2005gj obtained

from October 2005 to March 2006. We have split them in two figures to avoid

crowding. The position of the most conspicuous spectral features have been indicated

in this figure.

6.4. X-ray Observation

SN 2005gj was observed under Director’s Discretionary Time for 49.5 ks on

2005 Dec 11/12 (ObsID 7241) with the Chandra X-ray Observatory’s Advanced

CCD Imaging Spectrometer (ACIS). The data were taken in timed-exposure mode

with an integration time of 3.2 s per frame, and the telescope aimpoint was on the

back-side illuminated S3 chip. The data were telemetered to the ground in “very

faint” mode.

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Data reduction was performed using the CIAO 3.3 software provided by the

Chandra X-ray Center1. The data were reprocessed using the CALDB 3.2.2 set

of calibration files (gain maps, quantum efficiency, quantum efficiency uniformity,

effective area) including a new bad pixel list made with the acis run hotpix tool.

The reprocessing was done without including the pixel randomization that is added

during standard processing. This omission slightly improves the point spread

function. The data were filtered using the standard ASCA grades (0, 2, 3, 4, and 6)

and excluding both bad pixels and software-flagged cosmic ray events. Intervals of

strong background flaring were searched for, but none were found.

Absolute Chandra astrometry is typically good to ∼0.′′5, and we sought to

improve it by registering the Chandra image with an SDSS image. Chandra point

sources were found using the wavdetect tool, and their positions were refined using

ACIS Extract version 3.101. Fourteen X-ray sources had SDSS counterparts, which

we used to shift the Chandra frame by a small amount (0.′′15 in RA and 0.′′07 in

DEC). After the shift, the residual differences between the Chandra and SDSS

sources had rms values of 0.′′19 in RA and 0.′′12 in DEC.

We extracted counts in the 0.5–8 keV bandpass from the position of the

supernova using a standard extraction region (∼1′′ radius), and we constructed

response files with the CAIO tools and ACIS Extract. The background region is a

source-free annulus centered on the position of the supernova with inner and outer

1http://asc.harvard.edu

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http://asc.harvard.edu

radii of 6′′ and 32′′. Based on the 300 photons detected in this region, we expect 0.3

background counts in our source extraction region.

We detect only two counts from the location of the supernova, but neither may

be associated with the supernova itself. The counts had energies of 4.0 keV and

6.5 keV, but one would expect some emission in the 0.5–2.5 keV range since this is

where Chandra has the most collecting area. For example, the average effective area

in each of the 0.5–2.5 keV, 4.0–6.0 keV, and 6.0–8.0 keV bands is 470 cm2, 270 cm2,

and 90 cm2, respectively.

We calculate our upper limits using the Bayesian method of Kraft et al. (1991).

For the 0.5–4 keV band, the 68% (95.5%) upper limit to the source counts is 1.14

(3.05). For the 0.5–8 keV band, the 68% (95.5%) upper limit is 3.52 (6.14).

Since no Type Ia supernova has been conclusively detected in X-ray, we have

no a priori expectation of the spectral shape. We therefore adopt a simple absorbed

power law with a photon index of 2 and an absorbing column of nH = 7.08 × 1020

cm−2. For this choice of spectrum, the count rate to flux conversion is 5.2 × 10−12

erg/count in the 0.5–4 keV band and 6.9 × 10−12 erg/count in the 0.5–8 keV band.

We therefore arrive at 68% (95.5%) upper limits on the X-ray luminosity of

1.1 (2.9)× 1039 erg s−1 in the 0.5–4 keV band and 4.3 (7.6)× 1039 erg s−1 in the 0.5–8

keV band. Based on the above statements concerning the Chandra effective area, we

feel the 0.5–4 keV limit is a more appropriate limit.

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6.5. Results

6.5.1. Optical light curves and colors

Figure 6.4 shows the early SDSS and MDM ugriz light curves combined with

the late time coverage given by CSP u′g′r′i′Y JHKs photometry. They give excellent

multi-wavelength optical and NIR coverage and sampling of the first ∼150 days after

discovery (∼140 rest-frame days after the time of explosion).

In Table 6.5 we give important parameters derived from the light curves. We

have a good estimate of the time of explosion of the SN at JD 2,453,637.93±2.02

(September 24.4 UT, 2005) calculated as the average time between the epoch when

the supernova is detected at > 5σ in all filters at JD 2,453,639.94, and the last

pre-discovery observation of the field at JD 2,453,635.91. The times and observed

magnitudes at maximum in different filters are presented in Table 6.5. They are

calculated from a high order polynomial fit to each light curve. To estimate the

errors we assume a gaussian distribution for the magnitude uncertainty at each

epoch and filter (assuming they are not correlated from epoch to epoch), then we

draw randomly ∼1000 simulated light curves and fit each one with the same high

order polynomial. The 1σ uncertainties are taken as the rms deviation of the mean

values calculated from each simulated light curve.

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The risetimes, defined as the time between explosion and maximum light,

become longer at redder wavelengths. They are 13.5, 19.7, 33.7, and 46.9 days in

ugri filters, respectively, which correspond to 12.7, 18.5, 31.7, and 44.2 days in the

rest-frame of the supernova.

We also give in Table 6.5 the rest-frame magnitudes at maximum in different

filters. They have been corrected by Galactic extinction in the line of sight, using

E(B − V )Gal = 0.121 (SFD: Schlegel et al. 1998) and the Cardelli et al. (1989)

(CCM) reddening law assuming RV = 3.1, and K-corrections (see below), which are

not negligible in this object since it is at redshift z ∼ 0.062. We assume that the

host galaxy extinction is negligible (Aldering et al. 2006).

K-corrections have been calculated from the multi-epoch spectra presented in

§6.3. In order to estimate accurate K-corrections we need good spectrophotometric

calibration in all the wavelength range. Figure 6.5 shows the differences between the

observed g − r photometric colors obtained from the light curves and the synthetic

colors calculated directly from the spectra, as a function of the observed photometric

colors. We use an 8th order polynomial fit of the light curves to obtain the observed

g − r color at the epoch of a given spectrum to better than ∼ 0.03 mag. We can see

that most of the spectra have good spectrophotometric calibration in the wavelength

range of g , r filters (3800–7000 A), with a residual color of ∼0.05 mag rms (difference

between observed and synthetic colors), but there are some obvious outliers. To

correct the spectra and produce a better spectrophotometric calibration consistent

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with the observed spectral energy distribution (SED) obtained from broadband

photometry we warp the spectra multiplying by a smooth function to match the

observed colors, a technique commonly used for calculating K-corrections in Type Ia

SN (Nugent et al. 2002).

First, we extrapolate the continuum in the blue and red sides of each spectrum

presented in §6.3 with a low order polynomial to have complete coverage of the SDSS

ugriz filters (2000–11000 A). Using this extended version of the spectra we apply the

CCM reddening law iteratively until the synthetic g − r color matches the observed

color in each spectrum. This procedure does not ensure that the calibration is good

in the complete wavelength range; therefore we multiply by a smooth spline with

knots at the effective wavelength of the SDSS filters until the synthetic u − g, g − r,

r − i and i− z colors match the observed colors obtained, again by using polynomial

interpolation of the light curves.

The K-corrections for the same filter (Hamuy et al. 1993) calculated from the

modified, spectrophotometrically calibrated spectra using SDSS passbands are listed

in Table 6.7. The K-corrections are probably accurate to ±0.05 mag for g and r

filters, and to ±0.1 mag for u, i and z where we had to extrapolate the spectra.

However, our estimate is not precise and we can not exclude the possibility of

even larger errors. After fitting with a low order polynomial we use the results to

transform the observed SDSS magnitudes in Table 6.2 to the rest-frame.

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We corrected the CSP u′g′r′i′ magnitudes to rest-frame SDSS ugri magnitudes

using cross-filter K-corrections according to the prescription of Nugent et al. (2002).

These take into account the difference between the CSP and SDSS passbands

convolved by the SED of SN 2005gj and allow us to put all the rest-frame magnitudes

in the same system. We find that the differences between the same- and cross-filter

K-corrections are small (∼< 0.03 mag) at all epochs.

Figure 6.6 shows the evolution of the rest-frame colors of SN 2005gj as a

function of time after explosion. We have corrected the magnitudes by K-corrections

and Galactic extinction in the line of sight. As a comparison we also plot the color

evolution of the overluminous Type Ia SN 1991T, a typical Type IIn SN 1999el (Di

Carlo et al. 2002) (both obtained from spectral templates of P. Nugent 2 ), and two

previous cases that are thought to be Type Ia explosions in a very dense environment:

SN 2002ic and SN 1997cy. For SN 2002ic we use the published BV I photometry

(Hamuy et al. 2003) corrected by Galactic extinction (E(B − V ) = 0.06; SFD).

We calculate cross-filter K-corrections using the calibrated spectra of SN 2005gj

to transform observed magnitudes of the SN at z = 0.0667 (Kotak et al. 2004) to

rest-frame SDSS magnitudes: B → u, V → g, I → i. The time of explosion of

SN 2002ic is assumed to at ≈JD 2,452,581.5 (2002, Nov. 3 UT; Deng et al. 2004).

We obtain rest-frame colors in the SDSS system for SN 1997cy (Germany et al. 2000;

Turatto et al. 2000) by transforming the K-corrected V RI magnitudes in Germany

2http://supernova.lbl.gov/∼nugent/nugent templates.html

185

et al. (2000) to gri magnitudes using S-corrections calculated directly from spectra

of SN 1997cy available online in the SUSPECT database 3. We supplement this with

synthetic colors from the spectra. We also correct by Galactic extinction in the line

of sight (E(B − V ) = 0.02; SFD) and assume the time of explosion of SN 1997cy to

be JD 2,450,582.5 (1997, May. 14 UT; Germany et al. 2000), which is very uncertain,

since it is taken as the time of detection of the gamma-ray burst GRB 970514, which

may not have been associated with the SN. Magnitudes for SN 2002ic and SN 1997cy

are not corrected by extinction in their host galaxies, which is unknown.

Initially the evolution in rest-frame u − g and g − r colors of SN 2005gj,

up to ∼30-40 days after explosion (∼10-20 days after maximum for a SN Ia), is

roughly consistent with the colors of SN 1991T but ∼ 0.2 mag redder in g − r,

and evolves to redder colors at later times. SN 1991T reaches its maximum colors

of (u − g)=1.5 mag and (g − r)=1.0 mag at ∼50 days (30 days after maximum),

and after that it enters the nebular phase and becomes bluer. At late times the

u − g color of SN 2005gj has a slow linear increase and becomes systematically

redder than SN 1991T at > 70 days, while the g − r color stays approximately

flat at (g − r)=0.5 mag (and bluer than SN 1991T) between 60 − 110 days. The

evolution in u − g color of SN 2002ic is very similar and consistent with SN 2005gj.

SN 1997cy has a similar color evolution but it is ∼0.2 mag bluer in g − r color

between 60 − 100 days. However, Germany et al. (2000) give K-correction errors of

3http://suspect.nhn.ou.edu/∼suspect/

186

∼0.15 mag and we have further applied S-corrections, therefore this does not imply a

significant difference between SN 1997cy and the other two Type Ia/IIn supernovae.

The rest-frame r − i and i − z color evolution of SN 2005gj is very different

from SN 1991T and closely follows the evolution of a SN IIn. The r − i colors of

SN 1997cy are also consistent with SN 2005gj. From these comparisons it is clear

that the colors, a proxy for the temperature of the photosphere, of SN 2005gj and

two earlier cases of SN Ia strongly interacting with their circumstellar medium are

dominated by the radiation coming from the ejecta-CSM interaction.

In Figure 6.7 we present the ugri light curves of SN 2005gj in absolute

magnitudes. We also show the light curves of SN 1991T, SN 2002ic and SN 1997cy

obtained from the literature and corrected to SDSS rest-frame magnitudes as

explained above.

SN 2005gj has peak ugriz absolute magnitudes in the range −20.0 to

−20.3 mag, this is ∼ 0.7 − 1.4 mag brighter than the overluminous Type Ia

SN 1991T. The u light curve is consistent with a linear decline after peak

luminosity at a constant rate of 0.027 ± 0.001 mag day−1. The g light curve has a

∼ 20 day plateau with roughly constant luminosity after maximum (20 − 40 days

after explosion), then the light curve declines linearly between 40 − 100 days at

0.018 ± 0.001 mag day−1 and continues its linear decay at later times, but with a

shallower slope of 0.007 ± 0.002 mag day−1. The r and i light curves have a similar

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plateau shape between 20 − 60 days and a constant linear decay at later times of

0.013± 0.001 mag day−1. The change in slope observed in the g-band at late times is

less clear in ri, but still present. The secondary maximum present in ri light curves

of SN 1991T and other SN Ia is completely absent in SN 2005gj.

The light curves of SN 2002ic are fainter than SN 2005gj at all times by

0.3−0.6 mag depending on the filter (0.3−0.8 mag brighter than SN 1991T at peak).

The initial decline rates from maximum of the u and g light curves of SN 2002ic are

intermediate between those of SN 1991T and SN 2005gj until around 40 days after

explosion; after day 50 when the ejecta-CSM interaction had become dominant in

SN 2002ic (Hamuy et al. 2003), they closely resemble the decline rates of SN 2005gj

in the same bands. The i band light curve of SN 2002ic showed definite evidence

for a weak secondary maximum, which is again intermediate in morphology between

the strong secondary maximum observed in SN 1991T, and the absence of such a

feature in SN 2005gj. The gr light curves of SN 1997cy are consistent with linear

decay of ∼ 0.008 mag day−1, being ∼ 0.4 mag brighter than SN 2005gj at 150 days.

