Draft version July 22, 2020Typeset using LATEX twocolumn style in AASTeX62

Detecting Exoplanets Using Eclipsing Binaries as Natural Starshades

Stefano Bellotti,1, 2, 3 Ann I. Zabludoff,4 Ruslan Belikov,5 Olivier Guyon,4, 6 and Chirag Rathi4

1DARK, Niels Bohr Institute, University of Copenhagen, Lyngbyvej 2, 4. sal, 2100 Copenhagen Ø, Denmark2Universite de Toulouse, UPS-OMP, IRAP, 14 avenue E. Belin, Toulouse, F-31400 France

3CNRS, IRAP/UMR 5277, Toulouse, 14 avenue E. Belin, F-31400 France4Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson AZ 85721

5NASA Ames Research Center, Moffett Field, CA 94035, USA6NAOJ

(Received January 1, 2020; Accepted July 18, 2020)

ABSTRACT

We investigate directly imaging exoplanets around eclipsing binaries, using the eclipse as a natural

tool for dimming the binary and thus increasing the planet to star brightness contrast. At eclipse, the

binary becomes point-like, making coronagraphy possible. We select binaries where the planet-star

contrast would be boosted by > 10× during eclipse, making it possible to detect a planet that is

& 10× fainter or in a star system that is ∼ 2-3× more massive than otherwise. Our approach will

yield insights into planet occurrence rates around binaries versus individual stars. We consider both

self-luminous (SL) and reflected light (RL) planets. In the SL case, we select binaries whose age is

young enough so that an orbiting SL planet would remain luminous; in U Cep and AC Sct, respectively,

our method is sensitive to SL planets of ∼4.5MJ and ∼9MJ with current ground- or near-future space-

based instruments, and ∼1.5MJ and ∼6MJ with future ground-based observatories. In the RL case,

there are three nearby (. 50 pc) systems—V1412 Aql, RR Cae, RT Pic—around which a Jupiter-like

planet at a planet-star separation of & 20 mas might be imaged with future ground- and space-based

coronagraphs. A Venus-like planet at the same distance might be detectable around RR Cae and RT

Pic. A habitable Earth-like planet represents a challenge; while the planet-star contrast at eclipse and

planet flux are accessible with a 6-8m space telescope, the planet-star separation is 1/3 - 1/4 of the

angular separation limit of modern coronagraphy.

Keywords: eclipsing binaries — extrasolar planets — contrast — angular separation — catalogs

1. INTRODUCTION

Coronagraphs, nulling interferometry, and man-made

starshades are the existing strategies for imaging exo-

planets directly. Is there a way to dramatically improve

these techniques? Here we consider using the eclipse in

an eclipsing binary system to dim the observed bright-

ness of the primary and increase the planet to star flux

contrast, i.e., we explore the possibility of employing a

natural starshade as a tool to find additional exoplan-

ets around binaries via direct imaging, along with the

mentioned techniques.

Corresponding author: Stefano Bellotti

info.stefanobellotti@gmail.com

The NASA Exoplanet Archive (Akeson et al. 2013)

lists 270 binary systems with exoplanets, but only six

binaries have planets detected via direct imaging (Bur-

gasser et al. 2010; Kuzuhara et al. 2011; Currie et al.

2014; Kraus et al. 2014; Gauza et al. 2015; Janson et

al. 2019). All six are self-luminous planets with mini-

mum masses intermediate between 6 and 20 MJ . Our

alternate direct imaging method can explore a new pa-

rameter space by 1) targeting binaries or stars that are

unusual compared to previously observed exoplanet sys-

tems and 2) making different (fainter) types of planets

accessible.

Planets in binary systems could represent an impor-

tant fraction of planet demography, especially given that

∼ 45% of Sun-like stars in the Galactic field are part

of a multiple system (Raghavan et al. 2010). Parker

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& Quanz (2013) estimate statistically the percentage of

Solar System analogues, defined either as an individual

G dwarf or a binary system with separation > 100-300

AU and a G dwarf component, that host exoplanets.

This percentage declines from 65-95% to 20-65% from 1

to 100 AU for planets on circumprimary (S-type) orbits,

whereas it increases from 5-59% to 34-75% from 1 to 100

AU for planets on circumbinary (P-type) orbits.

The effect of binarity, relative to single stars, on planet

occurrence rates is uncertain. On one hand, the prox-

imity of a stellar companion could induce disk trunca-

tion (Jang-Condell 2015) and suppress planet formation

(Moe & Kratter 2019), reducing (by 0.3×) the occur-

rence rate compared to that in wider binary or indi-

vidual star systems (Kraus et al. 2016). On the other

hand, Matson et al. (2018) do not observe this suppres-

sion within ' 50 AU, and some planets have been dis-

covered orbiting in S-type configurations within tight

binary systems (e.g., Thebault & Haghighipour 2015).

Furthermore, the material-rich environments that form

massive stars and binaries may readily produce high-

mass protoplanetary disks and then gas giant planets

(Kennedy & Kenyon 2008).

Zorotovic & Schreiber (2013) analyze detached post-

common-envelope binaries and find that 90% of those

observed for ∼5 years have eclipse timing variations that

could be explained by a circumbinary companion. The

Search for Planets Orbiting Two Stars (SPOTS) survey

(Thalmann et al. 2014; Bonavita et al. 2016; Asensio-

Torres et al. 2018) constrains the frequency of wide (<

1000 AU) orbit substellar companions to between 0.9%

and 9%, consistent with that around single stars. The

combination of the low rate from the SPOTS survey

and the high frequency from the Zorotovic & Schreiber

(2013) study suggests a second generation scenario of

planet formation around post-common-envelope bina-

ries, i.e., the planet forms after the binary. In this con-

text, our approach has the potential not only to yield

further insights on planet occurrence rates around bina-

ries, but also to differentiate among theories of planet

formation in binary environments.

Using eclipsing binaries to image exoplanets directly

could also expand our knowledge of the kinds of binaries

around which planets can form and evolve. For exam-

ple, while planets have been discovered around eclipsing

binaries using the eclipse timing method (HW Vir, Lee

et al. 2009; DP Leo, Qian et al. 2010; NN Ser, Qian et

al. 2009; Beuermann et al. 2010; NY Vir, Qian et al.

2012a; RR Cae, Qian et al. 2012b), the host properties

are narrow and biased. The hosts are generally short-

period compact binaries with a low-mass star or a white

dwarf component, as the eclipse minimum can be timed

more precisely in these cases. Instead, for our method

to work efficiently, we require that the dip in magnitude

at eclipse is large enough to yield a substantial gain in

contrast, regardless of the type of components.

Here we present the application of the natural-

starshade method to both self-luminous (SL) and re-

flected light (RL) planets, in order to assess whether

eclipsing binaries represent competitive targets for po-

tential detections with current or future imaging tech-

nology. Figure 1 illustrates the observational space

and depicts previous, current, and future imaging fa-

cilities according to their actual or expected perfor-

mances. Some directly imaged self-luminous exoplanets

are shown in Figure 1 for comparison: β Pic b (Lagrange

et al. 2009, 2010), HD 95086 b (Rameau et al. 2013), κ

And b (Carson et al. 2013), HR 8799 b,c,d,e (Marois et

al. 2008, 2010), 51 Eri b (Macintosh et al. 2015), GJ 504

b (Kuzuhara et al. 2013), and Fomalhaut b (Kalas et

al. 2008), which could be surrounded by a cloud of dust

or a disk (Galicher et al. 2013; Lawler et al. 2015), or

due to a massive collision of two planetesimals (Gaspar

& Rieke 2020).

We also simulate the presence of an Earth-twin (i.e.,

same radius and albedo as Earth, and flux received from

the binary equal to the solar constant) around a sam-

ple of nearby (within 20 pc) stars ranging from M to

F type. No datapoint falls in the current technology

zone, meaning that the detection of an Earth-like planet

around these stars, by means of direct imaging of the re-

flected light, is not yet feasible. One problem is that an

Earth sibling around most M dwarfs would be located at

a planet to star separation well below the inner working

angle (IWA) of current coronagraphic instruments, i.e.,

below the detectable angular separation limit. Another

issue is that the combination of stellar luminosity and

simulated Earth luminosity places these planets below

achievable planet to star contrast levels.

Targeting an eclipsing binary has an advantage over a

regular binary due to the increased planet to star con-

trast at the moment of eclipse. This is particularly rel-

evant for planets observable in reflected light, as their

brightness would arise from both stars, while the ob-

served starlight would be only from the fainter compo-

nent of the system. In other words, the eclipse-induced

luminosity dip affects only the observed luminosity of

the binary system, whereas the reflected luminosity of

the planet remains unaltered.

This paper is structured so that we present a cata-

log of eclipsing binaries, along with the photometric,

astrometric, and spectroscopic properties compiled for

this study, in Section 2. We consider the detectabil-

ity of both self-luminous planets and of reflected light

Detecting Exoplanets with Eclipsing Binaries 3

-2 -1 0 1log Planet-Star Separation [arcsec]

-11

-9

-7

-5

-3

log

Plan

et-S

tar C

ontra

st

HD95086 bHR 8799 bHR 8799 c+d

HR 8799 e And b

GJ504 b51 Eri b

Fomalhaut b

Pic b

Proxima Cen "Earth"Barnard's Star "Earth"

Eri "Earth" Cet "Earth"

Altair "Earth"

Procyon "Earth"

Cen B "Earth" Cen A "Earth"Hypothetical

Earth twinsaround Mdwarfs...

...around Sun-like stars

2020s ground-based telescopesCurrent coronographsPrevious coronographs

2030s 4-m space-based telescopes2020s 1- to 2.4-m space-based telescopes30- to 45-cm space-based telescopes

mag-25 "Earth"mag-30 "Earth"Known SL planet

M

K

G

F

Figure 1. Planet-to-star contrast versus separation plane for the direct imaging method. The capabilities of existing orplanned instruments are bundled into different technology zones (colored regions): in pink, VLT-NACO (Lenzen et al. 2003;Rousset et al. 2003), Subaru-HiCIAO (Tamura et al. 2006), Keck-NIRC2 (McLean & Chaffee 2000); in yellow, VLT-SPHERE(Beuzit et al. 2008), Gemini-GPI (Macintosh et al. 2014), JWST -NIRCam (Krist et al. 2007); in green, EELT-EPIC (Kasperet al. 2008), TMT-PSI (Guyon et al. 2018), TMT-PFI and SEIT (Crossfield 2013); in blue, WFIRST -CGI (Bailey et al. 2018),LUVOIR-ECLIPS (Juanola-Parramon et al. 2019) and HabEx (Gaudi et al. 2020). The gap between 30- to 45-cm space basedtelescopes (lightest blue) and previous coronagraphs (pink) is not significant. The groups of instruments defining each coloredtechnology zone vary in their wavelength coverage, so our comparisons to them are rough guides to what is possible. Here,we plot self-luminous planets already discovered through direct imaging with infrared coronagraphs (black points) as well assimulated “Earths” that could be directly imaged in reflected light (colored points). The “Earths” are assumed to lie arounda sample of M–F type stars within 20 pc; their color encodes the spectral type of the host, while their size is scaled accordingto the apparent infrared magnitude of the planet. Not surprisingly, no “Earths” falls within the reach of existing instruments(yellow).

Jupiter-, Venus-, and Earth-like planets in Sections 3

and 4, respectively. We then describe the advantages

of our approach with respect to single stars and binary

systems that do not eclipse in Section 5 and report our

conclusions in Section 6.

2. ECLIPSING BINARY SAMPLE

We compile a list of eclipsing binary systems from

the Catalog of Algol Type Binary Stars (Budding et al.

2004) and the Catalog of Eclipsing Variables (Malkov et

al. 2006). To determine the best targets for directly

imaging exoplanets during eclipse, we consider those

binaries for which the depth of the primary minimum

(Dmag) is larger than 2.5, i.e., for which the luminosity

dimming factor is at least 10. Overall, 289 eclipsing bi-

nary systems satisfy the Dmag constraint: 58 from the

first catalog and 231 from the second (also considering

the latter’s updated version in Avvakumova et al. 2013).

In Table 3 in Appendix A, we list all these eclipsing bi-

nary systems, along with the photometric, astrometric,

and spectroscopic properties relevant for our work here.

