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Preprint typeset using LATEX style emulateapj v. 10/09/06
MEASUREMENTS OF STELLAR INCLINATIONS FOR KEPLER PLANET CANDIDATES II: CANDIDATE
SPIN-ORBIT MISALIGNMENTS IN SINGLE AND MULTIPLE-TRANSITING SYSTEMS
Teruyuki Hirano1, Roberto Sanchis-Ojeda2, Yoichi Takeda3, Joshua N. Winn2, Norio Narita3, and Yasuhiro H.
Takahashi3,4
ABSTRACT
We present a test for spin-orbit alignment for the host stars of 25 candidate planetary systems
detected by the Kepler spacecraft. The inclination angle of each star’s rotation axis was estimated
from its rotation period, rotational line broadening, and radius. The rotation periods were deter-
mined using the Kepler photometric time series. The rotational line broadening was determined from
high-resolution optical spectra with Subaru/HDS. Those same spectra were used to determine the
star’s photospheric parameters (effective temperature, surface gravity, metallicity) which were then
interpreted with stellar-evolutionary models to determine stellar radii. We combine the new sam-
ple with the 7 stars from our previous work on this subject, finding that the stars show a statistical
tendency to have inclinations near 90◦, in alignment with the planetary orbits. Possible spin-orbit mis-
alignments are seen in several systems, including three multiple-planet systems (KOI-304, 988, 2261).
Ideally these systems should be scrutinized with complementary techniques—such as the Rossiter-
McLaughlin effect, starspot-crossing anomalies or asteroseismology—but the measurements will be
difficult owing to the relatively faint apparent magnitudes and small transit signals in these systems.
Subject headings: planets and satellites: general – planets and satellites: formation – stars: rotation
– techniques: spectroscopic
1. INTRODUCTION
The angle of the stellar spin axis with respect to the
planetary orbital axis (spin-orbit angle) is an observ-
able quantity that may be important for understand-
ing the evolutionary history of exoplanetary systems. In
order to explain the existence of close-in giant planets
(hot Jupiters or Neptunes), various migration scenarios
have been proposed, which differ in their predictions for
the spin-orbit angle. Some theories, such as disk migra-
tion, predict that the stellar spin and planetary orbital
axes should be well aligned (e.g., Lin et al. 1996). Other
theories, such as planet-planet scattering or Kozai mi-
gration, predict a very wide range of spin-orbit angles
(see, e.g., Wu & Murray 2003; Nagasawa & Ida 2011;
Fabrycky & Tremaine 2007).
Most of the current measurements of the spin-
orbit angle have been based on observations of the
Rossiter-McLaughlin (RM) effect (e.g., Queloz et al.
2000; Ohta et al. 2005; Winn et al. 2005; Narita et al.
2007; Wolf et al. 2007; Hirano et al. 2011b) or pho-
tometric anomalies due to transits over starspots
(e.g., Sanchis-Ojeda et al. 2011; Nutzman et al. 2011;
De´sert et al. 2011). These measurements have revealed
a diversity of spin-orbit angles (e.g., He´brard et al. 2008;
Winn et al. 2009; Narita et al. 2009). This diversity has
inspired many theoretical studies of the possible rea-
sons for highly inclined planetary orbits (e.g., Lai et al.
Electronic address: hirano@geo.titech.ac.jp
1 Department of Earth and Planetary Sciences, Tokyo Insti-
tute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo 152-8551,
Japan
2 Department of Physics, and Kavli Institute for Astrophysics
and Space Research, Massachusetts Institute of Technology, Cam-
bridge, MA 02139
3 National Astronomical Observatory of Japan, 2-21-1 Osawa,
Mitaka, Tokyo, 181-8588, Japan
4 Department of Astronomy, The University of Tokyo, Tokyo,
113-0033, Japan
2011; Naoz et al. 2011). The measurements have also re-
vealed some possible patterns relating the spin-orbit an-
gle and the properties of the host stars (Winn et al. 2010;
Albrecht et al. 2012). However, the existing measure-
ments have been almost exclusively restricted to close-in
giant planets. This is simply because the preceding mea-
surement techniques are best suited to relatively large
planets, which produce stronger spectroscopic or photo-
metric signals during a planetary transit. Thus, the spin-
orbit relations for smaller planets have been unknown
until recently.
An important step was taken by Schlaufman (2010),
who demonstrated that the stellar inclination angle (the
angle between the stellar spin axis and the line of sight)
can be readily estimated for a large number of transiting
exoplanetary systems, and used to probe spin-orbit align-
ment. The basic idea is to use estimates of the rotation
velocity V and the projected rotation velocity V sin Is to
determine sin Is. Since the orbital axis of a transiting
planet must be nearly perpendicular to the line-of-sight
(sin Io ≈ 1), a small value of sin Is implies a spin-orbit
misalignment.
The pioneering analysis of Schlaufman (2010) was
based on spectroscopic determinations of V sin Is, as well
as statistical estimates of V based on the rotation-age-
mass correlations that are observed for main-sequence
stars. It is also possible to measure V more directly,
if accurate estimates of the stellar radius Rs and ro-
tation period Ps are available, using the relation V =
2piRs/Ps (see, e.g., Doyle et al. 1984). It has also be-
come possible to estimate sin Is using asteroseismology
(e.g., Chaplin et al. 2013).
