KOI-‐254 / Kepler 45 – A Hot Jupiter Orbi6ng a Metal-‐Rich M dwarf (Johnson et al. 2012)
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KOI-531 (M0V)
KOI-2764 (M0V)
KOI-2845 (M0V)
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M Giant KOIsKOI-977 (MIII)
KOI-3497 (MIII)
Characterizing the Cool KOIs An Infrared Spectroscopic Survey of Kepler M Dwarf Planet-‐Candidate Hosts
The Astronomical Journal, 143:111 (11pp), 2012 May Johnson et al.
-2 -1 0 1 2Hours from Mid Transit
0.88
0.90
0.92
0.94
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Rela
tive
Flux
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ffset
Figure 7. Kepler (upper, blue) and Nickel (lower, red) light curves, phased at thephotometric period. The Nickel light curve has been offset artificially for clarity.The best-fitting light curve models are shown for each data set (see Section 4),and the residuals are shown beneath each light curve.(A color version of this figure is available in the online journal.)
at Lick Observatory, and our 14 RV observations acquired withKeck/HIRES. The Nickel and HIRES observation timestampswere converted to BJDUTC to match Kepler MAST data usingthe techniques of Eastman et al. (2010).
We fitted the Kepler and Nickel light curves using version3.01 of the Transit Analysis Package (Gazak et al. 2011), whichuses the analytic eclipse model of Mandel & Agol (2002). Forthe Kepler transits we resampled the model to a cadence of60 s before rebinning to the 29.4 minute observing cadenceto account for long integration light curve distortions (Kipping2010). We determined the best-fitting parameters and their un-certainties using the same Metropolis–Hastings implementationdescribed in Section 3, with which we employ a Daubechiesfourth-order wavelet decomposition likelihood function (Carter& Winn 2009). Wavelet decomposition techniques provide in-creased confidence in derived MCMC uncertainties over thetraditional χ2 likelihood by allowing parameters which mea-sure photometric scatter (uncorrelated Gaussian σw, and 1/fcorrelated red σr ) to evolve as free parameters. The techniquerecovers the χ2 likelihood in the case where σr = 0 and σw islocked at a value characteristic to the observed data. For the RVdata we fitted a Keplerian model using the partially linearizedscheme of Wright & Howard (2009).
Of the fifteen parameters in this technique, thirteen vary freelywithin our MCMC analysis: the period P, Inclination i, thescaled semimajor axis aR, the radius ratio Rp/Rs, times of mid-transit Ttr, eccentricity e, argument of periastron ω, σw, σr , RVamplitude K, the systemic velocity offset γ , and two parametersto account for global linear trends in the data normalization.
Table 3KOI-254 Transit Mid-times and Ephemeris Residuals
Tmid (BJD-2450000.0) Tmid−Ephemeris Telescope
54964.5368 ± 0.0015 −0.00048 ± 0.0018 K54966.99228 ± 0.00081 −0.00024 ± 0.0013 K54969.44698 ± 0.00084 −0.00076 ± 0.0013 K54971.90303 ± 0.00092 0.000059 ± 0.0014 K54974.35835 ± 0.00095 0.00015 ± 0.0014 K54976.81357 ± 0.00092 0.00014 ± 0.0014 K54979.26833 ± 0.00083 −0.00033 ± 0.0013 K54981.72302 ± 0.00096 −0.00087 ± 0.0014 K54984.17923 ± 0.00086 0.00011 ± 0.0013 K
54986.6343 ± 0.0015 −0.000025 ± 0.0018 K54989.08909 ± 0.00091 −0.00049 ± 0.0014 K54991.54493 ± 0.00068 0.00012 ± 0.0012 K
54993.9996 ± 0.0010 −0.00043 ± 0.0014 K54996.45456 ± 0.00096 −0.00070 ± 0.0014 K
55003.8211 ± 0.0010 0.00015 ± 0.0014 K55006.27645 ± 0.00085 0.00027 ± 0.0013 K55008.73088 ± 0.00085 −0.00054 ± 0.0013 K55011.18647 ± 0.00081 −0.00017 ± 0.0013 K
55013.6420 ± 0.0010 0.00012 ± 0.0014 K55018.55210 ± 0.00098 −0.00023 ± 0.0014 K