Unfortunately, the light curves of SN 1997cy start at ∼ 60 days after explosion and

there is no information near peak to compare with the luminosities of SN 2005gj

and SN 2002ic. However, if we extrapolate the gr light curves to the time of peak

luminosity we find that the absolute magnitudes at maximum of SN 1997cy would

be within ∼ 0.2 mag of SN 2005gj. This has to be taken with caution because of the

extrapolation and uncertain time of explosion of SN 1997cy.

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6.5.2. NIR light curves

We used the CSP JHKs photometry obtained between 59-166 days after the

explosion (rest-frame) to construct the absolute magnitude light curves of SN 2005gj.

The observed magnitudes were corrected by Galactic reddening in the line of sight

(AJ = 0.108, AH = 0.070, and AKs = 0.044 mag) and K-corrections. We calculated

K-corrections for the same filter using the spectral templates of P. Nugent for the

Type IIn SN 1999el (which are derived from black-body curve fits to the photometry),

because as we showed in §6.5.1 the synthetic colors obtained from the spectral

templates approximate reasonably well the evolution of the redder optical colors of

SN 2005gj (i.e., r − i and i − z). The values of the K-corrections are consistent

with being constant in this time range: KJ ≈ −0.12 mag ,KH ≈ −0.14 mag,

KKs≈ 0.16 mag.

The JHKs absolute magnitude light curves of SN 2005gj are presented in

Figure 6.8. For comparison we show the NIR light curves of a normal Type Ia

obtained from synthetic photometry of spectral templates and the Type IIn SN 1999el

(Di Carlo et al. 2002). The Type Ia light curves have been shifted in magnitudes to

match the mean absolute magnitude at maximum of SN Ia (Krisciunas et al. 2004).

SN 2005gj is 1.7 − 3 mag brighter than a normal SN Ia and SN IIn at 60 days after

the explosion and declines in a slower fashion at later times. A linear decay is a good

fit in all NIR bandpasses at this late times (> 60 days after explosion), with decline

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rates of ∼ 0.014 mag day−1 in J , ∼ 0.013 mag day−1 in H, and ∼ 0.011 mag day−1

in Ks. These values are similar to the decline rates in the optical ri bands.

Since there are no template images of the host galaxy obtained in the NIR

bands, the light curves are preliminary and the analysis has to be taken with caution.

6.5.3. Bolometric light curve

The SDSS and CSP magnitudes were used to produce a quasi-bolometric

light curve of SN 2005gj covering the optical wavelength range from

3000− 10, 000 A(u → z). We corrected the magnitudes by Galactic extinction in the

line of sight and K-corrections to obtain magnitudes in the rest-frame ugriz filters

(see §6.5.1 for details). We applied small corrections to transform the magnitudes

based in the SDSS photometric system into the AB system obtained from the SDSS

website 4. The AB magnitudes derived in this way are transformed directly to

bandpass averaged fluxes using the definition of the AB system (Oke & Gunn 1983)

and they are assigned to the frequencies that correspond to the effective wavelengths

of the SDSS ugriz filters, calculated from the filters using the definition in Fukugita

et al. (1996): 3567, 4735, 6195, 7510, and 8977 A. We use the trapezoidal rule to

obtain the integrated flux from u to z, this is between λ1 = 3340 A and λ2 = 9596 A,

where the limits of the wavelength coverage are obtained from λ1 = λeff,u − ∆λu/2

4http://www.sdss.org/dr5/algorithms/fluxcal.html.

190

and λ2 = λeff,z + ∆λz/2. We extrapolate linearly the u and z light curves at late

times to fill in the lack of coverage of the SDSS in z, and CSP in u and z bands,

including the MDM z-band data at these epochs. The integrated fluxes are converted

to luminosity assuming a luminosity distance to SN 2005gj of 268.5 Mpc and a

spherically symmetric distribution of the output energy. We present the integrated,

quasi-bolometric luminosities from u to i (L(u→i)) and from u to z (L(u→z)) as a

function of time in Table 6.8.

In order to estimate bolometric UVOIR luminosities we calculate time dependent

bolometric corrections to include the energy output of the SN at wavelengths bluer

than u-band (λ < 3340 A) and redder than z-band (λ > 9596 A). We find that

black-body distributions with different temperatures are a reasonable approximation

of the spectral energy distribution of SN 2005gj. We use χ2-minimization to fit

the optical SED with a two parameter black-body function: temperature and a

multiplicative scaling factor. The scaling factor is a combination of fundamental

constants and the square of the angular radius of the spherical black-body: (Rbb/d)2,

where Rbb is the radius of the black-body and d is the distance to the SN. We

calculate time-dependent bolometric corrections by integrating the black-body

distributions in the ultraviolet and NIR/IR regions, and converting the integrated

fluxes to luminosities as explained above. For the CSP data we also include NIR

flux densities derived from the reddening and K-corrected Y JHKs magnitudes after

fitting the light curves with low order polynomials.

191

At early phases before peak the bolometric corrections account for ∼ 53 − 65%

of the total integrated luminosity, of which 85 − 93% is from the ultraviolet part

of the spectrum and only 7 − 15% from the NIR/IR. As the supernova evolves,

the ejecta expands and shocks the circumstellar gas. The energy emitted in the

ultraviolet/blue part of the spectrum declines quickly after maximum light and most

of the energy is emitted in the optical region, coinciding with the appearance in the

spectrum of emission/absorption features of the intermediate and iron-peak elements

(see §6.5.4). Between 25 − 150 days after explosion the bolometric corrections

account for ∼ 32 − 45% of the total output luminosity, with the NIR/IR correction

dominating completely over the blue/ultraviolet at > 60 days.

The bolometric UVOIR luminosities, black-body temperatures and radii

derived from the fits are presented in Table 6.8. The uncertainties in black-body

temperatures and radii are calculated using the diagonal terms of the covariance

matrix obtained from the χ2 minimization. We add a 10% error in the distance

to the SN to the error in the black-body radii, which comes from the random

and systematic uncertainties in the value of the Hubble constant (Freedman et

al. 2001). The uncertainties in the bolometric luminosities were estimated by

propagating errors through the trapezoidal integration of the SED, taking into

account: uncertainties in the photometry, light curve interpolation and fitting,

Galactic extinction, K-corrections, and distance to the SN. To approximately take

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into account the errors introduced by the bolometric corrections we multiplied these

values by√

χ2ν when the reduced χ2 is greater than 1.

In Figure 6.9 we show some examples of black-body fits to the optical SED at

different epochs. At early times, shortly after explosion, the SED is very well fit by a

hot ∼13000 K black-body. The temperature starts to decrease steadily close to the

time of explosion to a constant value of ∼ 6500 K at 60 days after. A black-body

is still a reasonable approximation of the SED at later times, but the fits become

poorer when emission/absorption features start to dominate the spectrum, which is

represented in the χ2ν of the fits (see Table 6.8).

Figure 6.10 shows the bolometric light curve of SN 2005gj in the top panel

and the evolution in temperature and radius from the black-body fits in the lower

panels. The early data of the bolometric light curve are well fit by an exponential

rise in luminosity, L(t) ∝ e0.17 t. The time of maximum bolometric luminosity

occurs between 6.6–18.8 days after explosion. After maximum, the bolometric light

curve is very well approximated by an exponential decay in luminosity, linear in

the logarithmic scale shown in Figure 6.10, L(t) ∝ e−0.013 t (0.014 mag day−1). This

is consistent with the exponential density distribution of the ejecta of Type Ia

SN (Dwarkadas & Chevalier 1998), whereas the distribution of ejecta around

core-collapse supernovae is better approximated by a power-law (Chevalier &

Fransson 2003). Extrapolating the pre- and post-maximum fits we get a maximum

bolometric luminosity of Lmaxbol = 5.6× 1043 ergs−1, which is ∼3 times more luminous

193

than the Type Ia SN 1991T at maximum light (Contardo et al. 2000; Stritzinger

et al. 2006). Assuming a bolometric correction of 50% at maximum, we find that

SN 2005gj was ∼1.5 times more luminous than SN 2002ic.

6.5.4. Optical spectroscopy

In Figure 6.11 we show a comparison of the spectra of SN 2005gj with spectra

of SN 2002ic and SN 1997cy obtained at similar times after explosion. The spectra

of SN 2005gj and SN 2002ic are very similar at all times. They are characterized

by strong and broad Hydrogen-Balmer lines Hα and Hβ in emission5 and a blue

continuum at early times that becomes redder and increasingly dominated by

absorption/emission P-cygni profiles from Fe-peak ions (e.g., Fe II, Fe III, Ni III,

Si II, S II).

Classification

Benetti et al. (2006b) proposed that SN 2002ic-like events are well explained by

the core-collapse of a massive star in a dense medium, casting doubt in the previous

classification of SN 2002ic as a Type Ia supernova. The authors found relatively

good agreement at all times between the spectra of SN 2004aw (Taubenberger et al.

2006), a Type Ic supernova, and SN 2002ic.

5As shown in Figure 6.2, Hγ is also visible in the early spectra of SN 2005gj.

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We used the SuperNova IDentification code, SNID (Matheson et al. 2005;

Miknaitis et al. 2007; Blondin & Tonry 2007b), to find the spectra that best match

SN 2005gj at different epochs. SNID cross-correlates an input spectrum with a

library of supernovae spectra. In the library we included spectra of 5 normal SN Ia,

two 1991T-like objects, two 1991bg-like objects, 4 broad-lined SN Ic (or hypernovae),

and 3 normal SN Ic (including SN 2004aw), that were chosen to span a wide range

of observed properties of SN Ia and SN Ic. In Table 6.9 we present the supernovae

and the epochs of the spectra in the library. We fixed the redshift of SN 2005gj at

z = 0.0616 and allowed for a range around the mean redshift of ∆z = 0.02 to find

cross-correlation peaks. The Balmer lines in emission were clipped from the input

SN 2005gj spectra to avoid spurious cross-correlation signal with library spectra that

contain emission lines from the host galaxy.

Figure 6.12 shows the spectra of SN 2005gj at four epochs and the library

spectra with the highest cross-correlation significance from SNID. Type Ia supernovae

spectra are a better match to SN 2005gj at most epochs, with 20 of the 26 (77%)

epochs having best matched a type Ia spectrum (45%-91T, 45%-normal, 10%-91bg)

and a similar fraction for the next-best matches. The broad-lined SN Ic 1997ef and

2002ap are the best matching spectra for 6 epochs, all of them between 26–46 days

after explosion. We repeated the same procedure using five spectra of SN 2002ic

obtained between 24–84 days after explosion (Hamuy et al. 2003). All the spectra of

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SN 2002ic are better matched with SN Ia, in contradiction with the results obtained

by Benetti et al. (2006b).

The continuum of SN 2005gj is well approximated at all times by the sum of

a scaled spectrum of the overluminous Type Ia SN 1991T, at the same epoch after

explosion as SN 2005gj, and a fourth order polynomial. A normal SN Ia does not fit

as well as SN 1991T. This procedure is very similar to the fits to SN 2002ic (Hamuy

et al. 2003) and SN 2005gj (Aldering et al. 2006) presented in previous studies. In

Figure 6.13 we show examples of the spectra decomposition at four epochs. We

excluded from the fit a region of ±100 A around the Hα and Hβ lines and obtained

a good fit for the remainder of the spectrum.

Balmer lines

We analyzed the Balmer emission features in the spectra using the sum of

two Gaussian components to model the line profiles. This decomposition gives

much better fits for Hα at all epochs than a single Gaussian and it is physically

motivated (Chugai 1997b,a). The spectra of Type IIn SN show Balmer features with

both a narrow and broad component that can be explained by radiation coming

from different regions of the ejecta/CSM, whether it is direct emission from the

shock-heated CSM (broad component) or emission from un-shocked gas photoionized

by the SN radiation (narrow component). The Hβ line is unresolved or only

marginally resolved for most of the spectra. Therefore a single Gaussian component

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was used to fit the line profile. We used a third order polynomial to model the local

continuum around each line that was included in the Gaussian fits. It is important

to stress that at late times there is a broad Fe II feature intrinsic to the supernova

spectrum in the region of Hα (see spectra in Figure 6.13) that makes the definition

of the continuum less reliable and may affect the line measurements.

The results of the Gaussian fits to the Hα and Hβ emission features, integrated

fluxes and FWHM, are shown in Table 6.10 as a function of epoch of the spectra.

We have excluded the two spectra with better resolution because they show P-Cygni

profiles (see below). We used the flux calibrated spectra corrected to match

the observed g, r magnitudes (as explained in §6.5.1) and corrected for Galactic

reddening in the line of sight. The FWHM of the Gaussian profiles listed in

Table 6.10 were corrected by the resolution of the spectrographs (from Table 6.6).

We do not present the values when the component is unresolved. The errors quoted

for the integrated fluxes are obtained by adding in quadrature an estimate of 10%

error assigned to the absolute flux calibration and the rms deviation of the Gaussian

fits.

Figure 6.14 shows the evolution in time of the FWHM (top left panel), Hα

and Hβ luminosities (top right and bottom left panels) and the Balmer decrement

(bottom right panel). The FWHM of Hα varies between ∼ 130 − 500 km s−1

(narrow component) and ∼ 1200 − 3800 km s−1 (broad component), with the broad

component showing a slow increase in time. The FWHM of Hβ varies between

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∼ 470 − 1700 km s−1 and does not show evident evolution. The luminosities of

Hα-narrow and Hβ lines evolve in a similar fashion, increasing at early times to

peak at ∼ 12 days with luminosities 6 − 6.5 × 1040 ergs−1, then they decay and stay

roughly constant after 50 days. The evolution of Hα-broad is similar during the

first 50 days, peaking at 1.1 × 1041 ergs−1, but it shows an increase at later times.

Compared with the Hα luminosities observed in SN 2002ic, both components are

∼ 4 times more luminous.