If these two catalogs report significantly different values

of Dmag in the same photometric filter, we exclude that

binary from our subsequent analysis, but report it for

completeness in Table 3.

The Dmag > 2.5 criterion selects mostly classical Al-

gols, i.e., with an evolutionary class of SA (Avvakumova

et al. 2013). These binaries have a B- or A-type main

sequence accretor and a G- or K-type subgiant or giant

donor that is large enough to completely eclipse the pri-

mary. Observationally, classical Algols are characterized

4 Bellotti et al.

by a deep primary eclipse, shallow secondary eclipse,

and ellipsoidal modulations between eclipses due the gi-

ant filling its Roche lobe and having a distorted (non-

spherical) shape (Budding et al. 2004; Moe & Di Stefano

2015).

2.1. Distances

Distance is a key factor in limiting direct planet imag-

ing. Because distance data in both catalogs are incom-

plete, we retrieve parallaxes from the Gaia Second Data

Release (DR2) (Gaia Collaboration 2016, 2018). As

explained in Gaia Collaboration (2018), the parallaxes

might be associated with either the photocenter of the

system or one of the two components, because all Gaia

DR2 targets were treated as individual sources.

To check that these DR2 parallaxes are generally con-

sistent with previous measurements, we compare them

with Gaia DR1 or Hipparcos measurements (Figure 2).

The RMS of the residuals with respect to the 1-to-1

line is 0.81 mas, roughly consistent with the mean of

the plotted DR1 and Hipparcos measurement errors (0.5

mas) and about 10× larger than the mean of the DR2 er-

rors (0.05 mas). There is a slight and expected increase

in the scatter at small parallaxes, but, overall, the two

datasets are consistent within the published measure-

ment uncertainties. There are no large systematics.

For binary systems characterized by an orbital period

longer than 2 yr, there might be a mismatch between the

parallaxes or proper motions listed in Gaia DR2 with

respect to the Tycho-Gaia astrometric solution (TGAS)

subset of Gaia DR1 (Lindegren et al. 2016). In our case,

the CI Cyg, AR PAV, V381 Sco, and V1329 Cyg sys-

tems have periods exceeding this threshold. Considering

this and their large distances, we exclude them from our

subsequent analysis.

2.2. Luminosities

Along with astrometry, Gaia DR2 provides stellar lu-

minosities (Andrae et al. 2018). The luminosities are

inferred via the FLAME module, which is part of the

Apsis data processing pipeline (Bailer-Jones et al. 2013).

As the authors specify, there are two potential sources

of systematic errors: the adopted bolometric correction

BCG,� and extinction. The former is estimated to be

+0.06 mag. The latter is assumed to be zero when cal-

culating the absolute magnitude, therefore resulting in

underestimated luminosity values (Andrae et al. 2018).

As we did for the distances in the previous section,

we test for consistency of the luminosities with previous

measurements (Figure 3). The comparison is performed

between DR2 luminosities and those inferred from B and

V magnitudes from the literature in SIMBAD. Most of

10 1 100 101 102

GAIA DR2, Parallax [mas]

10 1

100

101

102

GAIA

DR1

and

HIP

, Par

alla

x [m

as] 1:1 relation

Figure 2. Comparison of parallaxes reported in Gaia DR2(and used in our analysis; Table 3) with those in Gaia DR1 orHipparcos, where available. The two datasets are consistentto within the published measurement uncertainties and donot show any large systematics.

the SIMBAD values belong to the Tycho-2 catalog of

the 2.5 million brightest stars (Høg et al. 2000), whereas

other entries are taken from the Fourth US Naval Obser-

vatory CCD Astrograph Catalog (Zacharias et al. 2013)

and the fourth RAVE Data Release (Munari et al. 2014).

Both the B and V magnitudes are used to determine

the bolometric correction with the following empirical

calibration obtained from the catalog of nearby (< 8

pc) stars (Reid et al. 1995):

BC =

(B − V ) < 1.2 :

−0.121112 + 0.634846(B − V )− 1.01318(B − V )2+

0.125024(B − V )3;

(B − V ) > 1.0 :

−43.9614 + 115.958(B − V )− 110.511(B − V )2+

44.7847(B − V )3 − 6.74903(B − V )4.

(1)

The V magnitude is converted from apparent to abso-

lute scale using the distances retrieved from Gaia DR2.

We then convert the absolute magnitude (MV ) into

bolometric luminosity with

Lbol = 2.512−(MV −4.83)+(BC−BC�), (2)

where BC� = −0.076.

Considering Figure 3, the RMS of the residuals with

respect to the 1-to-1 line is 10.4 L�, which is roughly

consistent with the mean of the plotted SIMBAD

(Wenger et al. 2000) measurement errors (7.3 L�) and

> 2 times larger than that of the Gaia DR2 errors (4.1

L�). There are no obvious systematics. An additional

source of uncertainty in the luminosity would arise if the

measurements were made during the eclipse. However,

Detecting Exoplanets with Eclipsing Binaries 5

10 1 100 101 102 103

GAIA DR2, Luminosity [L ]10 1

100

101

102

103

Liter

atur

e Lu

min

osity

[L]

1:1 relation

Figure 3. Comparison of bolometric luminosities reportedin Gaia DR2, and listed in Table 3, with those inferred fromB and V magnitudes in SIMBAD. The two datasets are con-sistent within the measurement uncertainties, and there isno evident sign of systematics.

the consistency of the two datasets here suggests that

this possibility is unlikely.

3. DETECTING SELF-LUMINOUS PLANETS

To quantify the advantage of using an eclipsing binary

to directly image an orbiting SL planet requires that we

estimate the binary’s age and assume that the binary

and SL planet formed at the same time. We can then

model the fading of the SL planet as it cools (e.g. Mar-

ley et al. 2007; Mordasini 2013; Mordasini et al. 2017)

and ask if it is currently bright enough to be detected

(Fortney et al. 2010) by existing or near-future instru-

ments. In the following discussion, we constrain the ages

of our binaries using Algol models (van Rensbergen et

al. 2010; Mennekens & Vanbeveren 2017) in Section 3.1.

Some with large total masses are likely young enough to

host detectable SL planets. We then select two with the

youngest age ranges as potential targets for SL planetdirect imaging follow-up in Section 3.2.

3.1. Ages

Neither the eclipsing binary catalogs from which we

construct our sample nor the Gaia DR2 report ages,

which are expected with the third release of Gaia. De-

termining the ages of our binaries is challenging regard-

less; as discussed earlier, our Dmag > 2.5 criterion tends

to select classical Algols. Late B-dwarf classical Algols

are field blue stragglers, and, because they are rejuve-

nated by the mass transfer, are older than would be

expected from the main sequence lifetime of the current

primary (Paczynski 1971; Giannuzzi 1984; Iben & Tu-

tukov 1987). The total mass of the binary system is a

better proxy for age.

Mass transfer is nearly conservative for binaries with

initially A- or late-B type primaries that interact via

Case A Roche-lobe overflow (van Rensbergen et al. 2010;

Mennekens & Vanbeveren 2017). Under these circum-

stances, the total mass of the current system (Mtot) is

assumed equal to the sum of the components’ masses

when they reached the zero-age main sequence. Then,

the initial mass of the primary is

Mprimary =Mtot

q + 1, (3)

where q is the initial mass ratio. If the components are

initially the same, q = 1 and the minimum Mprimary is

0.5Mtot. For stable mass transfer to occur with a late-B

primary, q ≥ 0.4 and the maximum Mprimary is 0.7Mtot.

Given that the initial primary has now evolved into a

subgiant or giant, and that the time it spent on the

main sequence was much longer, its main sequence life-

time provides an estimate of the binary’s age. Assum-

ing 0.5Mtot ≤ Mprimary ≤ 0.7Mtot, with Mtot obtained

from the component masses in Budding et al. (2004),

and the relationship between stellar mass and main se-

quence lifetime, we convert this Mprimary range into an

age range for each of our binaries. This range could

extend to younger ages if mass loss affects the transfer

process (Mprimary > 0.7Mtot), but we proceed with the

conservative (older) age estimates above.

At present, there are few detections of SL planets

around binaries with massive stars. Observational stud-

ies such as the on-going BEAST survey (Janson et al.

2019) aim to address the question of exoplanets around

B-type stars. Our binaries have the potential to shed

more light on whether the incidence of massive planets

increases or declines with host stellar mass, thus con-

straining the stellar mass interval within which planet

formation is favorable (Janson et al. 2011, 2019).

3.2. Best targets

We shortlist the best eclipsing binaries for observa-

tional follow-up according to the criteria discussed pre-

viously: 1) depth of primary eclipse larger than 2.5 mag,

2) accessibility with current or near-future technology,

and 3) likelihood that the binary is young, and thus

that any orbiting SL planet is luminous, based on the

binary total mass. The two youngest eclipsing binaries

satisfying these criteria are U Cep and AC Sct, with age

intervals of 215-525 Myr and 640-1365 Myr, respectively

(Table 1).

An additional consideration for selecting suitable tar-

gets is whether the binary system has a tertiary com-

panion. Tokovinin et al. (2006) show that short period

(P ) binaries are more likely to have a third component

which, if orbiting at small separation, could suppress

planet formation (Moe & Kratter 2019). In particu-

lar, no circumbinary planet has been found around a

6 Bellotti et al.

Table 1. Best Eclipsing Binary Targets for Direct Detection of Self-Luminous Exoplanets

Name mV DmagV mJ DmagJ d Period tDmag abin M1 M2 Spec Type Age Interval

[pc] [days] [min] [AU] [M�] [M�] [Myr]

(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13)

U Cep 6.9 2.54 6.47 1.10 198.6 2.5 90 0.07 4.20 2.30 B7V+[G8III-IV] 215-525

AC Sct 10.0 2.60 9.72 1.62 985.6 4.8 168 0.11 2.80 1.40 B9+[G0IV] 640-1365

Note—(1) General Catalog of Variable Stars designation, (2) Magnitude at maximum brightness in V -band, (3) Depth ofprimary minimum in V -band, (4) Magnitude at maximum brightness in J-band, (5) Inferred depth of primary minimum inJ-band, (6) Distance inferred from parallax, (7) Binary period, (8) Duration of totality in primary eclipse, (9) Projectedseparation between binary components, (10) Mass of primary component, (11) Mass of secondary component, (12) Spectraltype, (13) Binary age estimated from total binary mass Mtot = M1 + M2. Columns (2), (3), (10), (11), and (12) are extractedfrom Budding et al. (2004), column (9) from Brancewicz, & Dworak (1980), column (4) from 2MASS Catalog (Cutri et al.2003), column (5) from Eq. 4, column (6) from Gaia DR2, columns (7) and (8) from Avvakumova et al. (2013). The values incolumn (13) are inferred from Mtot and Eq. 3, assuming q = 0.4 to 1. The evolutionary class of both targets is SA, i.e., classicalAlgols (Avvakumova et al. 2013).

P < 7 days binary, which may be due to the presence

of a tertiary (Hamers et al. 2016). In our case, U Cep

(P = 2.5 days) is known to have a third companion

(Tokovinin 2018) at ∼2800 AU, so its influence on a

potential planet’s dynamical stability is negligible. It is

not known whether AC Sct (P = 4.8 days) has a tertiary

component.

The projected separation between the binary stellar

components is 0.07 AU for U Cep and 0.11 AU for AC

Sct (Brancewicz, & Dworak 1980). Therefore, we would

expect any potential planet to lie on a P-type orbit. Be-

cause the eccentricities of the systems are not available,

we cannot assess the long-term stability of the planetary

orbits (Holman & Wiegert 1999; Quarles, & Lissauer

2016) at this time.

We add two more potential targets, V621 Cen and

RW Mon, if we relax our deep eclipse criterion from

Dmag > 2.5 to > 2.0 mag, a contrast improvement of

∼6×. The total mass of V621 Cen corresponds to an age

range of 196-478 Myr, i.e., comparable to U Cep’s, but

detecting SL planets around V621 Cen would require

future planned facilities due to its large distance (1.8

kpc). RW Mon is closer (505 pc) and has an age range

of 880-2150 Myr. RW Mon’s period variations may arise

from a close tertiary companion (Soydugan et al. 2011),

so targeting this system would not only test our direct

imaging method, but also reveal the nature of any third

component. There could be other good targets within

our sample, but some binaries have missing or conflict-

ing data, e.g., Dmag, which prevents us from evaluating

them.