55021.0084 ± 0.0010 0.00088 ± 0.0014 K55023.46242 ± 0.00084 −0.00036 ± 0.0013 K55025.91715 ± 0.00073 −0.00087 ± 0.0012 K55028.37311 ± 0.00091 −0.00014 ± 0.0013 K55030.82892 ± 0.00083 0.00044 ± 0.0013 K
55033.2832 ± 0.0020 −0.00051 ± 0.0022 K55035.7396 ± 0.0011 0.00069 ± 0.0015 K
55038.19468 ± 0.00084 0.00052 ± 0.0013 K55040.64897 ± 0.00062 −0.00042 ± 0.0012 K55043.10475 ± 0.00085 0.00013 ± 0.0013 K55045.55956 ± 0.00091 −0.00029 ± 0.0014 K55048.01603 ± 0.00081 0.00095 ± 0.0013 K55050.47094 ± 0.00093 0.00063 ± 0.0014 K55052.92573 ± 0.00088 0.00020 ± 0.0013 K
55055.3827 ± 0.0011 0.0019 ± 0.0015 K55057.83570 ± 0.00070 −0.00030 ± 0.0012 K55060.29115 ± 0.00093 −0.000073 ± 0.0014 K55062.74626 ± 0.00089 −0.00019 ± 0.0013 K55065.20262 ± 0.00087 0.00094 ± 0.0013 K55067.65679 ± 0.00091 −0.00012 ± 0.0014 K55070.11231 ± 0.00071 0.00017 ± 0.0012 K55072.56725 ± 0.00085 −0.00012 ± 0.0013 K55075.02311 ± 0.00089 0.00051 ± 0.0013 K55077.47754 ± 0.00097 −0.00029 ± 0.0014 K
55079.9343 ± 0.0010 0.0012 ± 0.0014 K55082.38913 ± 0.00094 0.00084 ± 0.0014 K55084.84365 ± 0.00086 0.00013 ± 0.0013 K55087.29916 ± 0.00084 0.00041 ± 0.0013 K55089.75411 ± 0.00099 0.00014 ± 0.0014 K55742.84449 ± 0.0027 −0.00047 ± 0.0029 N
Note. K—Kepler, N—Nickel Z-band.
The remaining two limb-darkening coefficients evolve undernormal priors. For the Nickel Z-band data we adopted fromClaret (2004): µ1 = 0.353 ± 0.35, µ2 = 0.255 ± 0.025.For the Kepler data we used the coefficients listed by Sing(2010): µ1 = 0.521 ± 0.056, and µ2 = 0.225 ± 0.052. Itis important to note that our joint fitting procedure alloweduncertainties in the orbital eccentricity to propagate into thedetermination of the Keplerian orbit parameters and the scaledsemimajor axis aR.
We ran 40 independent MCMC chains each with 5 × 105
links for a total of 1.4 × 107 total inference links after removing
7
KOI 3497 shows a deep CO (2-‐0) band head indicaHve of a giant star, but deep Na and Ca lines are consistent with a dwarf. As noted by Rojas-‐Ayala, an early + late M dwarf binary can create deep CO. Robo-‐AO image of KOI 3497 likely confirming the suspicion! But which star has the planet candidate? Image courtesy of the Robo-‐AO Team, including Christoph Baranec, Reed Riddle, and Nick Law.
Philip S. Muirhead (BU), JulieWe Becker (Caltech), Bárbara Rojas-‐Ayala (CAUP), Andrew Vanderburg (CfA), Jon SwiZ (Caltech) , Gregory Feiden (Uppsala), Ellen Price (Caltech), Rachel Thorp (Caltech), Katherine Hamren (UCSC), EvereW Schlawin (Cornell), Kevin R. Covey (Lowell), John Asher Johnson (CfA), James P. Lloyd (Cornell)
H-‐ and K-‐Band Spectra of Cool KOIs taken with the TripleSpec Spectrograph on the Palomar 200-‐inch Hale Telescope Ordered in increasing TEff
Science Highlights!
KOI 961 / Kepler 42 3 short-‐period sub-‐Earths orbiHng an M4V star! Muirhead et al. (2012)
KOI-‐961 / Kepler 42 system and Jupiter’s Galilean moons Orbital scale is 5 x object size scale
Spectra provided accurate stellar properHes for analysis of KOI-‐952’s 5 planets, aka Kepler 32: “A Prototype for the FormaHon of Compact Planetary Systems
Throughout the Galaxy” (SwiZ et al. 2013).
Spectra provided an accurate metallicity determinaHon for KOI-‐254, aka Kepler 45, the only M dwarf known to host a hot Jupiter (Johnson et al. 2012). KOI-‐254 is significantly metal-‐rich ([Fe/H] = +0.32), further supporHng the theory that metallicity plays a significant role in the formaHon of Jovian-‐mass planets, even around M dwarfs where Jovian-‐mass planets are rare (Johnson et al. 2010, Rojas-‐Ayala et al. 2010, 2012, Mann et al. 2012). LeF: Kepler and Lick Observatory transit measurements Johnson et al. (2012)
?