The Balmer decrement, the ratio between Hα (sum of narrow and broad

components) and Hβ fluxes, stays approximately constant during the first 30−40 days

(mean = 2.5 and rms = 0.5) and is consistent with the theoretical value in Case

B recombination of Hα/Hβ = 2.86 (Osterbrock 1989). At later times it shows an

steady increase, reaching Hα/Hβ ∼ 7 − 13 at ∼ 80 days. In Case B recombination a

Balmer decrement Hα/Hβ > 2.86 is usually interpreted as evidence for the presence

of internal extinction in the host; however, the large values observed at late times

would produce an Na I D interstellar absorption doublet easily detectable in the

spectra, which is not observed (see also Aldering et al. (2006)), and the evolution in

time is not expected. Aldering et al. (2006) proposed that the H level populations

are in Case C recombination, where the optical depth in the Hα line is high implying

high densities and greater importance of collisional processes. In this scenario, the

observed change in the Balmer decrement could indicate that collisional excitation

becomes increasingly important at later times (Branch et al. 1981; Turatto, et al.

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1993). SN 2002ic (Deng et al. 2004) and other SN IIn, like SN 1988Z (Aretxaga

et al. 1999) and SN 1995G (Pastorello et al. 2002), have also shown large values of

the Balmer decrement and therefore may have similar physical processes affecting

the formation of the Balmer lines.

In Figure 6.15 we show the regions around Hα and Hβ features in the best

resolution spectra from ISIS and LDSS-3, obtained at 44 and 115 days after explosion,

respectively. We clearly detect P-Cygni profiles in all these features, which indicates

the presence of an outflow moving at ∼150-200 km s−1; however, these measurements

are limited by the resolution of the spectra between ∼130-180 km s−1 (FWHM).

After correcting for the resolution we obtain an outflow velocity of 60-70 km s−1.

The detection of P-Cygni-like absorption rules out an H II region in the line of sight

that could be producing the narrow emission/absorption features; the line profiles

are intrinsic to the SN. Aldering et al. (2006) detected P-Cygni profiles in He I,

Hα and Hβ, in a high resolution spectrum obtained with LRIS+Keck 71 days after

the explosion. They derived a wind velocity of vw ≈ 60 km s−1 consistent with our

estimate. Kotak et al. (2004) also detected a P-Cygni profile in the a spectrum of

SN 2002ic obtained 256 days after explosion.

Line identification near maximum

We used the parameterized resonance-scattering code SYNOW (Fisher et al.

1997; Fisher 2000) to identify the lines in the spectra obtained near maximum light

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of SN 2005gj. SYNOW is a fast supernova spectrum-synthesis code used for direct

(empirical) analysis of supernova spectra, mainly to identify the lines, their formation

velocities and optical depths. The code is based on simple assumptions: spherical

symmetry, homologous expansion, a sharp photosphere that emits a black-body

continuous spectrum, and line formation by resonance-scattering, treated in the

Sobolev approximation. We have used the latest version of the code that includes a

Gaussian distribution of optical depths.

Figure 6.16 shows the spectrum of SN 2005gj at 17 days after explosion (2 days

before g maximum) and the best synthetic spectrum obtained with SYNOW. We

also show for comparison the spectrum of SN 1991T obtained at -3 days with

respect to the time of B maximum. The spectra have been locally normalized

as in Jeffery et al. (2006). The synthetic spectrum has a black-body continuum

temperature Tbb = 11000 K, photospheric velocity vphot = 10000 km s−1, and

excitation temperature Texc = 10000 K. We find a reasonably good match with

the spectrum of SN 2005gj using the following lines/multiplets: Fe III λ4404 and

λ5129, Si III λ4561, Ni III, S II λ5468 and λ5633, and Si II λ6355. These lines are

characteristic of the overluminous and spectroscopically peculiar, Type Ia SN 1991T

around maximum light with strong Fe III features and weak S II doublet and Si II

(Jeffery et al. 1992; Mazzali et al. 1995; Fisher et al. 1999).

The main discrepancy between the SYNOW modeling of SN 2005gj and

SN 1991T is in the optical depths of the lines. The fit to SN 2005gj needs

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unphysically small optical depths, approximately 1/10th of the values used for

SN 1991T around maximum light. We interpret this as an effect of the extra

continuum radiation that is added by the ejecta-CSM interaction, which is veiling

(Branch et al. 2000) the supernova lines (e.g., Hamuy et al. 2003; Aldering et al.

2006). This interpretation is supported by the good agreement obtained from fitting

the spectra of SN 2005gj using a simple polynomial continuum added to the spectra

of SN 1991T at the same epochs after explosion (see Figure 6.13).

6.6. Discussion

We have presented extensive spectroscopy and optical/NIR photometry of

SN 2005gj obtained by the SDSS-II and CSP supernova groups during the first

∼150 days after explosion, and also an X-ray observation at 74 days that gives

an upper limit on the X-ray luminosity. We have shown the remarkable similarity

in spectroscopic and photometric properties between SN 2005gj and SN 2002ic,

which is thought to be the first clear case of a thermonuclear supernova explosion

embedded in a dense CSM. The observational properties of SN 2005gj support this

interpretation, they are summarized as follows:

• Spectroscopic evidence for a shock propagating into an Hydrogen-rich medium

close to the site of the explosion inferred from the presence of Balmer lines with

narrow (FWHM∼ 200− 500 km s−1) and broad (FWHM∼ 1500− 3000 km s−1)

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components at all times. The Balmer lines show P-Cygni profiles in the highest

resolution spectra obtained at 44 and 115 days after explosion, these detections

show the presence of a slow (∼ 100 km s−1) moving outflow. Both observations

support the interpretation of the supernova ejecta interacting with a dense

circumstellar material.

• Spectrum evolves from a very blue continuum (13000 K black-body) similar

to SN IIn at ∼7 days after explosion to a redder continuum at later times

with P-Cygni absorption/emission profiles. The strongest lines present around

maximum are identified with singly and doubly ionized iron-peak elements

(especially strong Fe III, weak S II and Si II) and the spectra are well matched

by the overluminous Type Ia SN 1991T diluted with a polynomial continuum

at similar times after explosion.

• Very luminous and slowly declining bolometric light curve. The linear decay

in luminosity after peak (∼ 0.014 mag day−1) suggests an exponential density

distribution of the ejecta, which is consistent with the ejecta-density profiles

obtained from simulations of SN Ia.

The data presented here on SN 2005gj makes the interpretation of 2002ic-like

events as thermonuclear supernovae in a dense CSM, initially proposed by Hamuy

et al. (2003), stronger. In contrast with Benetti et al. (2006b), we find that the

overall shape of the spectra of SN 2005gj are more consistent with spectra of SN Ia

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at different epochs. Specifically, Type Ic SNe usually do not show the S II doublet

at ∼ 5500 A around maximum light; in fact SN 2004aw shows only two very weak

notches at the wavelengths of S II near maximum (Taubenberger et al. 2006). This

is one of the identifying features in SN Ia spectra, also present in the overluminous

SN 1991T (Phillips et al. 1992). In the spectrum of SN 2005gj obtained at 17 days

(see Figure 6.16) we detect a weak double absorption that we identify with S II,

that is much stronger in the spectrum of SN 2002ic around maximum light. We

can see on the top of Figure 6.11 that the spectrum of SN 2002ic obtained 24 days

after explosion clearly shows this feature. Other conspicuous features observed in

SN 2005gj and SN 2002ic around maximum are Fe III and Si II. These features are

present in SN 1991T, but Fe III is not observed and Si II is generally weaker in SN Ic.

SN 2005gj has stronger ejecta-CSM interaction than SN 2002ic. The peak

bolometric luminosity is ∼ 1.5 times brighter and the broad and narrow components

of Hα are ∼ 4 times more luminous in SN 2005gj. The fact that the SN 1991T

features are weaker in SN 2005gj compared with SN 2002ic at similar epochs is

consistent with this interpretation, because the supernova features are more diluted

by the stronger continuum. The absence of evidence for a secondary maximum in

SN 2005gj, whereas the i band light curve of SN 2002ic does show a hint of such a

feature, is likewise consistent with the ejecta-CSM interaction in SN 2005gj having

been stronger than in SN 2002ic.

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6.6.1. Structure of the CSM

The circumstellar interaction of core-collapse supernovae with a circumstellar

medium has been studied in detail in the literature (see Chevalier & Fransson

(2003) for a review). When the fast moving ejecta encounters the approximately

stationary CSM, a forward shock moving into the CSM (also called circumstellar

shock) and a reverse shock develops. The fast-moving shockwave implies large

post-shock temperatures, therefore radiating energy in the X-ray regime. The

density distribution of the ejecta and the CSM can be well described by power-laws

in radius, which leads to a set of self-similar analytical solutions for the evolution

of the shock radius in time (Chevalier 1982). The physics of the ejecta-CSM

interaction in the case of thermonuclear supernovae is basically the same, the main

difference is in the distribution of the ejected material which follows an exponential

function in velocity (Dwarkadas & Chevalier 1998). In this case the solutions are

no longer analytic. The density profile of the shocked region is different in the case

of exponential ejecta expanding into a constant density medium, but the similarity

increases for expansion into a wind profile whose density decreases as ∝ r−2.

A simple self-similar model of a SN shock expanding into a medium with a

power-law density decline, as suggested for core-collapse SNe by Chevalier (1982),

is ruled out for this object by several observations: the exponential decrease in

luminosity, suggesting an exponential ejecta density profile; the strange behavior of

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the broad and narrow Hα components; and the decrease in the blackbody radius at

later times. Detailed calculations of the SN-CSM interaction would require highly

detailed hydrodynamic modeling, which are beyond the scope of this paper. Instead

herein we focus on trying to explain the basic features of SN-CSM interaction as

deduced from the observational data.

The initial velocity of a SN shock wave as it breaks out from the surface is at

least of the order of 2 × 104 km s−1. The broad Hα velocities that are seen in the

first week or so are of the order of 1500 km s−1, and increase to more than twice this

value after ∼ 50 days. These velocities are almost an order of magnitude smaller

than expected SN blast wave velocities in the early stages, and a factor of few

smaller even after ∼ 50 days. Furthermore, the SN shock velocity would be expected

to gradually decrease as the shock moves outwards, whereas the Hα profile actually

indicates an increasing velocity after ∼ 50 days.

For these reasons, we suggest that the broad Hα lines do not indicate the

SN velocity. Instead, we put forward a scenario of a shock expanding into a

two-component ambient medium: a low density wind in which are embedded

high-density clumps. In this picture, there should theoretically exist three different

velocity components: a broad velocity component, which is not easily seen in this

case, and is related to the velocity of the blast wave itself; an intermediate velocity

component, which is what we have referred to as the broad Hα and is related to

the velocity of the shock driven into the clumped material; and a narrow velocity

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component, which may be related to the narrow Hα and is representative of the

velocity of the ambient medium. This scenario is like the scenario put forward by

(Chugai & Danziger 1994) to explain the origin of the broad, intermediate and

narrow line components in SN 1988Z. The large Hα luminosity of SN 2005gj at late

times is very similar to that seen in other Type IIn SNe, and is especially large

considering that this was a Type Ia. However, there are significant differences. We

do not see a really broad line component representative of the SN velocity, although

there are some suggestions that this may be appearing at late times. In particular,

the Hα profile of the spectrum obtained at ∼ 150 days is better fitted by three

components, including a very broad component with FWHM ≈ 7000 kms−1.

Our scenario envisions the Type Ia SN shock wave expanding in a clumped

medium presumably formed by mass-loss from a companion star. The broad

component is not easily visible in Hα initially because the forward shock is not

radiative. The density of the clumps is much higher than that of the interclump

(ambient) medium. When the SN shock wave interacts with a dense cloud or clump,

it drives a strong shock into the clump. A reflected shock is driven back into the

expanding ejecta (Klein et al. 1994). Assuming pressure equilibrium, the ratio of the

velocity of the clump shock to that of the blast wave is inversely proportional to the

square-root of the ratio of the clump density to that of the interclump medium. The

optical emission arises from behind the clump shock, probably by reprocessing of the

X-ray emission.

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In this model, the intermediate component represents the velocity of the clump

shock, which is probably radiative. If we assume that the initial velocity of the SN

shock wave is ∼ 20000 km s−1 and the broad Hα emission velocity is ∼ 1500 km s−1,

then the ratio of velocities is 13− 14. This indicates that the clump density is about

142, or ∼ 200 times the interclump density. Note that the optical emission, which

goes as density squared, will then be 2002 times, or about 40,000 times greater

compared to that from the interclump medium. This is consistent with the fact that

no broad line emission is seen from the interclump medium. If the initial velocity is

much higher, as is conceivable, the clump density could be up to ∼ 50% higher, and

the ratio between the emission from the dense clumps and interclump medium even

larger.

What value of the clump density is suggested? A shock wave traveling at

1500 km s−1 would be radiative if it were expanding in a medium whose density

is greater than ∼ 106 cm−3, whereas a 2500 km s−1 shock would require minimum

densities of the order of 107 cm−3 (Draine & McKee 1993) in order to be radiative.

The CSM density, being two orders of magnitude smaller, would then to be

∼> 104 cm−3. These are just minimum values, and it is conceivable that the actual

clump density is much higher. This result is consistent with the conclusion of

Aldering et al. (2006).

The observations show that the broad Hα width increases after 50 days,

suggesting an increase in the clump shock velocity at later times, which could

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perhaps be due to a decrease in the clump density. Conversely, however, the

luminosity of the Hα also increases, suggesting an increase in the electron density.

At the same time, we would expect the SN shock to be decreasing in velocity as it

continues its outwards expansion.