For U Cep and AC Sct, we calculate whether the

eclipse would increase the infrared SL planet-star con-

trast to that required by the instruments. We consider

planet masses between 0.5 and 10MJ , determining the

corresponding planet J-band (1.25 µm) magnitude at

the age of the binaries with the Sonora evolutionary

models (Marley et al. 2020, in prep.). We then esti-

mate the planet-star contrast at eclipse and compare it

with the technology regions as in Figure 1.

The planet-star contrast during the eclipse is the ratio

of the planet luminosity to the primary minimum. The

eclipsing binary catalogs report the Dmag value for U

Cep and AC Sct in the V -band. Thus, we estimate

Dmag in J-band based on the V −J color of both binary

components1. Formally, we write the fluxes normalized

to the total flux in the V -band, i.e., Fp,V + Fs,V = 1,

where Fp,V and Fs,V = 10−0.4DmagV are the primary and

secondary V -band fluxes, respectively. Then, we have

DmagJ = −2.5 log

(Fs,J

Ftot,J

)(4)

where Fs,J = Fs,V 100.4(V−J)s and Ftot,J = Fp,J + Fs,J ,

where Fp,J = Fp,V 100.4(V−J)p .

The J-band luminosity during the eclipse is then cal-

culated as

LJ,e = 10−0.4DmagJ LJ (5)

where LJ and LJ,e are the J-band luminosities at max-

imum brightness and primary minimum, respectively.

Figure 4 illustrates the planet-star contrast versus sep-

aration plane for U Cep and AC Sct. Around U Cep,

planets of & 4.5MJ reach contrast levels & 10−7 and

are thus detectable with current ground- or near-future

space-based instruments. Planets of 3-4MJ and roughly

1.5-2.5MJ achieve contrast levels associated with future

ground- (∼10−8) and space-based (∼10−9) facilities, re-

1 The V − J colors are retrieved fromhttp://www.pas.rochester.edu/ emamajek/spt/. For U Cep’sprimary and secondary, V − J is -0.24 and 1.57, respectively; forAC Sct, this color is -0.09 and 1.06.

Detecting Exoplanets with Eclipsing Binaries 7

-2 -1 0 1log Planet-Star Separation [arcsec]

-10

-9

-8

-7

-6

-5

-4

-3lo

g Pl

anet

-Sta

r Con

trast

1.6

2.1

2.63.13.74.24.75.26.37.38.49.410.5 MJ

3.1

3.74.24.75.2

6.37.38.49.410.5 MJ

U Cep (215 Myr)AC Sct (640 Myr)

Figure 4. Planet-star J-band contrast versus separationfor simulated SL planets around U Cep (red) and AC Sct(purple). The planet-star contrast is the ratio of the planetto binary luminosity at eclipse. The numerator is derivedfrom the Sonora evolutionary models (Marley et al. 2020,in prep.), assuming that the planet is as old as the esti-mated binary age. The denominator is equivalent to theJ-band Dmag, which is estimated from the V − J colors ofthe binary components (see Eq. 4). Current and future ob-servational facilities are denoted by the colored regions asin Figure 1. Around U Cep, current ground- or near-futurespace-based (yellow), future ground-based (green), and fu-ture space-based (blue) instruments can detect & 5MJ , 3-4MJ , and 1.5-2.5MJ SL planets, respectively. Around ACSct, an older binary, these limits rise to & 9MJ , 6-8MJ , and3-5MJ , respectively.

spectively. Around AC Sct, &9MJ planets can be de-

tected with current ground- or near-future space-based

instruments, while 6-8MJ and 3-5MJ planets require

future ground- and space-based observatories, respec-

tively.

In Figure 4, the technology regions are the same as

in Figure 1 and defined by the instrument performance

and inner-working angle. The location of the simulated

planets on the planet-star separation axis is arbitrary.

Current ground- or near-future space-based instruments

are characterized by observable separations of roughly

0.1-2 arcsec. At the distances of U Cep and AC Sct,

this range translates into orbital semi-major axes of 20-

400 AU and 100-2000 AU, respectively, consistent with

those of known directly imaged planets around binaries

(Schwarz et al. 2016).

Are the SL planets in Figure 4 bright enough to be

detectable? Figure 5 shows the change of the direct

imaging detection limits as a function of exposure time

and planet mass. For U Cep, in a <1 hour exposure,

& 4.5MJ planets are detectable with current ground-

or near-future space-based instruments, ∼1.5MJ planets

0 2 4 6 8 10 12 14 16Exposure time [hours]

24

26

28

30

32

34

Appa

rent

J M

agni

tude

1.6MJ

2.1MJ

2.6MJ

3.1MJ

3.7MJ

4.2MJ

4.2MJ

4.7MJ

5.2MJ

6.3MJ

7.3MJ

8.4MJ

9.4MJ

10.5MJ

8-m ground-based telescope30-m ground-based telescope6.5-m space-based telescope

U Cep low-limit (215 Myr)AC Sct low-limit (640 Myr)

Figure 5. Comparison of simulated SL planet J-band mag-nitude with observational detection limits, as a function ofexposure time and planet mass. We consider different obser-vational capabilities corresponding to some of the technologyregions in Figure 4: a current 8-m (yellow dashed line) and afuture 30-m (green dot-dashed) ground-based telescope anda near-future 6.5-m (yellow dashed) space facility. We as-sume photon-limited detections and require a signal-to-noiseratio of 10. For the ground-based observatories, we adoptan overall efficiency of 20% and a sky background contam-ination of 16.7 mag/arcsec2. For the space-based case, weassume a 30% efficiency and a zodiacal background contam-ination of 22.2 mag/arcsec2. SL planet magnitudes are cal-culated at the distances of U Cep (red) and AC Sct (purple).Around U Cep, within a 1 hour exposure, a & 4.5MJ planet isdetectable with current ground-based or near-future space-based instruments, a ∼1.5MJ planet with future ground-based facilities. Around AC Sct, these limits rise to ∼9MJ

and 6-8MJ , respectively.

with a future ground-based observatory. For AC Sct,

these limits are ∼9MJ and 6-8MJ , respectively.

The eclipse durations and short orbital periods of U

Cep and AC Sct allow for efficient scheduling of their ob-

servations. Ideally, the duration of totality would exceed

the exposure time required for planet detection and the

eclipses would be frequent enough to allow for repeat

observations during a typical observing run of several

nights. As discussed above, a ∼1 hour exposure is re-

quired to detect planets around U Cep and AC Sct with

current ground- or near-future space-based facilities. In

comparison, the totality of the primary eclipse for U Cep

lasts 90 minutes, during a primary eclipse of 9.0 hours

that occurs every 2.5 days. For AC Sct, totality lasts

approximately 168 minutes, during a primary eclipse of

16.1 hours every 4.8 days. Therefore, we can achieve the

detection limit within a reasonable time frame.

In Figure 6, we compare the properties of U Cep and

AC Sct (Table 1) with those of binaries known to host

8 Bellotti et al.

-2 -1 0 1 2 3 4log Binary semi-major axis [AU]

-1

-0.7

-0.3

0

0.3

0.7

1lo

g Bi

nary

tota

l mas

s [M

] Janson et al. 2019SPOTS

HIP 77911

Known host binaries Our eclipsing binaries

M

KGF

A

B

O

Figure 6. Binary total mass versus semi-major axis planecomparing our two best eclipsing binary targets, U Cep andAC Sct, with binaries known to host planets (Schwarz et al.2016). The datapoints are color-coded by primary spectraltype; black dots are systems whose primary type is not avail-able. As a consequence of our Dmag and binary age selectioncriteria, U Cep and AC Sct have earlier type primaries andlarger total masses than typical for known host binaries. Thisdistinction would expand the parameter space of SL hostingbinaries, were planets discovered in our target systems. Wehighlight HIP 79098 AB (gray), which has a candidate sub-stellar companion (Janson et al. 2019) and HIP 77911 (pink),which has a candidate planetary-mass companion (Asensio-Torres et al. 2018); these systems have primaries with similarspectral types to our targets, but lower total masses. In com-parison to the parameter boundaries of the SPOTS survey(Asensio-Torres et al. 2018) (pink box), our targets fall onthe low side of the semi-major axis distribution and the highend of total mass.

exoplanets (Schwarz et al. 2016), regardless of the detec-

tion technique. Our eclipsing binary targets have earlier

type primaries and larger total masses due to the Dmag

and age criteria applied above. Specifically, among the

97 binaries known to host planets and whose primaries

have been spectrally classified (Schwarz et al. 2016),

there are none with B- or A-type primaries or with to-

tal masses above 4M� like our targets. Thus, targeting

our binaries would expand the spectral and total mass

ranges of host binaries, were planets discovered there.

4. SEARCHING FOR REFLECTED LIGHT

PLANETS

Of the few direct imaging detections around binaries

so far, none has included a RL planet. Therefore, it

is interesting to test whether our approach would give

access to this unexplored territory. There is an added

benefit to using eclipsing binaries for RL planets rel-

ative to SL ones. At eclipse, the planet-star contrast

is improved, while the RL planet is still brightened by

the light of both binary components. RL planets have

evolved and cooled, so, unlike for SL planets, there is no

age constraint for shortlisting targets here.

In selecting potential RL targets, the main observa-

tional limitation is the distance of the binary from us.

Given that a typical inner working angle of current coro-

nagraphs (yellow region; Figure 7) is on the order of 100

mas, we do not expect to observe RL planets on tight

orbits in faraway systems. At 50 pc away, the separa-

tion corresponding to a planet on a 1 AU orbit is only 20

mas, whereas an orbit of 5 AU or larger is observable.

Therefore, we consider only the three deep (Dmag &2.5) eclipsing binaries within roughly 50 pc: V1412 Aql

at 22.9 pc, RR Cae at 21.2 pc, and RT Pic at 54.4 pc

(Table 2). In the case of RR Cae, a 4.2 MJ planet has

already been discovered via the eclipse timing method

at 5 AU (Qian et al. 2012b).

For RR Cae (11 min totality, ∼14 min eclipse dura-

tion, 7.2 hour orbital period), there is little time during

totality for observations, but we can build up a long ex-

posure by observing for a short fraction of each night

over multiple nights. For the other two targets, the to-

tality, eclipse duration, and binary period are not known

at present.

To determine whether an RL planet is detectable for

these binaries, we build the observational space in Fig-

ure 7 as in Figure 1. We assume that Jupiter- and

Venus-like planets (i.e., with the same sizes and albe-

dos as the originals) have observed planet-star separa-

tions of at least 20 mas, the detectable angular sepa-

ration limit of future technology (green and dark blue

regions in the figure). We implicitly assume that these

simulated planets could lie this physically close to their

binary. The separation between the binary components

is not available in the literature; therefore, we cannot

predict whether the planets lie on a P- or S-type orbit.

For each binary, we derive the planet-star contrast atmaximum binary brightness (Lpl/LV ) and during the

eclipse (Lpl/LV,e) in the V -band. The luminosity at

eclipse (LV,e) is estimated from Eq. 5 using the V -band

luminosity outside eclipse (LV ) and the observed Dmag,

which is reported in other photometric filters for RR Cae

(B) and RT Pic (p) (see Table 2). Unlike in Section 3.2,

we do not have enough information here about the pri-

mary and secondary (i.e., spectral type and/or color),

and so cannot convert Dmag to a common band. Given

that Dmag may decrease with increasing wavelength,

the plotted contrast boost could be overestimated for

RR Cae and RT Pic.

The luminosity of the planet due to reflection (Lpl) is

obtained from the definition of the albedo. Assuming

the planet to be a disk of radius Rpl and manipulating

Detecting Exoplanets with Eclipsing Binaries 9

Table 2. Best Eclipsing Binary Targets for Direct Detection of Reflected Light Exoplanets

Name m Dmag Filter mV d Period tDmag Spec Type

[pc] [days] [min]

(1) (2) (3) (4) (5) (6) (7) (8) (9)

RR Cae 14.88 3.30 B 14.40 21.21 0.3 11 WD+M5-6V

V1412 Aql 15.67 2.63 V 15.67 22.93 · · · · · · DC7

RT Pic 9.90 2.60 p 9.07 54.39 · · · · · · G8V

Note—(1) General Catalog of Variable Stars designation, (2) Magnitude at maximum brightness, (3) Depth of primary mini-mum, (4) Filter band for m and Dmag, (5) Magnitude at maximum brightness in V -band, (6) Distance, (7) Binary period, (8)Duration of totality in primary eclipse, (9) Spectral type. Columns (2), (3), (4), (7), (8), and (9) are extracted from Avvakumovaet al. (2013), column (6) is inferred from Gaia DR2 parallax, and column (5) is taken from SIMBAD. RR Cae’s evolutionaryclass is DW, i.e., white dwarf system Avvakumova et al. (2013).

the ratio between the incident and reflected flux yields

Lpl = Ltot

R2pl

a2

α

4, (6)

where a and α are the semi-major axis of the planetary

orbit and the albedo, respectively. The values of a, α,

and Rpl are taken from the NASA fact sheets for each

planet2.