3000 3500 4000 4500 5000KOI TEff from This Work
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LeF: Stellar effecHve temperature, metallicity and radius determinaHons for the stars in this sample. We determined stellar effecHve temperature and metallicity and using the calibraHons of Rojas-‐Ayala et al. (2010, 2012). We then interpolate those values onto new 5-‐Gyr Dartmouth isochrones calculated by Gregory Feiden (Uppsala University). The new Dartmouth isochrones include stars with effecHve temperatures less than 3000 K (Muirhead et al. in prep). Right: Comparison to Dressing & Charbonneau (2013), who used photometry to determine Cool KOI properHes. Our results generally show good agreement,; however, there is a slight metallicity dependence to our discrepancies. Dressing & Charbonneau (2013) assumed a strict prior for their metallicity determinaHons due to degeneracies with temperature when using photometry:
1.6 1.8 2.0 2.2 2.4Wavelength (µm)
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(erg
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M Dwarf KOIsKOI-4290 (M4V)
KOI-2842 (M4V)
KOI-961 (M4V)
KOI-1725 B (M4V)
KOI-2704 (M4V)
KOI-3749 (M3V)
KOI-249 B (M3V)
KOI-1702 (M3V)
KOI-3119 (M3V)
KOI-463 (M3V)
1.6 1.8 2.0 2.2 2.4Wavelength (µm)
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(erg
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KOI-256 (M3V)
KOI-2453 (M3V)
KOI-2542 (M2V)
KOI-2705 (M2V)
KOI-1422 (M2V)
KOI-1146 (M2V)
KOI-1686 (M2V)
KOI-249 A (M2V)
KOI-899 (M2V)
KOI-2626 (M2V)
1.6 1.8 2.0 2.2 2.4Wavelength (µm)
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(erg
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KOI-936 (M2V)
KOI-854 (M2V)
KOI-1907 (M2V)
KOI-1725 A (M2V)
KOI-2662 (M2V)
KOI-2715 (M1V)
KOI-1902 (M1V)
KOI-1681 (M1V)
KOI-3444 (M1V)
KOI-596 (M1V)
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KOI-3144 (M1V)
KOI-2862 (M1V)
KOI-3263 (M1V, EB)
KOI-781 (M1V)
KOI-1201 (M1V)
KOI-818 (M1V)
KOI-1843 (M1V)
KOI-886 (M1V)
KOI-3034 (M1V)
KOI-1867 (M1V)
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KOI-3090 (M1V)
KOI-952 (M1V)
KOI-739 (M1V)
KOI-247 (M1V)
KOI-1459 (M1V)
KOI-478 (M1V)
KOI-252 (M1V)
KOI-817 (M1V)
KOI-2058 (M1V)
KOI-571 (M1V)
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KOI-3284 (M1V)
KOI-947 (M1V)
KOI-2156 (M1V)
KOI-2006 (M1V)
KOI-2036 (M1V)
KOI-253 (M1V)
KOI-2650 (M1V)
KOI-255 (M1V)
KOI-1078 (M1V)
KOI-254 (M1V)
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KOI-2238 (M1V)
KOI-3010 (M1V)
KOI-251 (M1V)
KOI-4427 (M1V)
KOI-2347 (M1V)
KOI-1397 (M1V)
KOI-2329 (M1V)
KOI-1868 (M1V)
KOI-1879 (M1V)
KOI-2179 (M1V)
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KOI-248 (M1V)
KOI-314 (M1V)
KOI-2306 (M0V)
KOI-2191 (M0V)
KOI-4252 (M0V)
KOI-4875 (M0V)
KOI-1649 (M0V)
KOI-1427 (M0V)
KOI-250 (M0V)
KOI-2090 (M0V)
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KOI-2926 (M0V)
KOI-3282 (M0V)
KOI-898 (M0V)
KOI-2839 (M0V)
KOI-812 (M0V)
KOI-1880 (M0V)
KOI-1408 (M0V)
KOI-2130 (M0V)
KOI-1141 (M0V)
KOI-2057 (M0V)
Robo-‐AO image of KOI-‐3497 A false giant star!
The Astrophysical Journal, 764:105 (14pp), 2013 February 10 Swift et al.
1:2
2:3
Kepler 32 Planetary System
0.7d
2.9d
5.9d
8.8d
22.8d
f e b c d
0.13AU
0.07AU
0.05AU
0.03AU
0.01AU
planet size x 80
f e b c d
Figure 4. Depiction of the Kepler-32 planetary system with the star and orbits drawn to scale. The relative sizes of the planets are shown at the bottom of the figurescaled up by a factor of 80 in relation to their orbits.(A color version of this figure is available in the online journal.)
of M ∝ Rγp , with γp = 1.5–1.9, similar to the value of 2.06for the six solar system planets bounded by Mars and Saturn(Lissauer et al. 2011).