We suggest that the way to reconcile these observations is a scenario in which

the density within the clump medium starts out higher than 108 cm−3, probably

as high as 1010 cm−3 in the first few days, and decreases gradually outwards. The

almost constant behavior of the FWHM of the broad Hα suggests that the density

profile of the ambient medium is flatter than r−2. Since we want the clump shock to

be radiative even when the shock velocity is almost 3000 km s−1, this suggest that

the density at ∼ 150 days is greater than about 107 cm−3. And since the density

is decreasing outwards, we infer that the density close in is even larger. Over the

entire period of observations the clump density is large enough that the shock driven

into these clumps is always radiative. The density of the ambient medium is two

orders of magnitude smaller, as discussed above. The high bolometric luminosity is

consistent with these values.

For the first ∼ 50 days the Hα emission arises only from the radiative shock

driven into the dense clumps. However, by ∼ 50 days the SN forward shock, which is

decreasing in velocity, enters the radiative regime, and the cooling shell of material

begins to contribute to the Hα luminosity. The velocity of the SN shock is quite

large, and its contribution initially is not a large fraction of the total Hα luminosity.

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But as it expands outwards, its velocity decreases and the shock becomes more

radiative, and the contribution to the total Hα luminosity increases, more than

compensating for the decreasing density. If this conclusion is correct, then we

would expect that a broad velocity component would be visible in the Hα spectra,

whose intensity would gradually increase with time even as the FWHM decreases.

Although the underlying supernova contamination makes it hard to isolate a broad

component, it is suggestive that by day ∼ 150 the spectrum is best fit by a third,

much broader component of the velocity, thus providing support for this line of

reasoning.

Finally, in this model the narrow line emission arises from the unshocked slowly

expanding ambient material, presumably the outflow that we find expanding at ∼60

km s−1. We note that although the width of the narrow line Hα emission as listed is

higher, it is still unresolved, and it is possible that within the limits of resolution the

narrow line component and outflow velocity are indeed the same.

To summarize, in this model the Type Ia SN expands in a clumped ambient

medium, with the clump density about ∼ 200 times that of the surrounding

medium close in to the star, and decreasing as we go outwards. The Hα emission

initially arises mainly from the shock driven into the dense clumps. The SN shock

propagating into the interclump medium begins to enter the radiative regime around

day 50, and its contribution to the Hα emission gradually increases beyond that

coming from the clumped medium, leading to the gradual rise in the Hα emission.

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We note that several features of this model are similar to the model presented by

Chugai et al. (2004) for SN 2002ic, thus further supporting the similarity between

the two supernovae.

The upper limit on the X-ray luminosity obtained at 74 days after the explosion

can put a constrain on the mass loss rate from the precursor or companion (e.g.,

Immler et al. 2006). Assuming that the X-ray luminosity is dominated by emission

from the reverse shock we obtain M ∼< 2 × 10−4 M⊙ yr−1 (2σ) using Equation 3.10

in Fransson et al. (1996). This value has to be taken as an approximate estimate

because we are making several assumptions about the physical properties of the

ejecta-CSM interaction that should be calculated using detailed hydrodynamical

simulations: a constant velocity of the shock, Vsh ≈ 8000 km s−1; a solar composition

of the CSM material; an electron temperature at the reverse shock of Te = 107 K,

which comes from the modeling of SN 2002ic (Nomoto et al. 2005); a flat density

profile of the CSM, ρ ∝ r−2; and a power-law ejecta density profile with index n = 7

(Nomoto et al. 1984b).

We can also estimate a mass loss rate from the companion using the density

of the ambient medium (n ∼ 108 cm−3), the initial optical radius of the CSM

(R ≈ 1015 cm), and the velocity of the wind: M = 4π R2 v ρ, this is assuming a

flat density profile for the CSM. We obtain: M ≈ 2 × 10−4 M⊙ yr−1, which is in

agreement with the 2σ upper limit calculated from the X-ray luminosity.

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The presence of Balmer lines in emission in the first spectrum obtained 6.6 days

after explosion shows that the ejecta started to interact with the CSM at an earlier

epoch (Aldering et al. 2006). Extrapolating linearly to zero flux the early increase

of Hα and Hβ fluxes we find that the ejecta-CSM interaction started 3 ± 1 days

after explosion, which gives an internal radius of the CSM Ri ≈ 1.1 × 1015 cm. The

outer radius of the CSM can be estimated assuming a constant velocity of the shock

of Vsh ≈ 8000 km s−1 over the first year. We detect Hα in emission in a spectrum

obtained at 368 days after explosion, which will be presented elsewhere, putting a

lower limit on the outer radius of the CSM, Ro ∼> 3× 1016 cm. This is also consistent

with a Type Ia SN with an exponential ejecta density profile expanding outwards in

a medium of average density ∼> 107 cm−3.

In the interpretation above we assume that the broad component of the

Hydrogen Balmer lines originate in the dense clumps, while the narrow component

arises from the photoionized un-shocked gas. However, Thompson scattering of the

lines has been considered as an alternative mechanism that can explain relatively

well the symmetric line profiles of SN 2002ic (Wang et al. 2004) and SN 2005gj

(Aldering et al. 2006). In this scenario, both components would arise from a single

high density region. The total mass of hydrogen in the emitting region would

be MH ≈ 2 × 10−2 (1010/ne) M⊙, where ne is the average electron density in the

emitting zone, as calculated from the luminosity of the Hα line at maximum using

the Case B recombination coefficient. The electron density must be sufficiently

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high, ne ≈ 1010 cm−3, to be consistent with the line ratios of He lines observed in

the spectra (Aldering et al. 2006), and a high electron density would explain the

non-detection in X-ray and radio (Soderberg & Frail 2005). However, it is unlikely

that this model would be able to explain the initial constancy and then rise of the

broad Hα luminosity.

6.6.2. Rates, hosts galaxies and possible progenitors of

SN 2002ic-like supernovae

The SDSS-II Supernova Survey has a well understood discovery efficiency of

SN Ia at low redshift (z ∼< 0.1), which allows us to obtain an accurate supernova

rate measurement controlling systematic errors (Dilday et al. 2008). In the fall

2005 season there were a total of 16 spectroscopically confirmed SN Ia (including

1991T-like and 1991bg-like objects) at z < 0.12, one photometric identification, and

the spectroscopically confirmed peculiar events: SN 2005hk (Phillips et al. 2007)

and SN 2005gj. Since the detection efficiency of 2002ic-like objects has not been

carefully modeled, we can only put a lower limit on the fraction of these events.

The spectroscopic confirmation of one object at z < 0.12 puts a lower limit of 5+7−4%

(68% confidence) in the fraction of 2002ic-like events among SN Ia at low redshift.

From the previously known (2002ic) and probable events (1997cy and 1999E) the

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estimated fraction is ∼ 1% of SN Ia discovered between 1997 and 2002, which is

consistent with our limit.

In the fall of 2006 we obtained the spectrum of a slowly declining supernova that

was discovered in 2005, but did not have a spectroscopic classification, SN 70176.

To our surprise, the late spectrum of SN 7017 resembles that of SN 2005gj one year

after explosion and the early photometry also shows similarities which lead us to

classify it as the highest redshift SN 2002ic-like object observed to date, at z = 0.27.

Considering SN 7017 in the spectroscopically confirmed sample of SN Ia during the

2005 season, a total of 129 SNe at z ∼< 0.42, we have that 2/130 (1.5%) are SN 2002ic

like objects, which is consistent with the low limit on the fraction at low redshift

estimated before. However, this fraction has to be taken with extreme care and

probably does not reflect the true fraction. This is because the discovery efficiency

of SN Ia declines as a function of redshift and the total number of spectroscopically

confirmed SN Ia does not include SNe with good Ia-like light curves that were not

spectroscopically classified. A more careful study of the rates of SN 2005gj-like

supernovae in the SDSS-II is planned for a future publication.

The host galaxies of supernovae can provide important clues about their

progenitors. The host of SN 2005gj is a very blue, low-luminosity dwarf

(MB ≈ −17), and has an irregular morphology with no well defined core. Aldering

6This is the internal name given by the SDSS-II Collaboration. It was not announced in an IAU

circular because of the late spectroscopic classification.

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et al. (2006) combined the SDSS photometry with UV imaging from GALEX to

construct an SED of the galaxy. They constrained the metallicity to Z < 0.3 Z⊙,

with a burst of star formation ∼ 200 Myr ago. SN 2002ic has a late type (Sbc) spiral

host with a well defined core. The host of SN 1997cy is also a blue, low-luminosity

(MV ≈ −18.2), and low surface brightness dwarf irregular galaxy (Germany et al.

2000). GALEX has imaged the positions of SN 1997cy, 2002ic and 2005gj, and

their hosts galaxies are all detected in the Near-UV (NUV) band. Their absolute

magnitudes in the NUV (AB magnitudes) are between -16.6 (SN 2005gj) and -17.3

(SN 1997cy and 2002ic). They are low-luminosity late type galaxies, ∼ 1 − 1.7

magnitudes fainter than L∗ galaxies observed by GALEX at redshift z < 0.1 (Wyder

et al. 2005). The host galaxy of SN 1999E is a late spiral with a nuclear starburst

first observed by the IRAS satellite (Allen et al. 1991). From 2MASS photometry, its

absolute magnitude is 1 mag brighter than an L∗ galaxy in the K-band (Kochanek

et al. 2001). SN 7017 at redshift z = 0.27, has a blue, dwarf-like host galaxy with

absolute magnitude in B of −17.9.

The host galaxies of the five SN 2002ic-like objects known share some common

properties: they are late type galaxies, irregulars and late spirals, most likely with

recent star formation. Four of the host galaxies have low luminosities, similar to the

Magellanic clouds, which indicates they are low metallicity systems. For example,

a dwarf galaxy with intrinsic luminosity MB = −18 has an Oxygen abundance

of 12 + log(O/H) ≈ 8.4 (van Zee et al. 2006), which corresponds to 1/3 the solar

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Oxygen abundance (Delahaye & Pinsonneault 2006). On the other hand, the host

galaxy of SN 1999E has a K-band luminosity, that when converted to metallicity

using the luminosity-metallicity relationship derived by Salzer et al. (2005), makes

it consistent with the solar value. The host luminosities are only an approximate

indicator of their metallicities, therefore spectra of the hosts are needed to infer

the metallicities and star formation rates (SFRs) of these galaxies. However, it is

interesting to note that the range of host galaxy properties of SN 2002ic-like events

seem to be inconsistent with the host galaxies of GRBs associated with supernovae

(Stanek et al. 2006) and broad-lined type Ib/c SNe (Modjaz et al. 2008).

Type Ia supernovae are observed in all types of galaxies. There is a well

established correlation between the morphology of their host galaxies and the peak

luminosity of the SNe: brighter supernovae (1991T-like) tend to explode in late

type spirals and irregulars with recent star formation, while intrinsically fainter

events (1991bg-like) are observed mainly in early type galaxies with an old stellar

population (Hamuy et al. 1995, 1996a; Branch et al. 1996; Hamuy et al. 2000;

Gallagher et al. 2005). This environmental effect and observations of the local

supernovae rate as a function of host galaxy properties (Cappellaro et al. 1999;

Mannucci et al. 2005), motivated Scannapieco & Bildsten (2005) to parametrize the

delay time distribution, time between star formation and the appearance of SNe,

and the rates with a two-component model having a piece proportional to the SFR

of the host galaxy (or prompt, they explode ∼ few × 108 yr after an episode of

215

star formation), and a second piece proportional to the total stellar mass (delayed

component, they explode on scales of a few Gyr after the onset of star formation).

The difference in age of the stellar populations of these subclasses suggests that the

progenitors may also be different: prompt SN Ia would come from more massive

progenitors. The host galaxies of all five SN 2002ic-like events known are broadly

consistent with the properties of the hosts of prompt SN Ia, which suggest a real

association given that the best studied SN IIa to date, SN 2002ic and SN 2005gj,

have spectral characteristics similar to 1991T-like events.

Several progenitors have been discussed in the literature for SN 2002ic and

SN 2005gj. Livio & Riess (2003) proposed that SN 2002ic is a rare case of a

double-degenerate binary system, a white dwarf (WD) and the core of an AGB star

spiraling-in through gravitational wave losses, in which the explosion occurs during

or immediately after the common-envelope phase (a few hundred to a few thousand

years of duration). The difference in line strengths of the Balmer emission lines

observed for SN 2002ic and SN 2005gj makes this scenario unlikely. Also, as Aldering

et al. (2006) points out, in both SN 2002ic and SN 2005gj the mass loss stopped

only a few years before explosion, which is too short compared with the timescale

for gravitational wave radiation to produce the merger of the core and the WD.

Another possible progenitor initially proposed by Hamuy et al. (2003) and

favored by the models of Chugai et al. (2004), is the explosion of the Chandrasekhar-

mass Carbon-Oxygen core of a massive AGB star in a degenerate medium, a

216

supernova Type 1.5 (Iben & Renzini 1983), where the dense Hydrogen-rich CSM

would come from the outer layers of the AGB. In order for the core to grow to the

Chandrasekhar mass, the radiatively driven winds from the AGB have to be weak

enough, a condition that is only met in a very low-metallicity environment like the

Galactic halo (Zijlstra 2004). At face value, the range of host galaxy metallicities

for SN 2002ic-like events inferred from the luminosity-metallicity relation does not

support the SN 1.5 scenario, although admittedly these are average metallicities and

do not tell us the actual range of metallicities of the progenitors.