Figure 7 shows that for Jupiter- and Venus-like plan-

ets with planet-star separations within the technology

limits, the planet-star contrasts also will be accessible.

To further quantify the detectability of these simu-

lated RL planets, we estimate the required exposure

times for future ground- (Figure 8) and space-based

(Figure 9) observations, similarly to Figure 4. For

ground-based and space-based imaging, respectively, we

assume that the observations are carried out in the J-

(1.25 µm) and V -band (555 nm) to achieve the best

planet-star contrast and separation. For the ground-

based case, a large telescope aperture (∼30 m) with

extreme-AO in the near-IR leads to the best contrast,

given that extreme-AO in visible light is considerably

more challenging.

For a Jupiter-like planet at 20 mas separation, the

J-band magnitude is brighter than 33 mag around all

three eclipsing binaries, implying a ground-based detec-

tion within 2 hours (Figure 8). With a space-based tele-

scope (Figure 9), a Jupiter around RT Pic would be

detected in less than 1 hour with a 2-m aperture, in

∼4 hours around RR Cae with a 3-m aperture, and in

∼7 hours around V1412 Aql with a 5-m aperture. For

a Venus-like planet at 20 mas separation, the J-band

magnitude would be 37.1, 32.6, and 31.7 mag, respec-

tively, for V1412 Aql, RR Cae, and RT Pic, requiring at

most two hours to reach the detection limit around RR

2 For Jupiter, Venus, and Earth: /jupiterfact.html, /venus-fact.html, and /earthfact.html, respectively.

Cae and RT Pic with a ground-based telescope. With

a space-based telescope, the V -band detection limit of

a Venus would be accessible in less than 10 hours, even

with a 4-m aperture, but only around RT Pic.

To Figure 7, we add an Earth-like (same radius and

albedo) planet, assuming that it is habitable, i.e., that

the bolometric flux it receives from the binary is equal

to the solar constant. Therefore, the semi-major axis

of the planet is obtained from√

Ltot in AU. We do not

consider the eccentricity of the binary or the gravita-

tional interaction as the secondary moves in its orbit,

both of which alter the habitable zone boundaries over

time (Haghighipour, & Kaltenegger 2013; Kaltenegger,

& Haghighipour 2013; Jaime et al. 2014).

While the “Earth” in RT Pic lies at a planet-star con-

trast detectable by space-based telescopes planned for

the 2030’s, its separation from the binary is beyond the

capability of those instruments. Indeed, the planet-star

separation would be 13 mas, the planet-star V -band

contrast 1.5×10−9 during the eclipse, and the planet J-

band magnitude 32.4 and V -band magnitude 33.7. The

detection limit in the J-band would be achievable in ∼2

hours with a ground-based telescope and in the V -band

in <7 hours with a 6- to 8-m space-based telescope. For

the ground-based case, the 1.5 × 10−9 planet-star con-

trast is not achievable. For the space-based telescopes,

the diffraction limits are 23.3 mas (λ = 555 nm, D = 6

m) and 17.5 mas (λ = 555 nm, D = 8 m). Detection

at 13 mas would require that the 1.5× 10−9 planet-star

contrast be achieved at 0.56 λ/D and 0.74 λ/D, respec-

tively. Coronagraphs currently deliver deep contrast

levels at &2 λ/D (Guyon et al. 2006), so detecting an

Earth-like planet around an eclipsing binary will require

larger apertures and/or further advances in coronagraph

technologies.

10 Bellotti et al.

-9

-7

-5

-3

-1

V1412 Aql

Jupiters no eclipseJupiters with eclipse

Venuses no eclipseVenuses with eclipse

Earths no eclipseEarths with eclipse

-9

-7

-5

-3

log

Plan

et-S

tar C

ontra

st

RR Cae

-2 -1 0 1log Planet-Star Separation [arcsec]

-9

-7

-5

-3

RT Pic

Figure 7. Testing the detectability of RL planets during the eclipse of the binary. Once again, we plot the observational spacefrom Figure 1, but now for the closest eclipsing binaries in our sample. We consider simulated Jupiter- (gray) and Venus-like(orange) planets at increasing separations from the hosts, starting from 20 mas, the detectable angular separation limit of near-future facilities (green and dark blue regions). We plot the V -band contrasts both during (dotted lines) and before/after theeclipse (solid lines). The planet flux observed would decline from left to right. A habitable (i.e., flux received from the binaryequal to the solar constant) Earth-like (same radius and albedo) planet only appears on this plot for RT Pic, at eclipse (greencircle) and at other times (green triangle). The planet-star contrasts and separations of “Jupiters” and “Venuses” orbiting ourclosest binaries are accessible with future ground- and space-based coronagraphs. While the contrast of the habitable Earth-likeplanet around RT Pic is achievable with planned space-based instruments, the planet-star separation is too small.

Detecting Exoplanets with Eclipsing Binaries 11

0 2 4 6 8 10 12 14Exposure time [hours]

22

24

26

28

30

32

34

36

Appa

rent

J M

agni

tude

Detection limit in J bandRL Jupiter in RR CaeRL Venus in RR Cae

RL Jupiter in RT PicRL Venus in RT Pic

RL Jupiter in V1412 AqlRL Venus in V1412 Aql

Figure 8. Comparison of simulated J-band magnitude withobservational detection limit, as a function of exposure time.We assume a 30-m ground-based telescope with the same pa-rameters as in Figure 5 (green dot-dashed line). Also plottedare the apparent magnitudes of a “Jupiter” and “Venus” ata separation of 20 mas from V1412 Aql (dotted gray anddotted orange lines, respectively), RR Cae (dashed gray anddashed orange lines), and RT Pic (solid gray and solid or-ange lines). RT Pic is the most promising target given thehigher brightness of the planets there. A Jupiter-like planetaround all three binaries and a Venus-like planet around RRCae and RT Pic would be detected in two hours.

5. ADVANTAGES OF OUR METHOD AND

TARGETS

As discussed in the previous section, for reflected light

planets, the eclipse improves the observed planet-star

contrast by dimming the primary while the planet re-

mains illuminated by both stars. Observing eclipsing

binaries has several other advantages for directly imag-

ing both SL and RL planets: 1) The reduction of the

binary to a point-like source during eclipse makes coron-

agraphy feasible. 2) The increase in planet-star contrast

during eclipse makes fainter planets accessible. 3) The

contrast boost allows detection of planets in intrinsically

brighter, and thus more massive, stellar systems.

Eclipsing binaries are observable with a coronagraph,

which would block the point-like light of the superim-

posed stellar components during the eclipse. In this way,

our method incorporates both a natural starshade (the

eclipse) and a man-made coronagraphic measurement.

Light leakage during the coronagraphic measurement

would be minimized, because it is possible to predict

the time and duration of the binary eclipse accurately.

In fact, one of our eclipsing binaries, RR Cae, has been

successfully targeted in an exoplanet search with the

eclipse timing variation technique (Qian et al. 2012b),

for which timing accuracies on the order of seconds are

required (Sybilski et al. 2010)

Our method would work best during the total eclipse

of the primary, when the binary is point-like and darkest.

0 2 4 6 8 10 12 14Exposure Time [hours]

26

28

30

32

34

36

38

Appa

rent

V M

agni

tude

2-m3-m4-m5-m6-m7-m8-m

Detection limit in spaceRL Jupiter in RR CaeRL Venus in RR Cae

RL Jupiter in RT PicRL Venus in RT Pic

RL Jupiter in V1412 AqlRL Venus in V1412 Aql

Figure 9. Comparison of simulated V -band magnitude withobservational detection limit, as a function of exposure time,for the Venus- and Jupiter-like planets in Figure 8, assumingfuture space-based telescopes of different apertures and sameparameters as the space-based case in Figure 5 (blue solidlines). The Venus-like planet (orange) around RT Pic wouldbe detected in the V -band in less than 10 hours with a 4-mtelescope. A Jupiter-like planet (gray) around RT Pic wouldbe observable in less than 1 hour with a 2-m aperture, in ∼4hours around RR Cae with a 3-m aperture, and in ∼7 hoursaround V1412 Aql with a 5-m aperture.

However, we could still observe at partial eclipse, pro-

vided that the angular separation between the binary

components is below the angular stellar size tolerance

of the coronagraph (on the order of 0.1 λ/D). For our

best SL targets (U Cep and AC Sct), the projected bi-

nary separation is ∼0.01 λ/D in the J-band with a 30-m

telescope. Thus, observations during partial eclipse are

potentially useful. For our best RL targets (V1412 Aql,

RR Cae, and RT Pic), we lack sufficient information to

infer the binary separations and make a similar evalua-

tion.

The performance of a coronagraph with a small IWA

(<3 λ/D) degrades with increasing stellar angular size

(Guyon et al. 2006), even if the source size is well below

the diffraction limit; light leakage is significantly higher

for a partially resolved stellar disk than it would be for

an on-axis point source. A shaped pupil could be used

instead, but the IWA would be greater than ∼3 λ/D,

restricting the target distance at which observations of

close-in orbits could be made.

How might our method expand the parameter space

of detected exoplanet properties? We have considered

binaries that dim by at least 2.5 mag during eclipse.

As a result, at a given planet-star contrast and host

intrinsic stellar luminosity, our approach allows direct

imaging detection of a planet at least 10× fainter than

one around a single host star or non-eclipsing binary

system, at the same distance. Physically, for RL planets,

12 Bellotti et al.

10× fainter could imply 0.1× planet albedo or 0.33×planet size or 3× planet-star separation (see Eq. 6).

Does our method allow access to different types of host

stars or binaries than previously explored in exoplanet

systems? The Catalogue of Exoplanets in Binary Sys-

tems (Schwarz et al. 2016) lists detailed properties for

97 binaries with exoplanets, at least five of which were

directly imaged. Compared to these, our best target bi-

naries for SL planets are more massive and have earlier-

type primaries (i.e., B-type). This difference arises be-

cause we selected on higher total binary mass so that any

SL planets would be young and therefore bright enough

to be detectable.

In general, compared to single star and non-eclipsing

binary systems at the same distance, planet luminosity,

and planet-star contrast, our SL and RL target binaries

would have stellar luminosities at least 10× brighter and

thus stellar masses ∼2-3× greater, assuming the canon-

ical main sequence mass-luminosity relation. Therefore,

our approach could expand the boundaries of the host

stellar mass parameter space. If more exoplanets are

discovered around such massive primaries, there are im-

plications for how such planets form and evolve around

binaries (Janson et al. 2011, 2019).

6. CONCLUSIONS

We investigate the plausibility of a new approach for

directly imaging exoplanets. Our idea is to use the

eclipse event in eclipsing binary systems as a tool to

boost the planet to star flux contrast, i.e., to exploit

a natural starshade. During the eclipse, the binary is

reduced to a point-like source, making coronagraphic

observations possible.

We select 289 binaries where the depth of primary

minimum Dmag is > 2.5 mag, which boosts the planet-

star contrast by more than a dex. Thus, at a given

observed planet-star contrast and host intrinsic stellar

luminosity, we can detect a planet > 10× fainter in an

eclipsing binary, during eclipse, than in other star sys-

tems at the same distance. Likewise, we can detect plan-

ets of a given intrinsic luminosity around systems whose

intrinsic stellar luminosity is > 10× brighter and whose

stellar mass is ∼ 2-3× greater. In other words, we could

directly image exoplanets in a massive binary system at

the same contrast level as in a lower mass one.

We consider using this method to detect self-luminous

(SL) and reflected light (RL) planets around our bina-

ries.