The above stated densities imply that Kepler-32 b and care composed of a significant amount of volatiles. UsingEquations (7) and (8) of Fortney et al. (2007), we find thatif Kepler-32 b and c had no atmospheres, then they wouldbe expected to contain ∼96% and ∼56% volatiles. However,given the equilibrium temperatures of Kepler-32 b and c, a largefraction of their volatile content likely exists in the form of anatmosphere.
4.3. Atmospheric Evolution
The proximity of the Kepler-32 planets to their host starsuggest significant atmospheric evolution due to evaporation,outgassing, or both processes. The equilibrium temperatureof Kepler-32 f is ∼1100 K and its radius is measured to be0.81 R⊕. For a planet this small with such a high equilibriumtemperature, the atmospheric mass fraction would have to bevery small, ∼10−5 (Rogers et al. 2011). Using an extremeultraviolet luminosity of Kepler-32, LEUV ≈ 1026.6 (Hodgkin &Pye 1994), and following Lecavelier Des Etangs (2007) usinga conservative mass-loss efficiency of ϵUV = 0.1, we derive anatmospheric mass loss of ∼108 g s−1. Thus, the timescale to loseits atmosphere is more than 100 times shorter than the age of theKepler-32 system. We therefore conclude that the Kepler-32 fcontains no atmosphere.
Given the size and equilibrium temperature of Kepler-32 e, itsatmospheric mass fraction must also be small, Ma/Mp ∼ 10−4,while the present-day atmospheric mass-loss rate is between 107
and 108 g s−1. The timescale for the complete loss of the Kepler-32 e atmosphere is calculated to be between 0.2 and 2 Gyr.Therefore, Kepler-32 e must have lost a significant fraction ofany atmosphere it started with.
The total atmospheric mass loss for the other three planetsis at least ∼10−4 M⊕ for reasonable choices of planetary mass.
If these planets have relatively low density cores (ice and rock)and started out with large atmospheres, then they could havesuffered considerable atmospheric evolution due to the heatingby Kepler-32. Thus, the observed sizes of the Kepler-32 planetsare likely determined in part by the extreme ultraviolet andX-ray luminosity of their host star. However, the mass estimatesfrom Section 4.2 suggest that Kepler-32 b is 10% less massivethan Kepler-32 c while being 10% larger and 25% closer toKepler-32, hinting that the mass–radius relation for the Kepler-32 planets is not determined solely by a simple atmosphericevolution model.
4.4. Kepler-32 Planetary System Architecture
The physical characteristics of the Kepler-32 planets aresummarized in Table 3 and the remarkably compact and orderlyarchitecture of the system is shown schematically in Figure 4.As mentioned above, three of the planets lie within 2% of a1:2:3 period commensurability. Kepler-32 e and b have a periodratio of 2.038, which is 1.9% longward of commensurability,while Kepler-32 b and c have a period ratio of 1.483, or 1.1%shortward of commensurability.
Planets within a mean motion resonance can stray a fewpercent from commensurability and maintain the libration ofresonant angles (Murray & Dermott 1999). However, withoutdetailed knowledge of the individual orbits, it is not possibleto determine with certainty if a planet pair is in a resonantconfiguration. Therefore, we assess the significance of the nearcommensurability of Kepler-32 e, b, and c using a probabilisticargument.
We randomly populate five planet systems with periodsbetween the inner and outermost planets in the Kepler-32system, enforcing separations larger than 2
√3 for every pair
of neighboring planets (Gladman 1993) and larger than 9 forchains of planets (Chambers et al. 1996; Smith et al. 2009;Lissauer et al. 2011) in units of mutual Hill radii. In this section,a mass–radius relationship of M ∝ R2.06 (Lissauer et al. 2011) is
6
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Dartmouth 5-Gyr Isochrone
KOI 256
Muirhead et al. (2013)
KOI-‐256: An M Dwarf / White Dwarf Binary with Gravita6onal Microlenseing
KOI-‐3497: A False Giant Star
Comparison to Dressing & Charbonneau (2013)
KOI-‐961 / Kepler 42: A Mid-‐M dwarf with 3 Short-‐period Sub-‐earths
KOI-‐952 / Kepler 32: A Compact System of 5 Planets Swfit et al. (2013)
Stellar Proper6es using New Dartmouth Isochrones
Empirically derived (eclipsing binary)