Han & Podsiadlowski (2006) proposed that SN 2002ic could be produced

through the “super-soft channel”, the most common single-degenerate model for the

progenitors of SN Ia. In this scenario the white dwarf is accreting material from

a main sequence, or slightly evolved, relatively massive companion (∼ 3 M⊙) and

experiences a delayed dynamical instability that leads to a large amount of mass-loss

from the system in the last few × 104 yr before the explosion. Aldering et al. (2006)

notes that the estimated main-sequence mass of the progenitor of SN 2005gj of

∼ 2 M⊙, calculated using the age of the starburst of its host galaxy, is consistent with

the Han & Podsiadlowski (2006) model. Also, the predicted fraction of SN Ia that

would be produced through the “delayed dynamical” channel is 0.1-1%, consistent

with the limits we have obtain from the detection of SN 2005gj in the SDSS-II

Survey.

217

In general terms, the progenitor model proposed by Han & Podsiadlowski

(2006) successfully reproduces the observational properties of SN 2002ic and

SN 2005gj. However, it is still very early in the study of this new sub-class of SN Ia.

It would be interesting to see in the near future the results of theoretical modeling

exploring other single degenerate configurations (e.g., AGB donor) and detailed

hydrodynamical modeling of the ejecta-CSM interaction of SN 2005gj using the

observations of the early photometric and spectroscopic evolution presented in this

work.

218

SN

16

15

1413

12

11

10

9

8

7

6

5

4

3

2

1

N

E

Fig. 6.1.— r′-band image (3.5′×3.5′) of the field around SN 2005gj obtained with theSwope-1m telescope at LCO. North is up and east is to the left. Sixteen comparisonstars in common between SDSS and CSP used to derive differential photometry ofthe SN are labeled as in Table 6.1. The SN is close to the center of the field.

219

Fig. 6.2.— Spectra of SN 2005gj obtained from Oct. 1 (∼7 days after explosion)to Nov. 28 (∼61 days after explosion) of 2005. The sequence show the dramaticspectral evolution of the SN from a very blue continuum with strong Hydrogen-Balmer lines in emission in the early phases, resembling the spectrum of a Type IInSN, to a Type Ia supernova-dominated continuum with broad absorption and emissionfeatures (P-cygni profiles) of blended Fe II and Fe III profiles. The spectra areshown in logarithmic flux scale and a constant shift has been applied for clarity.The wavelength is in the rest-frame corrected using z = 0.0616 for the host galaxy.We show the UT date when the spectra were obtained and the epoch (rest-frame daysafter explosion) in parenthesis.

220

Fig. 6.3.— Late time spectra of SN 2005gj obtained between Dec. 8, 2005 (∼71 daysafter explosion) and Mar. 6, 2006 (∼152 days after explosion). The labels, axis andsymbols are the same as in Figure 6.2. The earth symbol shows the position of atelluric feature present in some of the spectra.

221

Fig. 6.4.— Observed light curves of SN 2005gj from SDSS (open circles), MDM(open squares) and CSP/Swope (filled triangles). The error bars are smaller than thesymbols. For clarity, the light curves have been shifted by an arbitrary constant.

222

Fig. 6.5.— Difference between the synthetic g − r color calculated from the spectraand the observed color from the photometry. We do not include the latest spectraobtained on Jan 24. and Mar. 6 because there is no contemporaneous photometricdata.

223

Fig. 6.6.— Time evolution of the colors of SN 2005gj (filled circles). For comparisonwe also show the color evolution of the overluminous Type Ia SN 1991T (solidline), the Type IIn SN 1999el (dashed line), and two previous cases of Type Iastrongly interacting with its circumstellar medium, SN 2002ic (stars) and SN 1997cy(triangles).

224

Fig. 6.7.— Absolute ugri light curves of SN 2005gj (filled circles). For comparisonwe also show the absolute light curves of the overluminous Type Ia SN 1991T (solidline), SN 2002ic (stars) and SN 1997cy (triangles). The error bar in the lower rightpannel represents the typical error in the absolute magnitudes dominated by a 10%uncertainty in the Hubble constant.

225

Fig. 6.8.— Absolute light curves of SN 2005gj in the NIR: J (top panel), H (middle

panel) and Ks (bottom panel). For comparison we also show the absolute light curvesof a normal Type Ia (solid line) and the Type IIn SN 1999el (open triangles). Theerror bar in the lower right of each panel represents the typical error in the absolutemagnitudes dominated by a 10% uncertainty in the Hubble constant.

226

Fig. 6.9.— Examples of black-body fits (solid line) to the SED of SN 2005gj obtainedby transforming the rest-frame ugriz magnitudes to monochromatic fluxes at theeffective wavelength of the filters (filled circles). These examples show the quality(i.e., goodness-of-fit) range of the black body fits at different epochs: χ2

ν =0.3(left panel), 2.4 (middle), 4.4 (right). The units of flux density in the y-axis aremJy = 10−26 erg s−1 cm−2 Hz−1.

227

Fig. 6.10.— Top panel: Quasi-bolometric (open circles) and bolometric light curvesof SN 2005gj (filled circles). The bolometric light luminosities were obtained afterapplying bolometric corrections calculated from black-body fits to the optical SEDobtained from the ugriz photometry. The dashed line shows the best-fit linear decayof 0.014 mag day−1. The middle and bottom panels show the evolution of the black-body temperature and radius, respectively.

228

Fig. 6.11.— Comparison of spectra of SN 2005gj at 26, 46, 62 and 83 days afterexplosion with comparable epoch spectra of SN 2002ic (Hamuy et al. 2003) andSN 1997cy (from SUSPECT database). The spectra are plotted on a logarithmicflux scale and shifted by an arbitrary constant. The wavelength was shifted to therestframe using z = 0.0616 of the host galaxy (Aldering et al. 2006).

229

Fig. 6.12.— Results of SNID. We show the spectra of SN 2005gj at four epochs(26, 37, 61, and 85 days after explosion) and their best three cross-correlation libraryspectra. The spectra are plotted on a logarithmic flux scale and shifted by an arbitraryconstant.

230

Fig. 6.13.— Fits to the spectra of SN 2005gj. We model the spectra as the sum of twocomponents: (1) SN 1991T spectrum at the same epoch after explosion as SN 2005gjscaled by an arbitrary constant (blue line); (2) fourth order polynomial (green line).The results of the fits are in red and the spectra of SN 2005gj, corrected by Galacticextinction in the line of sight are in black. The epochs of the spectra are shown inthe upper right of each panel.

231

Fig. 6.14.— Results from the Gaussian fits to Hα and Hβ emission features as afunction of time. Top left panel: FWHM of the Hα-broad and Hβ. Top right panel:

Luminosity of the narrow and broad Gaussian components of Hα. Bottom left panel:

Luminosity of Hβ. Bottom right panel: Balmer decrement, ratio of total fluxes in Hαand Hβ lines.

232

Fig. 6.15.— Line profiles of Hβ (top) and Hα (bottom) in the two highest resolutionspectra of SN 2005gj obtained with WHT+ISIS on day 44 (left) and Magellan+LDSS-3 on day 115 (right). The features show clear P-cygni profiles with weak absorptionminima at ∼ −200 kms−1, demonstrating the presence of a slowly moving outflow.

233

Fig. 6.16.— Identification of lines in the spectrum of SN 2005gj obtained at 17 daysafter explosion (2 days before g maximum). The red line shows the best fit syntheticspectrum generated with the SYNOW code. The lines of SN 2005gj are typicalof SN 1991T around maximum light (blue line), and very similar to the spectrumof SN 2002ic around maximum (green line). All the spectra have been locallynormalized. We have subtracted a constant value to the spectra of SN 2002ic andSN 1991T for clarity.

234

Star u g r i z

ID SDSS CSP SDSS CSP SDSS CSP SDSS CSP SDSS

1 18.639(027) 18.690(054) 17.367(018) 17.362(009) 16.837(022) 16.823(013) 16.643(018) 16.586(010) 16.549(020)

2 · · · · · · 19.122(027) 19.043(018) 17.609(014) 17.575(015) 16.191(017) 16.051(017) 15.403(018)

3 19.843(039) 19.899(097) 17.676(016) 17.653(023) 16.745(012) 16.726(015) 16.404(015) 16.362(010) 16.215(017)

4 · · · · · · 20.116(026) 20.038(045) 19.579(018) 19.556(062) 19.356(025) 19.113(067) 19.207(040)

5 · · · 20.262(436) 18.856(021) 18.770(015) 17.345(009) 17.323(009) 16.595(014) 16.505(010) 16.145(064)

6 16.074(030) 16.102(011) 14.556(010) 14.544(011) 13.965(016) 13.961(015) 13.774(007) 13.728(015) 13.711(023)

7 19.915(041) 20.347(285) 18.385(019) 18.384(017) 17.741(020) 17.735(010) 17.502(016) 17.451(013) 17.372(021)

8 18.486(027) 18.533(030) 17.214(011) 17.204(009) 16.694(017) 16.685(009) 16.508(018) 16.473(009) 16.418(016)

(cont’d)

Table 6.1. SDSS ugriz and CSP u′g′r′i′ photometry of comparison stars in common in the field of SN 2005gj.

235


Star u g r i z

ID SDSS CSP SDSS CSP SDSS CSP SDSS CSP SDSS

9 18.847(021) 18.923(041) 17.458(013) 17.453(009) 16.940(011) 16.936(011) 16.765(018) 16.732(011) 16.679(024)

10 18.285(025) 18.306(026) 16.956(018) 16.935(011) 16.412(011) 16.404(009) 16.208(017) 16.181(009) 16.125(021)

11 · · · · · · 20.227(033) 20.118(058) 18.717(014) 18.715(017) 17.861(021) 17.787(014) 17.347(020)

12 · · · · · · 18.927(020) 18.860(015) 17.433(014) 17.404(020) 16.669(018) 16.611(010) 16.243(024)

13 18.510(035) 18.529(031) 16.922(017) 16.897(009) 16.249(016) 16.226(011) 15.992(017) 15.958(009) 15.868(016)

14 20.030(065) 20.081(196) 17.833(019) 17.819(010) 16.882(007) 16.864(013) 16.515(017) 16.470(009) 16.306(015)

15 18.748(029) 18.314(545) 16.915(010) 16.888(009) 16.142(013) 16.117(009) 15.845(022) 15.798(009) 15.708(015)

16 18.767(032) 18.834(109) 17.458(013) 17.454(011) 16.896(015) 16.882(009) 16.692(013) 16.652(009) 16.599(018)

Note. — Uncertainties given in parentheses are in thousandths of a magnitude. For the CSP photometry with the Swope they

correspond to the rms of the magnitudes obtained on four photometric nights, with a minimum uncertainty of 0.015 mag for an

individual measurement.

236

JD Epocha

−2, 453, 000 (days) u g r i z Source

616.94 . . . 24.00±1.24 25.19±0.80 23.81±0.61 23.36±0.54 22.15±0.53 SDSS

626.91 . . . 21.12±0.70 27.83±0.86 22.85±0.45 24.92±0.93 22.28±1.77 SDSS

628.90 . . . 23.88±1.45 25.72±1.23 22.92±0.41 25.98±0.45 21.72±0.54 SDSS

635.91 . . . 22.00±0.49 26.38±1.11 23.32±0.62 23.27±0.59 22.15±0.53 SDSS

639.95 1.9 18.887(045) 18.592(025) 18.621(010) 18.718(016) 18.924(043) SDSS

641.95 3.8 18.468(034) 18.154(011) 18.256(020) 18.340(016) 18.511(041) SDSS

644.89 6.6 18.121(023) 17.795(011) 17.879(022) 17.935(010) 18.141(022) SDSS

656.94 17.9 . . . 17.355(026) 17.255(029) 17.360(064) . . . MDM

657.90 18.8 17.949(039) 17.343(012) 17.214(012) 17.275(016) 17.438(024) SDSS

663.89 24.5 18.149(039) 17.410(050) 17.158(012) 17.163(012) 17.280(027) SDSS

663.92 24.5 . . . 17.368(025) 17.160(020) 17.248(055) . . . MDM

665.91 26.4 . . . . . . 17.146(028) 17.215(054) . . . MDM

666.96 27.3 18.282(043) 17.408(012) 17.121(012) 17.133(013) 17.268(015) SDSS

668.87 29.2 . . . 17.393(014) 17.125(026) 17.186(045) . . . MDM

668.88 29.2 18.355(028) 17.401(009) 17.127(012) 17.096(016) 17.240(016) SDSS

669.97 30.2 18.347(045) 17.430(011) 17.101(014) 17.105(031) 17.231(017) SDSS

670.88 31.0 18.386(024) 17.419(007) 17.125(013) 17.092(014) 17.249(021) SDSS

673.85 33.8 18.479(028) 17.450(009) 17.123(012) 17.081(014) 17.229(015) SDSS

(cont’d)

Table 6.2. SDSS and MDM ugriz photometry of SN 2005gj

237


JD Epocha

−2, 453, 000 (days) u g r i z Source

675.85 35.7 18.551(025) 17.477(010) 17.134(029) 17.088(019) 17.182(019) SDSS

676.90 36.7 . . . 17.464(043) 17.116(026) 17.138(055) . . . MDM

680.86 40.4 18.663(027) 17.578(009) 17.133(011) 17.054(009) 17.156(016) SDSS

683.92 43.3 18.759(062) 17.656(034) 17.220(055) 17.080(015) 17.061(066) SDSS

686.86 46.1 18.887(047) 17.695(014) 17.160(011) 17.060(010) 17.138(021) SDSS

687.91 47.1 18.937(105) 17.687(013) 17.206(024) 17.042(018) 17.119(026) SDSS

693.86 52.7 19.018(081) 17.780(009) 17.207(015) 17.084(014) 17.130(018) SDSS

699.85 58.3 . . . 17.864(060) 17.271(034) 17.202(057) 17.166(024) MDM

699.88 58.4 19.083(042) 17.909(013) 17.273(009) 17.124(010) 17.135(020) SDSS

727.70 84.6 . . . 18.368(025) 17.639(043) 17.551(053) 17.407(015) MDM

737.75 94.0 . . . 18.525(038) 17.791(033) 17.709(063) 17.470(057) MDM

739.64 95.8 . . . 18.540(037) 17.822(035) 17.744(065) 17.548(027) MDM

Note. — Uncertainties given in parentheses in thousandths of a magnitude.

aRest-frame days since the time of explosion (JD 2,453,637.93).