For the SL planet case, we determine whether 0.5 to 10

MJ planets could be detected during eclipse with current

or future coronagraphs. SL planets are easiest to detect

in young systems, given that the thermal emission of the

planet decreases with time. Therefore, we select on the

age of the binary, as well as the infrared brightness of

the planet at the distance of the binary and the planet-

star contrast. Because we lack measured ages for our

best target binaries, which are classical Algols, we use

Algol models (van Rensbergen et al. 2010; Mennekens &

Vanbeveren 2017) to constrain a plausible age interval

from the binary total mass.

Using these criteria, we identify two targets: U Cep

and AC Sct. Around them, we might detect ∼4.5MJ

and ∼9MJ SL planets with current ground- or near-

future space-based instruments, respectively. With fu-

ture ground-based facilities, these limits reduce to 3-

4MJ and 6-8MJ . Because of our Dmag > 2.5 and

age criteria, these targets possess larger total masses

(> 4.2MJ) and earlier-type (B7-, B9-) primaries than

typical of known host binaries. As noted above, our

method puts such massive stellar systems within the

reach of coronagraphic observations; targeting these sys-

tems would expand the host stellar parameter space for

testing SL planet formation and evolution.

For RL planets, the advantage of using eclipsing bina-

ries is that both binary stars continue to illuminate the

planet while the planet-star contrast is increased dur-

ing the eclipse. To find the best targets in this case,

we focus on only the nearest (within ∼50 pc) eclipsing

binaries in our sample: RR Cae, V1412 Aql, and RT

Pic, for which the contrast boost during eclipse is 1.32,

1.05, and 1.04 dex, respectively. We assume that a large

(30-meter) ground-based telescope and intermediate (2-

to 8-meter) space-based telescopes will be available in

the future.

We consider a Jupiter-like, Venus-like, and habitable

Earth-like planet, estimating the change in detection

limit with exposure time in the J-band with future

ground-based telescope and in the V -band with spacetelescopes. For a “Jupiter” at 20 mas in all three target

binaries, a detection is achieved in less than ∼10 hours

with the ground- and space-based telescopes, whereas

for a “Venus” at 20 mas, detection is possible in the

J-band around RR Cae and RT Pic and in the V -

band around RT Pic. Thus, directly imaging these

Jupiter- and Venus-like planets is within the capabili-

ties of planned facilities.

Detection of a habitable Earth-like planet remains a

challenge. In a less than ∼10 hour exposure in the J-

band with the ground-based telescope or in the V -band

with a 6- to 8-meter space telescope during the eclipse,

this “Earth” would be bright enough to detect if it or-

bited RT Pic. The planet-star contrast of 1.5 × 10−9

would be achievable from space. The planet-star sep-

aration of 13 mas is equivalent to 0.56 λ/D and 0.74

Detecting Exoplanets with Eclipsing Binaries 13

λ/D for the future 6-m and 8-m space-based telescopes,

respectively. Given that current coronagraphs deliver

deep contrast levels at &2 λ/D (Guyon et al. 2006),

larger apertures and/or new coronagraph advances will

be required for detection.

This work has made use of data from the Euro-

pean Space Agency (ESA) mission Gaia (https://www.

cosmos.esa.int/gaia), processed by the Gaia Data Pro-

cessing and Analysis Consortium (DPAC, https://www.

cosmos.esa.int/web/gaia/dpac/consortium). Funding

for the DPAC has been provided by national institu-

tions, in particular the institutions participating in the

Gaia Multilateral Agreement. This research has made

use of the SIMBAD database, operated at CDS, Stras-

bourg, France.

We thank the anonymous referee for detailed com-

ments that substantially improved this paper. We are

deeply indebted to Maxwell Moe for sharing his exper-

tise on eclipsing binary evolution, the possibility of ter-

tiaries, and how to estimate the ages of Algols, Glenn

Schneider for important discussions on the feasibility of

this project, and Elizabeth Green, Andrew Odell, and

Joel Eaton for their knowledge about eclipsing binary

catalogs. Many thanks to Eric Mamajek, Andrej Prsa,

Guillermo Torres, Edwin Budding, and Oleg Malkov

for useful discussions on the discrepancies between the

values of primary eclipse depth, and to Jack Lissauer

for providing insight into the stability of planetary or-

bits around binary systems. AIZ thanks the Dark Cos-

mology Centre at the Niels Bohr Institute, University

of Copenhagen, for its hospitality and support, which

greatly furthered the development of this project.

APPENDIX

A. FULL SAMPLE OF ECLISPING BINARY SYSTEMS

14 Bellotti et al.

Table

3.

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

LQ

Cas

00

04

10.9

597

+61

42

07.8

75

14.1

03.1

0p

M0

0.6

25.2

1··

···

·V

0411

Cas

00

30

11.2

537

+56

07

47.7

34

13.9

02.8

0B

M0

0.5

011.6

4··

···

·C

VC

as

00

31

54.4

680

+71

41

38.6

38

13.9

03.6

0p

M0

0.8

68.2

3··

···

·G

WA

nd

00

35

09.4

069

+41

40

03.4

65

14.6

02.8

0B

M0

0.2

8··

···

···

·K

QC

as

00

38

37.6

964

+58

32

42.2

10

14.1

02.9

0p

M0

0.3

120.3

6··

···

·U

UA

nd

00

43

45.0

803

+30

56

19.6

64

11.2

03.0

0V

B0

1.0

610.3

5F

5+

[K7IV

]SA

YH

yi

00

45

50.6

974

-78

49

16.8

12

10.4

03.6

0V

TB

12.3

515.3

8A

6V

+[K

3IV

]SA

VT

uc

00

51

46.5

986

-71

59

52.9

27

10.6

08.2

0V

B1

1.8

613.3

5B

9/A

2IV

+[G

5IV

]SA

V0415

Cas

00

54

31.0

203

+59

24

00.0

27

14.0

02.8

0b

M0

0.4

024.2

0··

···

·V

0386

Cas

00

59

11.2

667

+55

57

20.0

37

14.0

02.5

0b

M0

0.3

435.0

9··

···

·U

Cep

01

02

18.4

416

+81

52

32.0

80

6.8

62.5

4V

HB

05.0

4··

·B

7V

e+

B8II

I-IV

SA

FK

And

01

07

03.1

786

+37

29

04.8

25

14.0

02.8

0p

M0

0.3

0··

···

···

·O

YP

er

01

34

53.7

662

+53

39

31.6

06

15.0

02.5

0p

M0

0.3

39.2

8··

···

·B

LH

yi

01

41

00.3

995

-67

53

27.4

69

14.9

03.1

0V

M0

7.6

5··

·W

D+

M3-V

S2C

V0367

Cas

01

41

11.2

580

+61

15

35.2

26

14.2

03.1

0b

M0

0.4

812.5

3··

···

·G

HC

as

01

49

03.3

762

+56

16

26.8

17

12.3

03.5

0p

M0

0.8

716.7

2F

2··

·H

SP

er

01

52

11.6

089

+57

06

42.7

15

13.0

02.6

0V

M0

0.8

912.5

1A

0II

-III

···

XZ

And

01

56

51.5

242

+42

06

02.1

81

10.1

52.5

4V

M1

1.8

421.9

7A

1V

+G

5IV

SA

XT

ri02

00

33.7

370

+27

53

19.2

05

8.5

52.7

2V

B1

4.8

2··

·A

5V

+[G

2IV

]SA

LL

Per

02

14

51.5

373

+57

29

34.4

56

17.0

02.5

0p

M0

0.2

2··

···

···

·M

OA

nd

02

23

53.6

676

+39

59

02.7

67

13.2

03.0

0p

M0

0.5

215.6

5··

···

·L

MP

er

02

23

56.6

594

+56

17

19.3

10

16.0

02.5

0p

M0

0.2

1··

···

···

·R

WT

ri02

25

36.1

555

+28

05

50.8

93

12.5

02.4

2V

M1

3.1

7··

·W

D+

M0V

S2C

WW

Hor

02

36

11.4

500

-52

19

13.5

00

18.4

04.6

0B

M0

3.0

7··

·W

D+

M6V

S2C

ZP

er

02

40

03.2

329

+42

11

57.6

94

9.7

02.7

0p

M0

1.4

042.9

4A

0V

+G

0IV

-VSA

VFor

02

40

51.4

000

-25

52

14.0

00

14.0

52.9

5B

M0

0.3

6··

···

···

·V

0366

Per

02

44

10.6

082

+36

35

11.8

56

13.7

02.6

0b

M0

0.8

95.5

0··

···

·V

0479

Per

02

46

05.6

524

+43

17

07.6

79

14.4

03.5

0b

M0

0.1

2··

···

···

·SS

Cet

02

48

36.2

797

+01

48

26.8

94

9.4

03.6

0V

B1

1.8

525.4

A0+

[K0IV

]SA

LS

Per

02

57

08.7

937

+37

53

40.4

77

11.5

03.1

0p

M0

0.8

830.2

3A

7IV

···

QU

Per

03

05

54.6

745

+40

41

12.6

25

12.2

02.6

0p

M0

0.5

433.6

0A

2V

e··

·E

FE

ri03

14

13.2

527

-22

35

42.9

17

14.5

02.8

0B

M0

6.2

7··

·W

D+

M9V

S2C

UZ

For

03

35

28.6

515

-25

44

21.7

65

18.2

02.7

0V

M0

4.1

7··

·W

D+

M4.5

VS2C

WY

Per

03

38

24.1

599

+42

40

39.0

50

11.5

02.7

0V

B0

0.8

321.2

2A

0+

[K25IV

]SA

GH

Cam

03

41

31.4

326

+67

04

43.5

71

13.8

03.2

0B

M0

0.4

315.7

0··

···

·H

UP

er

03

42

33.1

687

+39

06

03.6

25

10.9

03.0

0p

M0

0.5

466.2

8A

···

Table3continued

Detecting Exoplanets with Eclipsing Binaries 15Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

