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

0.96

0.98

1.00

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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-0.4

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his

Wor

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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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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)

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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)

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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)

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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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-0.5 0.0 0.5[M/H]

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)