238

JD Epoch

−2, 453, 000 (days) u′ g′ r′ i′

698.69 57.2 19.159(042) 17.884(017) 17.243(017) 17.106(017)

699.67 58.2 19.196(041) 17.910(017) 17.261(017) 17.100(017)

702.68 61.0 19.286(042) 17.959(017) 17.281(017) 17.130(017)

706.69 64.8 19.312(040) 18.031(017) 17.345(017) 17.173(017)

712.64 70.4 19.619(060) 18.157(017) 17.418(017) 17.249(017)

720.66 77.9 19.721(131) 18.283(023) 17.542(017) 17.363(017)

725.65 82.6 19.964(079) 18.365(017) 17.602(017) 17.412(017)

728.71 85.5 19.938(104) 18.430(017) 17.658(017) 17.453(017)

736.62 93.0 20.104(084) 18.554(017) 17.794(017) 17.586(017)

740.64 96.8 20.229(090) 18.606(017) 17.856(017) 17.641(017)

741.59 97.6 20.238(103) 18.635(017) 17.872(017) 17.651(017)

746.61 102.4 20.532(220) 18.641(025) 17.926(017) 17.725(017)

754.58 109.9 20.509(143) 18.770(017) 18.021(017) 17.803(017)

761.64 116.5 · · · 18.847(017) 18.100(017) 17.910(018)

763.56 118.3 · · · 18.841(017) 18.107(017) 17.960(017)

764.58 119.3 · · · 18.859(017) 18.143(017) 17.947(017)

768.60 123.1 · · · 18.882(019) 18.200(017) 17.978(021)

773.55 127.7 · · · 18.884(026) 18.239(019) 18.050(021)

774.56 128.7 · · · 18.946(031) 18.221(017) 18.079(022)

783.55 137.2 · · · 19.017(018) 18.322(017) 18.117(021)

786.53 140.0 · · · 19.015(017) 18.363(017) 18.178(020)

795.55 148.5 · · · 19.099(019) 18.416(021) 18.227(031)


Table 6.3. CSP u′g′r′i′ photometry of SN 2005gj

239

JD Epoch

−2, 453, 000 (days) Y J H Ks Instrument

700.71 59.1 16.565(015) 16.484(034) 16.253(030) · · · Retrocam

704.68 62.9 16.591(015) 16.537(020) 16.315(033) · · · Retrocam

709.66 67.6 16.628(015) 16.550(023) 16.271(029) · · · Retrocam

714.58 72.2 16.673(015) 16.594(020) 16.389(028) · · · Retrocam

718.65 76.0 16.725(016) 16.658(022) 16.364(037) · · · Retrocam

722.62 79.8 16.832(016) 16.725(032) 16.490(017) 16.384(096) WIRC

724.68 81.7 16.781(016) 16.716(028) · · · · · · Retrocam

727.69 84.6 16.872(024) 16.757(036) · · · · · · Retrocam

732.71 89.3 16.920(019) 16.839(025) · · · · · · Retrocam

750.61 106.1 17.307(016) 17.152(016) 16.812(016) · · · WIRC

755.57 110.8 17.337(024) 17.263(024) 16.891(025) 16.745(039) PANIC

756.59 111.8 17.428(016) 17.269(016) 16.928(017) · · · WIRC

757.60 112.7 17.423(024) 17.270(024) 16.938(025) 16.807(042) PANIC

773.57 127.8 17.677(016) 17.476(016) · · · · · · WIRC

776.54 130.6 17.538(026) 17.399(070) 17.081(068) · · · Retrocam

777.55 131.5 17.648(038) 17.417(056) 17.129(119) · · · Retrocam

782.55 136.2 17.513(039) 17.436(073) 17.072(114) · · · Retrocam

783.55 137.2 17.766(024) 17.586(024) 17.271(025) 17.045(042) PANIC

785.53 139.0 17.765(025) 17.570(022) 17.192(049) 17.152(177) WIRC

788.54 141.9 17.869(016) 17.594(016) 17.335(019) · · · WIRC

797.52 150.3 17.855(061) 17.563(118) 17.313(158) · · · Retrocam

800.50 153.1 17.840(041) · · · · · · · · · Retrocam

808.52 160.7 17.871(058) 17.551(079) · · · · · · Retrocam

814.49 166.3 17.864(084) · · · · · · · · · Retrocam


Table 6.4. CSP Y JHKs photometry of SN 2005gj

240

Parameter Value

Time of explosiona 637.93 ± 2.02

Time of umax 651.77 ± 0.48

Time of gmax 657.80 ± 1.28

Time of rmax 672.49 ± 1.39

Time of imax 684.51 ± 1.00

umax 17.85 ± 0.05

gmax 17.35 ± 0.01

rmax 17.12 ± 0.01

imax 17.05 ± 0.01

u0max

b 17.11 ± 0.12

g0max 16.94 ± 0.09

r0max 16.83 ± 0.07

i0max 16.89 ± 0.11

M0g,max –20.21

E(B − V )Galc 0.121 ± 0.019

Ag(Gal) 0.45 ± 0.07

aJD-2, 453, 000

bMagnitudes at maximum in the

rest-frame, they have been corrected

by Galactic

extinction and K-corrections. We

assume a negligible extinction in the

host galaxy.

cFrom Schlegel et al. (1998)

Table 6.5. Light-curve parameters for SN 2005gj

241

JD Epoch Wavelength Resolutiona Exposure

−2, 453, 000 (days) Instrument Range (A) (A) (s)

644.92 6.6 MDM-CCDS 3850 – 7308 15 1200

646.95 8.5 ARC-DIS 3824 – 10192 7 1800

650.84 12.2 ARC-DIS 3600 – 9597 7 1000

655.87 16.9 MDM-CCDS 3823 – 7283 15 1800

665.92 26.4 MDM-CCDS 3883 – 7341 15 2700

668.83 29.1 MDM-CCDS 3886 – 7346 15 2700

676.79 36.6 MDM-CCDS 3882 – 7338 15 3600

684.73 44.1 WHT-ISIS 3924 – 8901 4 1800

686.79 46.0 MDM-CCDS 3858 – 7315 15 2700

698.67 57.2 duPont-ModSpec 3780 – 7290 7 2700

699.67 58.2 duPont-ModSpec 3780 – 7290 7 2700

700.76 59.2 MDM-CCDS 3933 – 7391 15 2700

702.73 61.1 MDM-CCDS 3856 – 7310 15 2700

712.73 70.5 MDM-CCDS 3831 – 7286 15 2700

722.71 79.9 NTT-EMMI 4000 – 10200 9 2700

(cont’d)

Table 6.6. Spectroscopic observations of SN 2005gj

242


JD Epoch Wavelength Resolutiona Exposure

−2, 453, 000 (days) Instrument Range (A) (A) (s)

724.66 81.7 duPont-WFCCD 3800 – 9235 6 2700

725.65 82.6 duPont-WFCCD 3800 – 9235 6 2700

726.66 83.6 duPont-WFCCD 3800 – 9235 6 3600

727.67 84.5 duPont-WFCCD 3800 – 9235 6 3600

728.67 85.5 duPont-WFCCD 3800 – 9235 6 3600

729.67 86.4 MDM-CCDS 3915 – 7373 15 2700

737.70 94.0 MDM-CCDS 3909 – 7368 15 2700

751.60 107.1 NTT-EMMI 3200 – 10200 9 2700

755.62 110.9 MDM-CCDS 3844 – 7299 15 3600

759.61 114.6 Magellan-LDSS-3 3788 – 9980 3 3600

799.52 152.2 duPont-WFCCD 3800 – 9235 6 1200

Note. — Most of the spectra are the combination of multiple observation, the total

exposure is given.

aAverage resolution obtained from the FWHM of arc-lamp lines.

243

Epoch

(days) Ku Kg Kr Ki Kz

6.6 0.039 –0.077 –0.053 –0.141 –0.133

12.2 0.162 –0.097 –0.005 –0.196 –0.047

16.9 0.226 –0.052 –0.055 –0.135 –0.223

26.4 0.276 –0.015 –0.058 –0.116 –0.228

29.1 0.255 0.008 –0.056 –0.104 –0.239

36.6 0.296 0.034 –0.030 –0.118 –0.096

46.0 0.251 0.063 –0.025 –0.079 –0.201

57.2 0.195 0.106 0.049 –0.123 –0.152

58.2 0.186 0.101 0.013 –0.068 –0.216

59.2 0.313 0.090 0.006 –0.072 –0.215

61.1 0.241 0.091 0.011 –0.096 –0.136

70.5 0.505 0.086 0.055 –0.113 –0.054

79.9 0.439 0.112 0.046 –0.090 –0.115

81.7 0.328 0.118 0.063 –0.096 –0.057

82.6 0.340 0.115 0.063 –0.104 –0.102

83.6 0.378 0.105 0.079 –0.099 –0.101

84.5 0.326 0.104 0.071 –0.104 –0.026

85.5 0.391 0.112 0.078 –0.110 –0.036

86.4 0.330 0.124 0.066 –0.094 –0.071

94.0 0.249 0.126 0.094 –0.081 –0.092

107.1 0.266 0.141 0.082 –0.084 –0.089

110.9 0.337 0.138 0.097 –0.091 –0.007

Table 6.7. K-corrections of SN 2005gj

244

Epoch log L(u→i) log L(u→z) log Lbol Tbba Rbb

b

(days) (erg s−1) (erg s−1) (erg s−1) (K) (1015 cm) χ2ν

c

1.9 42.822 42.862 43.285(0.109) 12898(1193) 1.02(0.16) 0.7

3.8 42.990 43.029 43.485(0.105) 13552(1316) 1.16(0.18) 0.3

6.6 43.141 43.180 43.637(0.098) 13569(1288) 1.38(0.21) 0.2

18.8 43.344 43.389 43.713(0.096) 10955(831) 2.32(0.34) 1.2

24.5 43.335 43.387 43.644(0.125) 9468(612) 2.87(0.42) 1.2

27.3 43.335 43.388 43.628(0.135) 9069(720) 3.06(0.50) 2.2

29.2 43.335 43.389 43.617(0.140) 8772(683) 3.22(0.53) 2.3

30.2 43.334 43.388 43.617(0.137) 8761(692) 3.23(0.54) 2.3

31.0 43.332 43.386 43.611(0.122) 8697(705) 3.25(0.55) 2.6

33.8 43.325 43.381 43.595(0.140) 8392(685) 3.43(0.59) 2.9

35.7 43.316 43.374 43.581(0.160) 8100(582) 3.61(0.59) 2.4

40.4 43.299 43.361 43.559(0.152) 7684(526) 3.92(0.64) 2.5

43.3 43.270 43.338 43.532(0.143) 7284(335) 4.24(0.59) 0.6

46.1 43.269 43.336 43.526(0.101) 7142(477) 4.36(0.73) 2.8

47.1 43.263 43.332 43.523(0.139) 7115(450) 4.38(0.70) 2.1

52.7 43.247 43.317 43.506(0.123) 6889(425) 4.58(0.74) 2.4

57.2 43.219 43.294 43.480(0.095) 6958(247) 4.32(0.54) 2.0

58.2 43.212 43.288 43.473(0.096) 6913(248) 4.34(0.55) 2.0

58.4 43.215 43.290 43.480(0.133) 6656(322) 4.78(0.70) 1.6

61.0 43.198 43.274 43.458(0.107) 6840(261) 4.34(0.56) 2.4

64.8 43.176 43.254 43.438(0.087) 6830(249) 4.27(0.54) 2.2

70.4 43.133 43.214 43.394(0.109) 6570(278) 4.33(0.58) 3.4

77.9 43.088 43.171 43.355(0.098) 6633(280) 4.10(0.54) 2.8

82.6 43.057 43.143 43.321(0.127) 6380(293) 4.18(0.58) 4.2

(cont’d)

Table 6.8. Derived integrated luminosity and black-body fits.

245


Epoch log L(u→i) log L(u→z) log Lbol Tbba Rbb

b

(days) (erg s−1) (erg s−1) (erg s−1) (K) (1015 cm) χ2ν

c

85.5 43.038 43.125 43.305(0.096) 6459(282) 4.04(0.54) 3.5

93.0 42.987 43.076 43.256(0.103) 6360(272) 3.92(0.53) 3.6

96.8 42.963 43.054 43.232(0.101) 6305(279) 3.87(0.53) 3.9

97.6 42.955 43.047 43.226(0.105) 6318(279) 3.83(0.52) 3.8

102.4 42.935 43.026 43.205(0.133) 6372(332) 3.68(0.52) 4.4

109.9 42.896 42.989 43.165(0.117) 6339(313) 3.53(0.50) 4.4

116.5 42.859 42.953 43.135(0.123) 6549(350) 3.25(0.46) 3.8

118.3 42.855 42.947 43.128(0.138) 6566(362) 3.21(0.46) 4.0

119.3 42.847 42.940 43.122(0.117) 6548(357) 3.20(0.46) 4.0

120.2 42.839 42.932 43.114(0.120) 6536(352) 3.18(0.45) 3.9

123.1 42.831 42.924 43.104(0.127) 6533(370) 3.14(0.45) 4.3

127.7 42.818 42.909 43.087(0.155) 6523(399) 3.06(0.46) 5.0

128.7 42.809 42.901 43.078(0.153) 6486(399) 3.07(0.46) 5.2

137.2 42.776 42.869 43.043(0.169) 6404(418) 2.99(0.46) 6.3

140.0 42.766 42.857 43.030(0.197) 6377(421) 2.96(0.46) 6.5

148.5 42.738 42.830 42.999(0.205) 6246(438) 2.94(0.48) 8.0

aBlack-body temperature from the fits to the broadband photometry; 1σ

uncertainty are given in parenthesis.

bBlack-body radius from the fits to the broadband photometry; 1σ uncertainty

are given in parenthesis and include a 10% uncertainty in the distance to

SN 2005gj.

cχ2 per degree of freedom of the black-body fits.