RW

Tau

04

03

54.3

166

+28

07

33.4

98

8.0

03.8

0V

HB

13.2

7··

·B

8V

e+

K0IV

SA

RV

Per

04

10

37.9

975

+34

15

54.8

83

10.2

33.7

5V

B1

0.9

437.7

1A

2+

[G7IV

]SA

V0407

Tau

04

10

58.4

464

+26

18

21.6

20

13.5

03.0

0b

M0

0.4

0··

···

···

·B

NT

au

04

15

44.8

807

+30

40

08.0

25

10.1

74.2

1V

TB

10.7

316.4

3A

7+

[K2IV

]S

RR

Cae

04

21

05.5

631

-48

39

07.0

62

14.8

83.3

0B

M0

47.1

6··

·W

D+

M5-6

VD

W

TZ

Eri

04

21

40.3

306

-06

01

09.2

01

9.8

02.8

0V

B0

3.3

010.4

1F

3+

[K5IV

]SA

AO

Eri

04

32

00.9

396

-17

44

47.6

51

10.3

03.2

7V

TB

11.4

019.1

7A

8+

[G8IV

]S

AC

Tau

04

37

06.3

566

+01

41

31.1

58

10.3

02.8

0V

B0

1.7

710.4

4F

0+

[K6IV

]SA

AQ

Cam

04

51

17.8

661

+54

55

59.9

50

12.0

04.0

0p

M0

0.2

3··

·B

7··

·E

TO

ri04

55

34.3

335

+01

22

49.5

64

10.1

33.1

4V

TB

12.7

55.8

7G

3+

[K3IV

]SA

EQ

Ori

04

57

14.4

652

-03

36

04.6

57

10.2

03.1

0P

B0

1.6

5··

·A

0+

[G2IV

]SA

V0723

Tau

05

06

46.1

891

+28

35

23.4

10

14.5

02.6

0b

M0

0.6

04.3

0··

···

·F

LO

ri05

07

46.6

517

-02

44

38.2

47

11.4

03.2

0P

B0

1.9

85.9

1A

3V

+[K

0IV

]SA

XZ

Cam

05

17

12.6

742

+75

50

05.3

54

11.4

03.0

0V

TB

00.8

142.4

4A

0+

[K0.5

SA

V0666

Ori

05

54

06.5

077

+17

56

18.1

52

14.0

03.0

0b

M0

0.7

59.8

4··

···

·R

TP

ic06

00

41.2

987

-44

53

50.1

06

9.9

02.6

0p

M0

18.3

90.6

2G

8V

···

LP

Gem

06

05

05.1

282

+26

40

53.3

86

12.5

02.5

0p

M0

0.5

335.8

8··

···

·D

XO

ri06

05

44.7

518

+08

20

55.2

89

13.5

02.5

0p

M0

0.4

117.1

1··

···

·C

GP

up

06

07

02.9

478

-46

44

47.7

17

10.2

02.8

0p

M0

2.0

55.8

1··

···

·V

0347

Pup

06

10

33.6

585

-48

44

25.3

63

13.5

52.4

0B

M0

3.3

80.2

8··

···

·R

UC

Ma

06

22

33.5

874

-22

41

29.1

51

11.3

02.5

0V

M1

1.3

6··

···

···

·B

OG

em

06

25

01.2

975

+17

58

12.7

77

11.3

03.8

0P

B1

1.1

221.1

0A

2+

[K3IV

]SA

NS

CM

a06

29

09.5

867

-31

15

33.8

00

14.0

02.7

1V

M0

0.4

010.9

7··

···

·V

0639

Mon

06

37

20.1

591

+11

36

31.4

85

16.0

02.5

0b

M0

0.3

1··

···

···

·M

PG

em

06

48

33.4

000

+19

37

15.0

00

15.4

02.6

0p

M0

0.4

1··

···

···

·E

FG

em

06

51

01.8

159

+17

29

54.9

87

12.6

02.8

0p

M0

0.4

136.4

2··

···

·C

HM

on

06

51

38.1

214

+05

56

12.2

68

13.1

02.8

0p

M0

0.5

717.7

3··

···

·E

HM

on

06

52

08.7

658

-07

03

52.9

71

13.8

02.5

0p

M0

0.1

9··

···

···

·R

VL

yn

06

56

11.4

299

+50

51

45.5

09

13.0

02.6

0p

M0

0.6

612.3

8··

···

·B

PM

on

06

56

55.4

325

+05

01

43.8

20

13.1

02.6

0p

M0

0.7

411.0

6··

···

·F

SG

em

06

57

21.2

259

+16

30

13.6

81

13.7

02.9

0p

M0

0.2

5··

···

···

·U

UC

Ma

06

58

41.4

546

-18

48

48.1

87

10.0

02.5

0p

M0

1.7

99.9

3··

···

·A

CC

Ma

07

08

23.2

400

-19

40

01.5

47

12.8

02.5

0p

M0

0.7

517.9

1··

···

·A

BC

Mi

07

07

57.3

118

+11

58

18.8

83

13.4

02.5

0p

M0

0.2

6··

···

···

·H

OM

on

07

10

19.2

100

+00

25

28.9

49

11.4

02.8

0P

B0

1.0

522.0

2A

5+

[K2.5

IV]

SA

AQ

Mon

07

14

17.6

748

-07

13

45.0

60

9.7

92.7

1V

TB

11.4

821.3

9A

0+

[A4.5

]SA

FP

Mon

07

15

08.8

358

-09

57

47.6

39

13.3

03.0

0p

M0

0.8

57.0

5F

6V

···

AF

CM

a07

15

22.7

708

-23

40

42.6

24

12.3

03.2

0p

M0

0.6

733.7

6··

···

·

Table3continued

16 Bellotti et al.Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

HS

Cam

07

19

14.5

151

+65

57

44.3

05

19.4

04.0

0B

M0

2.1

9··

·W

D+

M4-6

VS2C

RY

CM

i07

22

58.1

591

+06

46

35.0

19

11.9

03.0

0p

M0

0.5

5··

···

·S

CX

Pup

07

41

53.5

171

-22

53

36.9

24

12.2

03.0

0p

M0

0.8

4··

···

···

·V

0681

Mon

07

52

21.9

573

-01

17

25.1

49

13.0

02.5

0b

M0

0.8

116.1

8··

···

·T

UM

on

07

53

19.7

487

-03

02

31.1

37

9.0

02.7

0V

HB

10.6

1··

·B

5V

+A

5II

I:SH

MT

Pup

07

54

11.8

928

-14

39

16.6

33

15.0

04.0

0p

M0

1.0

1··

·em

···

FG

Pup

08

04

31.6

979

-24

02

53.4

06

13.0

03.0

0p

M0

0.4

622.1

4··

···

·IO

Pup

08

06

57.9

454

-25

49

08.2

88

13.5

02.5

0p

M0

0.4

8··

···

···

·A

VV

el

08

07

01.3

411

-47

38

46.0

07

11.9

02.7

0p

M0

0.8

914.3

2··

···

·F

IP

up

08

08

14.7

308

-20

17

14.8

82

14.0

02.5

0p

M0

0.2

4··

···

···

·X

ZP

up

08

13

31.0

553

-23

57

11.3

87

7.7

52.5

1V

M1

2.5

4··

·A

0SA

FQ

Pup

08

18

10.5

357

-23

48

40.5

71

14.0

02.5

0p

M0

0.4

38.1

1··

···

·F

RP

up

08

21

19.9

553

-22

18

45.7

03

13.5

03.0

0p

M0

0.2

7··

···

···

·D

RL

yn

08

24

24.4

675

+50

00

50.8

91

11.6

02.7

0V

M0

1.2

46.4

2··

···

·D

EH

ya

08

27

47.5

470

+05

38

58.9

39

11.0

03.0

0V

B0

···

···

A2+

[K2IV

]SA

SY

Hya

08

29

51.6

786

-09

23

57.4

96

10.7

02.9

0V

TB

00.9

3··

·A

5+

[K2IV

]SA

VW

Hya

08

33

51.0

321

-14

39

53.8

91

10.5

03.6

0V

B1

0.8

0··

·A

3+

[G4IV

]SA

RY

Cnc

08

39

54.6

211

+19

49

18.8

64

12.9

92.5

3V

B0

0.8

57.4

6B

8+

[G5IV

]S

TY

Cnc

08

47

09.6

962

+08

24

24.5

82

12.7

03.0

0p

M0

0.6

9··

···

···

·T

UC

nc

08

52

16.6

529

+09

05

18.7

81

9.9

02.5

0P

B0

1.4

228.5

4A

0+

[G8

SA

AC

UM

a08

55

54.1

228

+64

58

14.9

99

10.3

03.7

0V

B0

1.2

441.1

4A

2+

[K6IV

]SA

RX

Hya

09

05

41.1

657

-08

15

39.7

27

8.9

02.7

0V

B0

3.7

18.9

8A

8+

[K5IV

]SA

TY

Hya

09

29

02.3

991

+05

34

28.2

69

10.5

03.0

0P

B0

0.7

637.7

9A

0+

[K1IV

]SA

YL

eo

09

36

51.8

058

+26

13

57.6

43

10.0

93.1

1V

M1

2.4

911.9

4A

5V

+K

5V

SA

DX

Vel

09

51

45.9

871

-50

53

23.7

16

10.0

83.4

2V

TB

11.8

512.1

2A

5+

[G1IV

]SA

IVC

ar

10

02

36.6

688

-58

57

22.4

60

11.3

02.8

0p

M0

1.2

9··

···

···

·H

LC

ar

10

04

42.7

497

-61

50

35.0

84

11.0

03.0

0p

M0

0.7

753.0

9··

···

·T

TV

el

10

20

16.4

606

-46

14

04.4

36

10.1

03.2

4V

TB

11.9

512.5

4A

5+

[F8.5

]SA

AW

Car

10

32

35.0

208

-60

46

10.6

53

12.2

04.3

0p

M0

0.3

439.1

1A

3··

·F

VC

ar

10

42

31.9

233

-61

57

55.0

72

12.9

02.9

0p

M0

0.5

760.5

5··

·S

PT

Car

11

03

17.6

000

-74

56

37.0

00

12.8

03.2

0p

M0

···

···

···

···

EH

Car

11

04

23.8

163

-61

18

23.0

20

12.5

02.5

0p

M0

0.3

047.3

9··

···

·A

NU

Ma

11

04

25.6

556

+45

03

13.9

41

14.9

05.3

0B

M0

3.1

0··

···

·S2C

DE

Car

11

06

13.6

515

-60

47

32.6

72

11.2

02.6

0p

M1

0.9

417.8

5A

SA

V0443

Cen

11

25

22.5

000

-59

37

40.0

00

12.8

02.5

0b

M0

···

···

···

···

AB

Cen

11

26

21.2

641

-58

24

58.1

36

10.6

02.6

0p

M0

0.6

285.9

3G

0··

·V

0646

Cen

11

36

58.7

768

-53

12

35.4

14

9.0

02.7

0b

M0

1.4

0··

·B

8IV

S

V0348

Cen

11

48

28.9

566

-43

46

53.2

64

10.7

02.7

0b

M0

1.2

614.7

3A

···

Table3continued

Detecting Exoplanets with Eclipsing Binaries 17Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

ZC

rv12

29

41.7

139

-23

38

03.8

78

14.0

02.5

0V

M0

1.2

11.3

2··

···

·R

RC

rv12

30

26.9

456

-17

58

53.4

49

11.1

02.8

0p

M0

1.3

312.1

5··

···

·B

PM

us

12

50

37.7

277

-71

46

18.6

98

9.6

03.0

0V

M0

1.7

521.6

5A

0,5

/1,5

V+

G5II

I··

·C

TC

en

13

10

43.0

552

-58

16

38.6

29

10.3

02.7

0p

M0

1.5

641.8

7A

3(e

a)

SA

CY

Cen

13

12

13.0

891

-52

52

51.8

52

13.3

02.8

0p

M0

0.7

06.8

9··

···

·U

WV

ir13

15

20.7

356

-17

28

16.9

24

8.8

43.1

6V

B0

4.2

710.3

5A

4+

[K3IV

]SA

BM

Cen

13

17

38.4

543

-56

17

09.7

65

13.2

02.8

0p

M0

0.6

97.5

1··

···

·B

PC

en

13

19

08.6

963

-49

55

04.7

51

12.2

03.8

0p

M0

1.1

4··

···

···

·SX

Hya

13

44

37.3

594

-26

46

48.3

76

8.6

04.0

0V

HB

13.8

616.2

3A

3+

[G9IV

]SA

IQC

en

14

00

34.4

391

-54

16

15.6

69

12.7

03.0

0p

M0

0.5

128.3

2··

···

·A

NC

ir14

04

21.9

000

-65

38

05.0

00

13.8

02.8

0p

M0

···

···

···

···

GK

Vir

14

15

36.4

126

+01

17

18.2

27

17.0

15.9

9V

M0

2.1

1··

·D

A0+

M3-5

VD

W

TU

Lup

14

43

40.6

725

-49

50

24.7

52

12.4

03.8

0p

M0

0.7

211.7

3··

···

·W

Cir

15

11

47.0

000

-55

38

12.0

00

13.8

02.7

0b

M0

···

···

···

···

EV

Lup

15

16

08.0

026

-44

19

20.3

80

9.8

02.7

0p

M0

1.0

533.6

5··

···

·T

WD

ra15

33

51.0

598

+63

54

25.7

06

8.0

02.5

0p

M1

6.0

122.7

5A

8V

+K

0II

ISA

LX

Ser

15

38

00.0

940

+18

52

03.2

46

13.3

04.1

0B

M0

2.0

3··

·W

D+

M32V

S2C

HH

Nor

15

43

30.1

810

-51

50

47.5

42

9.7

03.0

5V

TB

12.2

215.3

1A

3+

[G8IV

]S

YY

Nor

15

47

58.4

618

-57

24

42.6

26

13.2

02.9

0p

M0

0.7

612.1

3··

···

·H

KN

or

15

58

54.7

069

-51

41

54.9

60

12.6

02.6

0p

M0

0.8

83.6

9··

···

·H

SN

or

16

01

44.9

307

-57

27

33.1

67

12.6

02.6

0p

M0

0.4

125.4

1··

···

·B

HA

ps

16

08

40.2

241

-76

50

28.0

71

10.8

02.6

0p

M0

1.0

0··

···

···

·C

EN

or

16

11

13.2

040

-60

02

20.9

28

11.5

03.0

0p

M0

0.6

0··

···

···

·C

CH

er

16

17

38.8

937

+08

56

02.6

10

10.2

02.9

0p

M0

2.8

210.1

8A

1V

(pSr)