246

SN Name Class Epochs Reference

1990N Ia normal –14, –13, –8, –7, –6, 0, 4, 8, 15, 18, 39 1

1991T Ia 91T –9, –8, –7, –6, –5, –4, –2, –1, 9, 10, 11, 12,

15, 16, 17, 18, 19, 20, 21, 22, 23, 27, 43, 44,

47, 48, 51, 69, 77 2, 3

1991bg Ia 91bg 1, 3, 16, 18, 25, 32, 33, 46, 54, 85 4

1992A Ia normal –5, –1, 3, 5, 6, 7, 9, 11, 16, 17, 24, 28 5

1994I Ic normal –6, –4, –3, 0, 1, 2, 21, 22, 23, 24, 26, 30, 36, 38 6

1997ef Ic broad –14, –12, –11, –10, –9, –6, –5, –4, 7, 13, 14, 16, 17,

19, 22, 24, 27, 41, 45, 47, 49, 75, 80, 81 7

1998aq Ia normal –9, –8, –3, 0, 1, 2, 3, 4, 5, 6, 7,

19, 21, 24, 31, 32, 36, 51, 55, 58, 60, 63,

66, 79, 82, 91, 211, 231, 241 8

1998bu Ia normal –3, –2, –1, 9, 10, 11, 12, 13, 14, 28, 29, 30, 31, 32,

33, 34, 35, 36, 37, 38, 39, 40, 41, 42, 43, 44, 57 9

1998bw Ic broad –9, –7, –6, –3, –2, –1, 1, 3, 4, 6, 9, 11, 12, 13, 19,

22, 29, 45, 64, 125, 200, 337, 376 10

1999aa Ia 91T –11, –7, –3, –1, 5, 6, 14, 19, 25, 28, 33, 49, 47, 51 11

(cont’d)

Table 6.9. Library of spectra used in SNID

247


SN Name Class Epochs Reference

1999ee Ia normal –9, –7, –2, 0, 3, 8, 10, 12, 17, 20, 23, 28, 33, 42 12

1999ex Ic normal –1, 4, 13 12

1999by Ia 91bg –4, –3, –2, –1, 2, 3, 5, 6, 7, 8, 10, 11, 25, 29, 31, 33, 42 13

2002ap Ic broad –5, –4, 3, 8, 10, 17, 19 14

2004aw Ic normal 1, 5, 6, 8, 15, 21, 22, 26, 28, 39, 35, 44, 49, 63, 64, 236, 260, 413 15

2006aj Ic broad –6, –5, –4, –3, –2, –1, 0, 2, 3 16

References. — (1) Leibundgut et al. (1991); (2) Jeffery et al. (1992); (3) Schmidt et al. (1994); (4)

Leibundgut et al. (1993); (5) Kirshner et al. (1993); (6) Millard et al. (1999); (7) Iwamoto et al. (2000);

(8) Branch et al. (2003); (9) Jha et al. (1999); (10) Patat et al. (2001); (11) Garavini et al. (2004); (12)

Hamuy et al. (2002); (13) Garnavich et al. (2004); (14) Gal-Yam et al. (2002); (15) Taubenberger et al.

(2006); (16) Modjaz et al. (2006).

248

Hα (narrow) Hα (broad) Hβ

JD Epoch FWHMa fluxb FWHMa fluxb FWHMa fluxb

−2, 453, 000 (days)

644.92 6.6 . . . 0.24(0.02) 1575 0.58(0.06) 776 0.50(0.05)

646.95 8.5 137 0.37(0.04) 1481 0.83(0.08) 1307 0.52(0.06)

650.84 12.2 314 0.69(0.07) 1731 1.25(0.13) 1462 0.75(0.10)

655.87 16.9 . . . 0.57(0.06) 1555 1.11(0.11) 1339 0.75(0.09)

665.92 26.4 . . . 0.53(0.05) 1569 1.03(0.10) 523 0.48(0.06)

668.83 29.1 . . . 0.37(0.04) 1234 0.99(0.10) 1275 0.55(0.07)

676.79 36.6 . . . 0.45(0.05) 1513 0.77(0.08) 884 0.28(0.04)

686.79 46.0 . . . 0.41(0.04) 1836 0.71(0.07) . . . 0.15(0.02)

698.67 57.2 . . . 0.45(0.04) 2115 0.91(0.10) . . . 0.15(0.03)

699.67 58.2 . . . 0.34(0.04) 2053 0.76(0.08) . . . 0.14(0.03)

700.76 59.2 . . . 0.36(0.04) 1830 0.77(0.08) 620 0.16(0.02)

(cont’d)

Table 6.10. Results of the Gaussian fits to Hα and Hβ features

249


Hα (narrow) Hα (broad) Hβ

JD Epoch FWHMa fluxb FWHMa fluxb FWHMa fluxb

−2, 453, 000 (days)

702.73 61.0 . . . 0.41(0.04) 1978 0.67(0.07) . . . 0.11(0.02)

712.73 70.5 . . . 0.41(0.04) 2357 1.02(0.11) . . . 0.11(0.02)

722.71 79.9 . . . 0.34(0.04) 2413 1.03(0.11) 490 0.12(0.03)

724.66 81.7 . . . 0.34(0.04) 2137 0.99(0.10) 1067 0.19(0.03)

725.65 82.6 . . . 0.34(0.03) 2260 1.02(0.11) 568 0.11(0.02)

726.66 83.6 . . . 0.35(0.04) 2322 1.10(0.11) . . . 0.08(0.02)

727.67 84.5 160 0.32(0.03) 2364 1.02(0.10) . . . 0.10(0.02)

728.67 85.5 . . . 0.20(0.02) 1802 1.24(0.14) 1127 0.19(0.03)

729.67 86.4 . . . 0.37(0.04) 2687 1.15(0.12) 680 0.14(0.02)

737.70 94.0 . . . 0.29(0.03) 1941 0.85(0.09) 459 0.09(0.01)

755.62 110.9 . . . 0.32(0.03) 2236 1.14(0.12) 1031 0.16(0.02)

799.52 152.2 525 0.51(0.05) 3809 2.18(0.22) 1669 0.23(0.03)

aUnits of FWHM are in km s−1; FWHM is not presented when the spectral resolution is bigger

than the measured value.bUnits of flux are in 10−14 erg s−1 cm−2; 1σ uncertainties are given in parentheses.

250

Chapter 7

LBT Discovery of a Yellow Supergiant

Eclipsing Binary in the Dwarf Galaxy

Holmberg IX

7.1. Introduction

Although small in number, massive stars are critical to the formation and

evolution of galaxies. They shape the ISM of galaxies through their strong winds

and high UV fluxes, and are a major source of the heavy elements enriching the ISM

(e.g, Massey 2003; Zinnecker & Yorke 2007, and references therein). A large fraction

of massive stars are found in binaries (e.g., Kiminki et al. 2007). Eclipsing binaries

are of particular use because they allow us to determine the masses and radii of the

components and the distance to the system. Many young, massive eclipsing binaries

have been found and studied in our Galaxy, the LMC, and the SMC, primarily in OB

associations and young star clusters (e.g., Bonanos et al. 2004; Peeples et al. 2007;

Gonzalez et al. 2005; Hilditch et al. 2005). The study of massive eclipsing binaries

beyond the Magellanic clouds has been limited until very recently, when variability

searches using medium-sized telescopes with wide-field CCD cameras, coupled with

251

spectroscopy using 8-meter class telescopes, have yielded the first systems with

accurately measured masses in M31 (Ribas et al. 2005) and M33 (Bonanos et al.

2006).

We conducted a deep variability survey of M81 and its dwarf irregular

companion, Holmberg IX, using the Large Binocular Camera (LBC) mounted on the

Large Binocular Telescope (LBT), between January and October 2007. Holmberg IX

is a young dwarf galaxy (age ∼< 200 Myr), with a stellar population dominated by

blue and red supergiants with no signs of old stars in the red giant branch (Makarova

et al. 2002). The dwarf may have formed during a recent tidal interaction between

M81 and NGC 2976 (e.g., Boyce et al. 2001). The gas-phase metal abundance of

Holmberg IX of between 1/8 and 1/3 solar (e.g., Miller 1995; Makarova et al. 2002)

is consistent with this hypothesis (e.g., Weilbacher et al. 2003). A normal, isolated

dwarf on the luminosity-metallicity relationship would have a metallicity of ∼ 1/20

solar (Lee et al. 2006).

In this paper, we report on the discovery of a 271 day period, evolved, massive

eclipsing binary in Holmberg IX using data from the LBT. The overcontact system

is the brightest periodic variable discovered in our LBT variability survey. It has

an out-of-eclipse magnitude of Vmax = 20.7 mag and is located at α = 09h57m37.s14,

δ = +6902′11′′ (J2000.0). In §7.2 we discuss the observations and data reduction,

in §7.3 we present the light curve and the best-fit eclipsing binary model, and note a

similar eclipsing binary in the SMC. In §7.4 we discuss the results and their possible

252

implications for Type II supernovae. Throughout this paper, we assume the HST

Key Project distance to M81 of µ = 27.80 mag (3.6 Mpc; Freedman et al. 2001) as

the distance to Holmberg IX, and correct only for a foreground Galactic extinction

of E(B − V ) = 0.08 mag (Schlegel et al. 1998).

7.2. Observations

Holmberg IX was observed as part of a variability survey of the entire M81

galaxy conducted between January and October 2007 with the LBT 8.4-meter

telescope (Hill et al. 2006), using the LBC-Blue CCD camera (Ragazzoni et

al. 2006; Giallongo et al. 2008) during Science Demonstration Time. The survey

cadence and depth (1 min single exposures, with ≥ 3 consecutive exposures) are

optimized to detect and follow-up Cepheid variables with periods between 10–100

days (V ∼< 24 mag), getting better than 10% photometry in the B and V filters. We

obtained 168 V -band images on 24 different nights, and 87 B-band images on 13

nights. We coadded the B-band images from each night (usually 3–4) to improve

the signal-to-noise in the combined images. The seeing (FWHM) varied between

0.′′7 − 2.′′0 in V (median 1.′′4), and between 1.′′0 − 3.′′3 in B-band (median 1.′′9). Our

program did not request especially good image quality for these queue scheduled

SDT observations.

253

We also observed Holmberg IX as part of a variability survey of M81

conducted with the 8K Mosaic imager mounted on the MDM 2.4-meter telescope.

The observations were obtained in 5 one-week runs between February 2006 and

February 2007. All the images were obtained in V -band using 15 min exposures.

Due to weather loses and bad seeing, we ended up using only 36 images from 12

different nights. The typical seeing was ∼ 1.′′1.

7.3. Light Curve

We used the ISIS difference image analysis package (Alard 2000; Hartman et

al. 2004) to obtain the V -band light curves of all the point sources detected in the

LBT reference image. The detection of all point sources and the transformation

of difference-flux light curves to instrumental magnitudes were done using the

DAOPHOT/ALLSTAR package (Stetson 1987, 1992). After visual inspection of

all the light curves of variable point-sources selected by standard criteria (rms and

AoV significance; Hartman et al. 2007), we detected ∼ 20 periodic variables in the

field of Holmberg IX. These include Cepheids with periods of 10 − 60 days and one

long-period variable. The analysis of the Cepheid PL-relation and the distance to

Holmberg IX will be presented in a future paper. The brightest periodic variable is

the peculiar, long-period (P = 270.7 days) eclipsing binary we discuss here (hereafter

V1).

254

After discovering the binary in the LBT data, we also ran ISIS and

DAOPHOT/ALLSTAR on the MDM data to extract the light curve of the

long-period binary. The variability data from LBT and MDM were complemented

with single-epoch archival imaging of the field obtained from the SDSS Data Release

6 (Adelman-McCarthy et al. 2008) in the gri bands (UT Nov. 30, 2003), and the

HST/ACS Wide Field Camera (GO proposal 10605, PI E. Skillman) in the F555W

and F814W filters (UT Mar. 23, 2006). The high resolution HST/ACS images

(FWHM ∼ 0.′′1, corresponding to ∼ 2 pc at the distance of Holmberg IX) show that

the binary is spatially coincident with a stellar association in the dwarf galaxy.

Figure 7.1 shows the phased V -band light curve and B − V color curve of

the eclipsing binary system. We include all the LBT, MDM, SDSS and HST/ACS

V -band photometry. The LBT and MDM photometry have been calibrated using

SDSS photometry of several relatively bright (r ∼< 21.0 mag) and unsaturated stars

in the field, transforming the gr magnitudes to standard BV magnitudes with the

transformations presented in Ivezic et al. (2007). The rms deviations of the absolute

calibration are 0.02 − 0.03 mag for LBT-BV , and ∼ 0.05 mag for MDM-V . The

SDSS gri photometry of the binary was extracted in the same way as for LBT and

MDM, using the DAOPHOT/ALLSTAR package to obtain instrumental magnitudes

calibrated using absolute photometry of the bright stars in the field. Our g and

r magnitudes of the binary from the SDSS data are 0.2 and 0.5 mag brighter,

respectively, than the magnitudes reported in the SDSS-DR6 catalog, while the

255

i-band magnitude agrees at the 1% level. We think this is due to problems in the

SDSS photometry for faint sources in a crowded field (e.g., Smolcic et al. 2007). The

details of the HST photometry can be found in Weisz et al. (2008, in preparation).