SA

UU

Nor

16

20

21.5

635

-53

44

52.2

36

11.6

02.7

0p

M0

0.3

778.9

1··

···

·V

0704

Ara

16

56

16.9

143

-55

27

58.3

61

12.9

02.7

0b

M0

0.4

613.0

9··

···

·U

UO

ph

16

57

22.6

406

-25

47

58.5

63

10.0

02.5

0V

HB

00.7

6··

·A

0+

[K2IV

]SA

V0588

Sco

16

57

49.7

872

-40

29

59.9

53

14.0

02.5

0b

M0

0.5

87.0

9··

···

·V

0416

Ara

16

58

06.6

680

-55

30

12.2

52

12.6

02.9

0b

M0

0.3

628.3

9··

···

·IM

Sco

17

00

13.2

531

-31

27

57.6

09

13.3

03.2

0p

M0

0.6

44.2

9G

2V

···

V1194

Oph

17

00

21.1

176

-21

51

22.3

43

17.6

02.7

0p

M0

0.2

4··

···

···

·V

735

Oph

17

07

48.9

680

+09

33

09.2

06

10.3

02.9

6V

TB

01.7

811.2

5A

0+

[G1IV

]SA

V0460

Sco

17

08

50.4

000

-32

54

05.0

00

13.4

02.6

0b

M0

···

···

···

···

V1453

Oph

17

08

52.7

155

-19

01

08.9

33

17.4

03.0

0p

M0

0.3

0··

···

···

·F

TSco

17

10

55.2

000

-32

41

09.0

00

13.8

02.7

0p

M0

···

···

···

···

V1578

Oph

17

12

14.2

156

-20

14

54.8

64

16.4

03.4

0p

M0

0.1

1··

···

···

·V

1590

Oph

17

12

22.2

904

-16

24

34.0

19

16.0

02.9

0p

M0

0.2

6··

···

···

·T

UH

er

17

13

35.3

658

+30

42

36.0

42

10.8

82.8

2V

M1

1.9

69.1

0F

0II

I/IV

SA

Table3continued

18 Bellotti et al.Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

V0621

Sco

17

14

36.8

000

-41

04

15.0

00

12.3

02.7

0b

M0

···

···

···

···

V0441

Oph

17

20

52.7

094

-17

20

05.2

61

11.6

03.0

0b

M0

1.2

616.8

3A

0··

·LT

Ara

17

22

30.1

256

-47

04

02.9

89

10.8

04.6

0p

M0

0.8

739.2

5A

···

V0506

Sco

17

23

12.4

000

-39

11

40.0

00

13.3

02.7

0b

M0

1.0

2··

···

···

·V

0510

Sco

17

27

37.2

557

-42

58

53.2

76

13.1

03.0

0b

M0

0.4

235.9

2··

···

·V

0511

Sco

17

28

22.5

406

-42

02

39.0

60

12.5

03.5

0b

M0

0.7

918.6

4··

···

·V

0532

Oph

17

32

42.6

073

-21

51

40.7

60

12.6

02.9

0b

M0

0.1

7··

···

···

·R

WA

ra17

34

49.2

233

-57

08

50.5

66

8.8

52.6

0V

M1

2.5

3··

·A

1IV

SA

V0525

Sco

17

39

57.7

671

-40

29

13.3

85

13.8

02.6

0b

M0

0.5

811.5

1··

···

·V

0529

Oph

17

41

56.6

000

-28

18

10.0

00

13.2

03.3

0b

M0

···

···

···

···

AK

Ser

17

42

04.4

088

-13

53

11.9

71

10.8

02.8

0p

M1

1.3

3··

·A

5SA

V0755

Sgr

17

48

46.1

804

-25

53

09.9

93

13.5

02.5

0b

M0

0.6

1··

···

···

·V

0765

Sgr

17

51

54.0

858

-28

26

32.0

28

12.8

03.7

0b

M0

0.3

746.2

9··

···

·V

913

Oph

17

55

03.5

918

+14

10

39.2

54

11.5

03.0

0P

B0

2.1

45.1

9A

5+

[G5IV

]SA

V391

Oph

17

58

09.1

304

+04

39

27.6

10

11.5

03.5

0P

B1

0.7

012.9

3A

1+

[G5IV

]SA

V2301

Oph

18

00

35.5

318

+08

10

13.9

20

16.0

06.0

0V

M0

8.2

4··

·W

D+

M6V

S2C

V0573

Oph

18

06

05.3

800

+02

05

43.8

28

13.1

03.0

0b

M0

0.4

819.5

1··

···

·V

1178

Sgr

18

11

03.8

499

-30

30

56.8

84

12.8

02.5

0p

M0

0.0

8··

···

···

·A

GP

av

18

12

25.6

276

-62

33

21.8

59

10.2

02.8

0p

M0

1.6

210.4

8··

···

·V

1282

Sgr

18

13

34.6

857

-34

28

42.9

44

11.9

04.3

0p

M0

1.0

813.0

3··

···

·V

1959

Sgr

18

16

43.0

576

-25

31

20.2

68

12.6

03.2

0p

M0

0.5

213.0

3··

···

·A

ISgr

18

16

45.0

943

-21

34

48.5

29

11.8

03.2

0p

M0

0.4

387.4

1··

···

·V

0710

Sgr

18

16

46.8

091

-36

25

32.1

60

13.3

03.0

0b

M0

0.5

1··

···

···

·V

3171

Sgr

18

22

44.2

273

-34

13

32.2

05

16.0

03.2

0p

M0

0.1

6··

···

···

·V

2537

Sgr

18

23

00.3

07

-32

34

03.0

30

12.3

02.5

0p

M0

···

···

···

···

V1654

Sgr

18

23

17.1

000

-23

15

10.0

00

13.2

02.6

0p

M0

···

···

···

···

V0586

Oph

18

27

13.9

427

04

17

15.3

19

13.3

02.5

0b

M0

0.7

714.1

9··

···

·V

3421

Sgr

18

29

11.6

135

-32

23

55.7

11

16.5

02.6

0p

M0

0.2

8··

···

···

·V

3476

Sgr

18

30

49.5

100

-32

57

51.0

00

14.9

02.5

0p

M0

···

···

···

···

EY

Ser

18

30

49.8

137

+05

30

17.6

29

14.5

03.0

0p

M0

0.5

412.6

0··

···

·V

4727

Sgr

18

33

30.7

141

-28

58

51.2

39

10.6

03.2

0V

M0

1.1

514.6

1··

···

·V

3622

Sgr

18

35

46.2

499

-34

52

30.6

16

14.8

03.0

0p

M0

0.2

1··

···

···

·V

0372

Sct

18

38

03.6

563

-04

01

06.1

18

14.2

02.6

0b

M0

0.3

554.5

8··

···

·V

3696

Sgr

18

38

46.8

768

-34

39

03.8

02

16.3

03.0

0p

M0

0.2

3··

···

···

·B

OH

er

18

40

30.0

965

+24

55

42.7

68

10.7

03.1

0V

M1

1.2

917.0

8A

7V

SA

V3752

Sgr

18

40

43.9

834

-35

32

26.3

40

15.9

02.6

0p

M0

0.2

9··

···

···

·R

RD

ra18

41

47.4

012

+62

40

34.9

44

10.0

03.3

0V

B0

2.1

720.2

3A

2+

[G8IV

]SA

V1933

Sgr

18

45

10.1

332

-19

47

19.9

60

12.5

02.5

0p

M0

0.6

222.2

0··

···

·

Table3continued

Detecting Exoplanets with Eclipsing Binaries 19Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

AC

Sct

18

46

01.4

709

-10

14

55.8

15

10.0

02.6

0V

B0

1.0

163.8

2B

9+

[G0IV

]SA

V0795

Aql

18

47

06.8

639

+11

40

11.0

52

13.4

03.0

0b

M0

0.4

020.7

0··

···

·D

HH

er

18

47

34.5

584

+22

50

45.7

95

9.4

02.6

0V

M1

0.3

3··

·A

5SA

V0370

Sct

18

48

31.2

299

-05

42

37.4

49

15.3

03.2

0b

M0

0.4

68.3

5··

···

·A

DH

er

18

50

00.3

000

+20

43

16.5

09

9.3

810.7

4V

B1

1.7

137.5

8A

4V

+[G

9IV

]SA

FN

Sct

18

50

46.7

234

-05

12

07.9

16

12.7

02.5

0B

M0

0.3

5108.3

8B

3V

e··

·C

TSct

18

54

21.6

365

-06

00

16.6

96

10.0

24.1

7V

B1

0.2

6631.5

3B

7V

n+

[A3]