In Figure 7.2 we show the position of the binary in the color magnitude diagrams

(CMD), obtained from calibrated LBT and HST/ACS photometry. The CMDs show

the well-populated blue and red supergiant sequences in Holmberg IX. The binary is

among the most luminous stars in this dwarf galaxy, with MV ∼ −7.1 mag, and it

has clearly evolved from the main-sequence. With such a high intrinsic luminosity,

the binary is bound to be massive. After correcting for Galactic foreground

extinction, the B − V and V − I colors are consistent with an effective temperature

of Teff = 4800 ± 150 K (Houdashelt et al. 2000). Both components seem to be

G-type yellow supergiants given the equal depths of the eclipses and the lack of color

variations (see Figure 7.1).

We used the eclipsing binary model-fitting program NIGHTFALL1 to model

the V -band light curve. As shown in Figure 7.1, we obtain a good fit to the light

curve with an overcontact configuration where both stars are overflowing their Roche

lobes. We fixed the effective temperature of the primary at T1 = 4800 K obtained

from the colors. We assumed equal masses for the stars, a linear limb-darkening law,

circular orbits, and synchronous rotation. We fit for four parameters: the Roche

lobe filling factors, the inclination, and the temperature of the secondary. The time

1http://www.hs.uni-hamburg.de/DE/Ins/Per/Wichmann/Nightfall.html

256

of the primary eclipse and the period were determined externally and were fixed for

these fits. The main parameters of the binary are listed in Table 7.1. The light curve

shows a hint of the O’Connell (1951) effect, in which the maxima (out-of-eclipse

regions) show a difference in brightness (e.g., Pilecki et al. 2007).

We searched the literature for other examples of evolved, massive eclipsing

binaries in the yellow supergiant phase and found none2. We also searched the

available catalogs of eclipsing binaries in the LMC and SMC. The MACHO catalog

of eclipsing binaries in the LMC (Derekas et al. 2007) contains 25 contact systems

with red colors (i.e., evolved), (V −R) > 0.5 mag, and periods > 200 days. However,

these systems have absolute magnitudes MV ∼> −4 (V ∼> 14.5 mag), that are

∼ 3 mag fainter than the yellow supergiant eclipsing binary in Holmberg IX3. The

All Sky Automated Survey (ASAS; Pojmanski 2002) contains complete Southern

sky coverage for V < 15 mag. To our surprise, we found in the ASAS catalog a

luminous (MV ∼ −7.5, Vmax ∼ 11.5 mag), 181 day period contact eclipsing binary

in the SMC. The star, SMC R47 (α = 01h29m17.s26, δ = −7243′20.2′′), had been

spectroscopically classified as an FO supergiant (Teff ≃ 7500 K) with emission lines

2Note, however, that a possible Galactic counterpart is the contact binary BM Cas (Fernie &

Evans 1997), composed by an A7 Iab supergiant (MV ≃ −6.3) and a late-type giant.3Mennickent et al. (2006) obtained spectroscopy of 17 “peculiar” periodic variables in the

SMC from the OGLE database, and found a 184 day period eclipsing binary composed by two

yellow supergiants (F5Ie + G5-K0I). However, this system is ∼ 2 mag fainter than the binary in

Holmberg IX.

257

by Humphreys (1983). The ASAS V -band light curve of SMC R47, obtained between

December 2002 and June 2006, and the fit obtained with NIGHTFALL are shown in

Figure 7.3. The best-fit eclipsing binary model requires a contact configuration, with

non-zero eccentricity to account for the difference in timing between the eclipses.

Even though we selected a clean part of the full ASAS light curve, there seems to

be intrinsic variability from the binary components. The main parameters of this

eclipsing binary are in Table 7.1. While substantially hotter than the Holmberg IX

binary, it does not lie on the SMC blue supergiant sequence (Grieve & Madore 1986).


An eclipsing binary is the best explanation for the light curve of the brightest

variable we have discovered in our LBT variability survey of the dwarf irregular

companion of M81 Holmberg IX. The other possible explanation for the periodic

variability of V1 is a long-period (P = 135 days) Cepheid. Such long-period

Cepheids (P > 100 days) have been observed in dwarf galaxies like the LMC and

SMC (e.g., Freedman et al. 1985), NGC 55 ( e.g., Pietrzynski et al. 2006), NGC 300

(e.g., Pietrzynski et al. 2002), NGC 6822 (e.g., Pietrzynski et al. 2004), IC 1613

(e.g., Antonello et al. 1999), and I Zw 18 (Aloisi et al. 2007). The magnitude of

V1 is consitent with the magnitude of a Cepheid with P = 135 days (MV ≃ −7.0),

extrapolating the period-luminosity relationship of Fouque et al. (2007). However,

258

while a few of these long period Cepheids have quasi-sinusoidal light curves that

are nearly symmetric under a time reversal, they all have larger amplitudes in

bluer bands (B amplitudes 1.3 − 1.6 times V amplitudes) due to the changes in

the effective temperature as the star pulsates (e.g., Freedman et al. 1985; Madore

& Freedman 1991). Spectroscopy of V1, while challenging, would eliminate any

remaining ambiguities in classifying this system.

We can safely rule out the possibility that the eclipsing binary is in our Galaxy.

Using the period, estimated effective temperature, magnitude at maximum, and

assuming that the stars are of similar size in a contact configuration, we can use

Kepler’s law and their total surface brightness to estimate a distance to the system

of D ≃ (Mtotal/M⊙)1/3 Mpc, where Mtotal = M1 + M2 (e.g., Gaposchkin 1962).

Conversely, we can estimate the total mass of the binary system by assuming the

distance, Mtotal ≃ 45 (D/3.6 Mpc)3 M⊙. This should be taken as a rough estimate

because of the overly simple model − to accurately constrain the total mass of the

system, and its components, we need radial velocity measurements. Another piece of

evidence that puts the binary system in Holmberg IX is its spatial coincidence with

a stellar overdensity in the dwarf, observed in the HST/ACS images.

259

We expected that such systems were rare4, but were surprised to find none

in the literature. However, we found a similar eclipsing binary system in the SMC

(SMC R47) searching through the ASAS catalog. From the absolute magnitudes

of both binaries and their colors, we estimate that at least one of the stars in each

binary is ∼ 15 − 20 M⊙ (main-sequence age ∼ 10 − 15 Myr) using the evolutionary

tracks for single stars of Lejeune & Schaerer (2001) (see Figure 7.4).

The stellar evolutionary path of stars of a given mass in binary systems can differ

significantly from their evolution in isolation (e.g., Paczynski 1971). In particular,

binary interactions through mass loss, mass accretion, or common-envelope evolution,

play a very important role in the pre-supernova evolution (e.g., Podsiadlowski et

al. 1992). Most of the massive stars with masses 30M⊙ ∼> M ∼> 8 M⊙ are expected

to explode as supernova when they are in the red supergiant stage, with a small

contribution from blue supergiants (e.g., SN 1987A; West et al. 1987). Surprisingly,

Li et al. (2005) identified the progenitor of the Type IIP supernova 2004et in

pre-explosion archival images and determined that it was a yellow supergiant with

a main-sequence mass of ∼ 15 M⊙. Also, the position in the CMD of the likely

progenitor of the Type IIP supernova 2006ov (see Fig. 10 in Li et al. 2007) is

4While the relative numbers of eclipsing binaries is a much more complicated problem, we note

that the relative abundances of red, blue and yellow supergiants is 4:13:1 for the Geneva evolutionary

track (Lejeune & Schaerer 2001) of a single, non-rotating star with M = 15M⊙ and Z = 0.004.

260

remarkably similar to the position of the eclipsing binary in Holmberg IX (see

Figure 7.4).

We propose that the binary we discovered in Holmberg IX and the binary found

in the SMC5 are the kind of progenitor objects of supernovae like SN 2004et and

SN 2006ov that appeared to be the explosions of yellow supergiants. A close binary

provides a natural means of slowing the transition from blue to red, allowing the

star to evolve and then explode as a yellow supergiant. As the more massive star

evolves and expands, the Roche lobe limits the size of the star forcing it to have a

surface temperature set by the uncoupled core luminosity and the size of the Roche

lobe. It can expand further and have a cooler envelope only by becoming a common

envelope system, which should only occur as the secondary evolves to fill its Roche

lobe. This delayed temperature evolution allows the core to reach SN II conditions

without a red envelope.

5A possible earlier stage of these yellow supergiant binaries might be represented by the B-

supergiant pair HD 1383 (Boyajian et al. 2006).

261

Fig. 7.1.— Phased V -band light curve of the 270.7 day period eclipsing binary inHolmberg IX (top panel) and its B − V color evolution (lower panel). The differentsymbols correspond to photometry from different telescopes: LBT (filled circles),MDM (open triangles), HST/ACS (filled square) and SDSS (open square). The solid

line shows an overcontact eclipsing binary model that best fits the light V light curve.

262

Fig. 7.2.— Color Magnitude Diagrams of stars in the field of Holmberg IX obtainedfrom the LBT-BV reference images (left panel) and HST/ACS-V I single-epochobservations (right panel). The ACS CMD shows well-defined stellar sequences forthe main sequence (MS) and the evolved blue (BSG) and red supergiants (RSG). Theeclipsing binary (pentagon) lies between the blue and red supergiant sequences.

263

Fig. 7.3.— Phased V -band light curve of the long-period, evolved eclipsing binarySMC R47 obtained from the ASAS catalog (Pojmanski 2002). The star was classifiedspectroscopically as an F0 supergiant by Humphreys (1983). The solid line shows thecontact eclipsing binary model that best fits the light curve. The lower panel showsthe residuals of the fit (observed-calculated).

264

Fig. 7.4.— CMD of Holmberg IX from HST/ACS V and I photometry. Theconnected filled symbols show the position of the evolved eclipsing binaries inHolmberg IX (filled pentagons) and the SMC (filled triangles) at maximum andminimum. The lines show evolutionary tracks with extended mass-loss from theGeneva group (Lejeune & Schaerer 2001) for single stars with masses between 12–25 M⊙, assuming two different metallicities: 1/3 solar (dashed) and 1/5 solar (solid).We use a distance modulus of µ = 27.80 mag to Holmberg IX, and Galactic color-excess E(B − V ) = 0.08 mag, to put the evolutionary tracks in the diagram.

265

Parameter Holmberg IX V1 SMC R47

Period, P 270.7 ± 2.3 days 181.58 ± 0.16 days

Time of primary eclipse, Tprim 2454186.0 ± 0.6 2452073.1 ± 0.2

Inclination, i 55.7 ± 0.6 82.2 ± 0.2

Primary temperature, T1 4800 ± 150 K 7500 ± 100 K

Temperature ratio, T2/T1 1.05 ± 0.05 1.17 ± 0.02

Eccentricity, e 0.00 0.039 ± 0.002a

Roche Lobe Filling factorsb 1.23 ± 0.02 1.02 ± 0.02

Semi-major axis, ac 547 R⊙ 418 R⊙

Note. — The mass ratio was fixed at q = 1 for fitting both light curves.

aA non-zero eccentricity is required to fit the difference in timing

between the primary and secondary eclipses. The best-fit longitude of

the periastron is w = 168.6 ± 1.1 .

bRatio of stellar to Roche lobe polar radius for each star.

cSeparation between the stars assuming a total mass of 30M⊙ for each

system.

Table 7.1. Best-fit Binary Model Parameters.

266

Appendix A

A. Extreme-AGB Star Variability

Section 4.4.4 presents a discussion of the variability of the 18 reddest sources

selected as EAGB stars (see Table 4.2 for their photometry). In this Appendix (in

Figs. A.1 & A.2) we present the lightcurves for all 18 sources. See the large open

circles in Figures 4.2, 4.4, and 4.7, as well as the left panel of 4.5 for a summary of

their colors, SEDs, and RMS variability.

267

Fig. A.1.— Lightcurves for sources S1−S9 of the 18 reddest sources (sorted by color)with M4.5 < −10 and [3.6] − [4.5] > 1.5 in the M33 MIR color-magnitude diagram(open circles in Fig. 4.2; remaining source lightcurves are shown in Fig. A.2). Foreach source, the top and bottom panels show the absolute 4.5µm magnitude and the[3.6]− [4.5] color variation, respectively, as a function of time. More than a magnitudevariation on a timescale of 100 − 1000 days is common.

268

Fig. A.2.— Same as Figure A.1, but for sources S10-S18 of the 18 reddest sourceswith M4.5 < −10 in the M33 MIR color-magnitude diagram (see Figs. 4.2 & 4.3).

269

Appendix B

LBV Candidate Variability

Like Appendix A, here we present the lightcurves for the 16 LBV candidates

from Massey et al. (2007) that have been matched to the MIR point source catalog,

as described in §4.4. Table 4.3 lists photometry for these sources. See Figures 4.2,4.4,

4.5, and 4.7 for a summary of their colors, SEDs, and RMS variability properties.

270

Fig. B.1.— Lightcurves at 4.5µm for 9 of the 16 LBV candidates of Massey etal. (2007). Although there are several exceptions (notably J013335.14+303600.4above and J013429.64+303732.1 in Fig. B.2), the LBV candidates do not varysignificantly in absolute magnitude or color (see also Fig. 4.7). As shown inFig. 4.2 (open triangles), the LBV candidates with high luminosities at 4.5µm arecharacteristically more red than those with larger M4.5.

271

Fig. B.2.— Same as Figure B.1, but for the remaining 7 LBV candidates of Masseyet al. (2007), matched to our 4.5µm catalog.

272

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