SA

LP

Her

18

55

41.2

099

+12

15

28.4

00

13.0

03.2

0p

M0

0.6

516.5

6··

···

·E

HSct

18

57

39.7

600

-06

52

59.7

00

14.8

02.7

0p

M0

···

···

···

···

BN

Sct

18

58

42.8

687

-08

20

28.3

64

11.4

03.2

0p

M0

0.8

244.0

6A

0Ib

···

HY

Aql

19

02

47.6

109

-06

13

33.0

67

12.8

02.8

0p

M0

0.6

017.2

1··

···

·F

KA

ql

19

04

18.5

775

+02

46

47.2

44

11.1

02.4

0V

TB

01.3

819.7

7B

9+

[G6IV

]SA

V1074

Sgr

19

05

23.6

885

-18

47

06.8

64

12.7

02.5

0p

M0

1.2

67.4

9F

4II

I··

·V

1109

Aql

19

05

40.1

991

+14

14

43.6

86

17.0

02.5

0p

M0

0.1

9··

···

···

·E

PD

ra19

07

06.1

869

+69

08

43.8

74

17.6

02.9

0V

M0

2.5

9··

···

·S2C

V0413

Lyr

19

07

16.4

316

+30

19

25.5

90

14.5

02.6

0b

M0

0.2

321.6

5··

···

·V

0449

Lyr

19

07

37.0

709

+44

00

19.7

97

12.5

03.5

0b

M0

0.1

8··

···

···

·V

0415

Sgr

19

07

37.0

938

-23

37

08.5

43

12.3

03.0

0b

M0

0.5

538.8

6··

···

·V

1112

Aql

19

07

39.1

500

-00

23

07.1

00

14.0

02.5

0p

M0

···

···

···

···

KK

Dra

19

07

56.6

643

+59

23

52.0

65

11.8

03.0

0V

M0

1.6

93.6

3··

···

·V

0407

Aql

19

11

10.8

162

+01

08

52.1

45

13.3

02.5

0b

M0

1.2

44.3

3··

···

·N

SL

yr

19

11

13.9

662

+36

10

35.3

21

14.0

02.8

0p

M0

0.4

89.3

2··

···

·V

354

Sgr

19

13

23.4

427

-18

29

00.0

14

10.7

03.0

0V

TB

11.5

019.0

5F

8+

[K9IV

]SA

V1103

Cyg

19

14

57.6

721

+46

10

01.6

22

15.0

02.5

0p

M0

0.3

010.2

9··

···

·R

VL

yr

19

16

17.9

644

+32

25

14.9

69

11.5

03.1

0P

B0

1.0

222.0

1A

5+

[K3IV

]SA

YZ

Aql

19

16

46.2

253

-00

36

17.1

52

10.5

03.7

0V

B1

0.9

245.9

4A

3+

[K5IV

]SA

V342

Aql

19

17

03.4

768

+09

20

38.5

46

9.5

03.4

0P

B0

3.2

125.5

9A

4II

+[K

0IV

]SA

USge

19

18

48.4

083

+19

36

37.7

20

6.5

02.6

0V

B1

3.6

2··

·B

8II

I+K

SA

FR

Vul

19

36

24.8

396

+26

45

56.5

58

9.9

12.5

4V

TB

11.8

016.2

0A

2+

[G7IV

]SA

V0418

Aql

19

36

39.4

426

+03

57

00.4

70

12.8

03.2

0b

M0

1.0

66.7

2··

···

·LT

Aql

19

38

49.7

896

+06

34

59.0

39

12.4

02.5

0p

M0

0.5

930.2

7B

9II

I/IV

···

V0932

Cyg

19

39

05.8

610

+33

29

01.6

74

14.0

02.6

0b

M0

0.2

238.8

4··

···

·E

NC

yg

19

40

09.4

635

+29

16

23.1

00

12.9

03.2

0p

M0

1.0

67.7

6··

···

·V

2168

Sgr

19

42

37.5

266

-38

39

55.6

48

12.5

02.5

0p

M0

0.7

67.2

3··

···

·SY

Cyg

19

46

34.3

155

+32

42

18.4

63

10.7

03.5

0P

B0

0.9

439.9

5A

3+

[K6IV

]SA

V1033

Aql

19

46

44.7

712

+14

21

57.3

79

15.0

02.5

0p

M0

0.2

0··

···

···

·V

0814

Cyg

19

49

50.7

520

+36

34

27.6

59

15.0

03.0

0b

M0

0.1

7··

···

···

·V

0689

Cyg

19

51

06.5

892

+36

51

32.7

21

14.0

02.5

0b

M0

0.6

56.3

5A

5V

···

Table3continued

20 Bellotti et al.Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

V524

Sgr

19

53

14.5

011

-14

54

38.8

82

9.9

62.5

4V

TB

12.1

612.2

5F

8II

I+[K

3IV

]SA

V0340

Aql

19

55

56.4

890

+15

51

07.1

67

11.5

02.6

0b

M0

0.8

816.3

0F

···

BO

Vul

19

56

29.0

711

+23

54

45.0

11

10.5

02.8

0p

M0

2.7

99.7

6F

0+

G0IV

SA

PV

Cyg

19

56

29.1

748

+37

43

09.1

24

12.7

02.8

0p

M0

0.5

813.7

1A

1··

·V

0691

Cyg

19

57

18.4

468

+40

02

55.7

12

15.5

02.5

0b

M0

0.3

56.5

0··

···

·Q

TC

yg

19

58

07.3

701

+38

49

27.8

87

14.8

02.6

0p

M0

0.2

615.7

7A

5··

·V

4140

Sgr

19

58

49.7

022

-38

56

13.2

25

15.5

02.5

0p

M0

1.6

7··

·W

D+

M6-7

S2C

QX

Sge

19

59

36.7

480

+20

48

14.5

99

20.4

02.6

0V

M0

···

···

···

···

V0698

Cyg

19

59

53.3

495

+36

16

39.9

98

12.2

02.7

0b

M0

0.2

8··

·B

2··

·V

1174

Cyg

20

03

55.9

791

+31

15

39.9

47

14.5

02.5

0p

M0

0.3

138.6

7··

···

·W

WC

yg

20

04

02.7

111

+41

35

16.4

62

10.0

23.2

4V

B0

0.8

0··

·B

7V

+[G

1IV

]SA

EX

Vul

20

04

47.4

096

+22

19

21.8

61

12.4

92.6

1B

M0

0.5

152.4

7ea

···

SW

Cyg

20

06

57.9

310

+46

17

58.1

47

9.2

42.5

9V

M1

2.1

233.4

6A

2e+

K0

SA

V1037

Cyg

20

08

49.5

433

+35

14

55.5

07

14.7

02.6

0p

M0

0.5

18.0

5··

···

·Q

YA

ql

20

09

28.8

293

+15

18

44.7

18

11.4

03.2

0V

HB

00.7

857.1

8F

0+

/K

3IV

/SA

V1412

Aql

20

13

55.6

789

+06

42

44.8

26

15.6

72.6

3V

M0

43.6

2··

·D

C7

···

V0445

Aql

20

19

38.8

300

+06

08

47.4

00

13.0

02.5

0b

M0

···

···

···

···

VSge

20

20

14.6

910

+21

06

17.1

78.6

05.3

0V

M0

0.4

2··

·W

N5

S2C

V1320

Cyg

20

21

22.0

535

+39

19

57.9

24

15.9

02.6

0p

M0

0.4

414.0

6··

···

·V

0445

Cyg

20

28

18.9

582

+38

17

43.2

26

11.7

03.8

0b

M0

1.2

87.3

6··

···

·V

1051

Cyg

20

31

00.5

678

+56

46

31.1

59

14.4

02.8

0p

M0

0.5

25.5

2··

···

·L

LV

ul

20

31

43.3

028

+25

38

36.3

03

16.0

03.0

0p

M0

0.0

9··

···

···

·H

ND

el

20

33

45.1

355

+11

03

40.8

75

14.5

02.5

0p

M0

0.5

94.9

6··

···

·SY

Vul

20

36

21.8

687

+23

51

51.0

38

13.0

03.0

0p

M0

0.2

829.2

6··

···

·K

KD

el

20

36

42.7

618

+17

28

48.0

78

15.0

02.5

0p

M0

0.9

8··

···

···

·W

Del

20

37

40.0

857

+18

17

03.7

52

9.6

92.6

4V

M1

1.2

3··

·B

9.5

Ve+

G5IV

SA

V1204

Cyg

20

44

10.7

490

+46

49

03.7

56

14.5

03.0

0p

M0

0.4

86.9

3··

···

·V

1843

Cyg

20

46

27.7

000

+34

03

53.0

00

15.7

02.8

0B

M0

0.2

0··

···

···

·V

0398

Cyg

20

46

50.7

555

+34

12

03.7

17

12.5

02.5

0b

M0

0.7

014.2

0··

···

·D

WC

ep

20

51

39.6

899

+62

48

50.2

92

10.2

52.6

5V

TB

11.2

120.0

2B

8+

[F5]

SA

V1870

Cyg

20

51

41.0

469

+35

44

08.1

01

14.3

02.9

0B

M0

0.8

64.3

7··

···

·V

1884

Cyg

20

56

07.1

165

+33

39

07.1

45

15.2

02.5

0B

M0

0.2

527.2

5··

···

·V

1718

Cyg

21

02

01.2

000

+41

31

57.0

00

14.7

02.8

0p

M0

···

···

···

···

V0377

Cyg

21

03

12.1

952

+29

07

11.8

78

14.4

02.6

0b

M0

0.2

833.7

1··

···

·LY

Del

21

06

26.0

978

+19

24

36.4

22

10.4

03.1

0V

M0

0.5

4129.2

7··

···

·V

1539

Oph

21

11

20.0

900

-20

29

04.9

67

15.6

03.2

0p

M0

0.4

05.3

3··

···

·V

1960

Cyg

21

12

41.9

861

+37

32

28.2

12

15.3

02.5

0B

M0

0.2

612.0

1··

···

·V

0534

Cyg

21

21

08.4

383

+44

15

04.3

82

13.7

03.0

0b

M0

0.6

08.4

9··

···

·

Table3continued

Detecting Exoplanets with Eclipsing Binaries 21Table

3(continued)

Nam

eR

AD

ec

mD

mag

Fil

ter

Ref

Fla

gπ

Ltot

Sp

ec

Typ

eE

vol

Cla

ss

[hh

mm

ss]

[dd

mm

ss]

[mas]

[L⊙ ]

(1)

(2)

(3)

(4)

(5)

(6)

(7)

(8)

(9)

(10)

(11)

(12)

V1727

Cyg

21

31

26.2

132

+47

17

24.5

14

15.6

02.9

0V

M0

0.5

5··

·p

ec(e

)S2L

UG

ru21

31

48.7

723

-45

02

42.3

24

11.0

04.0

0P

B1

1.4

518.6

7A

5+

[K0IV

]SA

V1667

Cyg

21

32

33.3

186

+34

27

06.5

81

14.9

02.5

0B

M0

3.2

81.5

0··

···

·B

QP

eg

21

34

20.6

566

+20

57

18.1

18

13.7

02.6

0p

M0

0.6

18.7

0··

···

·V

0705

Cyg

21

34

56.6

008

+43

01

28.6

16

13.6

02.8

0b

M0

0.4

718.3

0··

···

·V

1618

Cyg

21

45

54.1

084

+38

45

26.7

37

15.4

02.6

0p

M0

0.3

08.3

1··

···

·C

WP

eg

21

48

27.6

008

+28

06

29.1

91

11.8

03.4

3V

M1

1.0

2··

···

···

·D

OP

eg

22

07

30.6

309

+06

10

16.5

510.6

02.9

0V

B0

0.9

0··

·B

8+

[G4IV

]SA

EL

Lac

22

08

53.7

317

+42

16

20.9

67

12.4

03.6

0p

M0

0.6

412.5

3··

···

·T

VC

ep

22

09

53.9

531

+63

07

16.4

09

12.2

02.5

0p

M0

0.8

66.4

7K

5V

···

ER

Lac

22

19

03.1

800

+51

41

08.9

00

14.0

02.5

0p

M0

···

···

···

···

BS

Lac

22

19

37.2

753

+44

17

03.1

17

13.5

02.7

0p

M0

0.9

46.6

9··

···

·B

RC

ep

22

27

17.1

259

+66

10

00.3

00

12.5

02.5

0p

M0

0.8

520.5

1A

3··

·D

YL

ac

22

47

19.2

341

+53

59

06.3

04

14.6

03.0

0p

M0

0.3

016.2

5··

···

·E

HL

ac

22

51

55.6

894

+51

25

14.1

20

13.6

03.2

0p

M0

0.3

414.9

2··

···

·H

IL

ac

22

56

48.6

620

+53

47

47.6

10

15.0

03.0

0p

M0

0.7

42.0

4··

···

·B

OA

nd

22

58

38.0

000

+45

31

52.0

00

13.4

02.9

0P

B0

0.3

437.6

B8+

[G6IV

]SA

V0341

Cas

22

59

34.7

018

+56

23

09.7

59

14.7

02.6

0b

M0

0.5

28.7

7··

···

·C

UA

nd

23

01

01.5

967

+49

58

25.3

18

12.5

03.5

0p

M0

0.9

26.7

1··

···

·V

0570

Cas

23

16

27.5

700

+59

48

18.1

94

13.8

02.7

0B

M0

1.2

70.6

7F

-G··

·X

Gru

23

19

42.3

634

-55

36

41.4

85

10.6

43.6

4V

B1

1.4

7··

·A

0SA

YP

sc23

34

25.3

848

+07

55

28.5

24

10.1

03.0

0p

M1

2.3

026.9

7A

3+

K0IV

SA

V0442

Cas

23

40

14.7

999

+53

57

33.9

90

13.2

03.8

0b

M0

0.5

621.8

5··

···

·V

ZScl

23

50

09.2

550

-26

22

52.7

01

15.6

02.8

0V

M0

1.7

8··

·p

ec(e

)S2C

QS

Cas

23

51

58.6

213

+56

02

31.8

91

13.6

02.9

0p

M0

0.3

417.3

2··

···

·

Note—

Eclipsi

ng

bin

ary

syst

em

sw

ith

pri

mary

min

imum

depth

gre

ate

rth

an

or

equal

to2.5

.(1

)G

enera

lC

ata

log

of

Vari

able

Sta

rsdesi

gnati

on,

(2)

Rig

ht

asc

ensi

on,

(3)

Declinati

on,

(4)

Magnit

ude

at

maxim

um

bri

ghtn

ess

,(5

)D

epth

of

pri

mary

min

imum

,(6

)P

hoto

metr

icfilt

er

use

dto

obta

inth

elight

curv

e,

(7)

Refe

rence

of

targ

et,

maxim

um

bri

ghtn

ess

,pri

mary

min

imum

,and

filt

er

(Bis

Buddin

get

al.

(2004)

and

Mis

Malk

ov

et

al.

(2006))

,(8

)D

iscre

pancy

flag,

(9)

Para

llax,

(10)

Maxim

um

bolo

metr

iclu

min

osi

ty,

(11)

Sp

ectr

al

typ

e,

(12)

Evolu

tionary

cla

ss.

Colu

mns

(2)

and

(3)

are

extr

acte

dfr

om

SIM

BA

Dand

refe

rto

ep

och

J2000,

colu

mn

(11)

isfr

om

Buddin

get

al.

(2004),

Malk

ov

et

al.

(2006),

and

Avvakum

ova

et

al.

(2013),

and

colu

mn

(9)

and

(10)

are

from

theGaia

DR

2data

base

.If

,fo

ra

giv

en

targ

et,

the

cata

logs

rep

ort

signifi

cantl

ydiff

ere

nt

valu

es

of

Dm

ag

inth

esa

me

photo

metr

icfilt

er,

we

exclu

de

that

targ

et

from

the

analy

sis,

but

we

rep

ort

them

here

for

com

ple

teness

.A

valu

eof

0and

1fo

rth

edis

cre

pancy

flag

repre

sents

consi

stent

and

dis

cre

pant

Dm

ag

valu

es

resp

ecti

vely

.

22 Bellotti et al.

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