An important advantage of this technique is that the
difficulty of measuring stellar inclinations is independent
of the size of the transiting planet, and therefore the
spin-orbit relation may be investigated even for smaller
planets (such as Earth-sized planets). One shortcoming
2 Hirano et al.
of this technique is that the relative uncertainty in Is be-
comes large when Is approaches 90
◦. Another is that it
is often difficult to obtain accurate and precise measure-
ments of V sin Is for cool stars (Teff < 6000 K), for which
the rotational line broadening is often comparable to the
effects of instrumental broadening and macroturbulence.
This is contrast with measurements of the RM effect,
by which the sky-projected spin-orbit angle λ can often
be measured to within 5-10 degrees (e.g., Triaud et al.
2010; Hirano et al. 2011a). For these reasons, it may be
best to regard this technique as an efficient method for
identifying low-inclination hot stars; and for identifying
candidate low-inclination cool stars that can be followed
up with complementary techniques.
In the precursor to this paper, Hirano et al. (2012a) de-
termined stellar inclinations for 7 host stars of transiting-
planet candidates. To measure rotation periods Ps, they
used a periodogram analysis of the light curve modu-
lations seen with the Kepler telescope. They also un-
dertook new spectroscopic measurements of V sin Is and
stellar radii Rs via Is = arcsin(Ps · V sin Is/2piRs), for
several KOIs (Kepler Objects of Interest). They found
that most of the systems are consistent with Is = 90
◦,
suggesting good spin-orbit alignment, but at least one
system (KOI-261) may have a spin-orbit misalignment.
The planet Kepler-63b was also found to have a tilted
orbit using the same technique, and also through the
measurement of the sky-projected obliquity using the
RM effect (Sanchis-Ojeda et al. 2013). More recently,
Walkowicz & Basri (2013) applied the same technique
and found candidate spin-orbit misalignments for several
KOI’s including a multiple transiting system (Kepler-
9). Even more recently, a robust spin-orbit misalign-
ment around multiple systems was reported for Kepler-
56 based on the asteroseismic determination of the stellar
inclination (Huber et al. 2013).
In this paper, we continue the effort by Hirano et al.
(2012a) to examine the stellar inclinations for KOI sys-
tems. In the next section, we describe the new spectro-
scopic observations with the Subaru telescope to obtain
basic spectroscopic parameters for 25 KOI systems, in-
cluding 10 systems with multiple transiting planets. We
then present the analyses of stellar rotational periods and
spectroscopic parameters such as V sin Is and Rs in Sec-
tion 3. Section 4 presents a statistical analysis of the
observed distribution of Is. We try to test some hy-
potheses such as whether the observed values of Is are
drawn from an isotropic distribution (§4.3). Section 5
summarizes our results and their implications.
2. TARGET SELECTION AND OBSERVATIONS
We composed a list of KOIs for measurements of stel-
lar inclinations based on the following criteria: (1) a pre-
liminary light curve analysis shows a peak power in the
Lomb-Scargle periodogram larger than 1000, (2) the es-
timated rotation velocity at the stellar equator is larger
than about 3 km s−1, and (3) the apparent magnitude
in the Kepler bandpass is mKep . 14. The rotational ve-
locity needed for the second criterion was estimated from
the stellar radius in the Kepler Input Catalog (KIC) and
the preliminary estimate of the rotation period. We ex-
cluded slow rotators because the measurement of V sin Is
for slow rotators (V sin Is . 3 km s
−1) has a large frac-
tional uncertainty, as shown below.
 0
 0.05
 0.1
 0.15
 0.2
 0.25
 0.3
 0.35
-6 -4 -2  0  2  4  6
R
el
at
iv
e 
In
te
ns
ity
Velocity Shift [km s-1]
Image Slicer #1
Image Slicer #2
Fig. 1.— Instrumental profiles of Subaru/HDS for Image Slicer
#1 (blue solid line) and #2 (green dashed line). These profiles were
extracted from the same spectral region that was used to determine
the V sin Is of the program stars.
In order to estimate the basic spectroscopic param-
eters, we conducted high dispersion spectroscopy with
Subaru/HDS on 2012 June 30, July 1, 2, and September
4; and on 2013 June 20 and 21. All together we obtained
spectra for 25 KOIs. During the 2012 observations, we
employed the standard “I2a” setup with Image Slicer #1
(Tajitsu et al. 2012), attaining a spectral resolution of
R ∼ 110, 000. For the 2013 observations we used Image
Slicer #2 (R ∼ 80, 000). On each night of observations,
we obtained a spectrum of the flat-field lamp through the
iodine cell, to determine the instrumental line broadening
function. In the subsequent analysis, the line broaden-
ing due to the instrumental profile (IP) for each setup
was deconvolved as shown in Figure 1, and taken into
account when we estimated the rotation velocity of each
star.
Each spectrum was subjected to standard IRAF proce-
dures to extract a one-dimensional (1D) spectrum. The
wavelength scale was set with reference to a spectrum
of the thorium-argon lamp. The resultant signal-to-
noise ratio (SNR) in the 1D spectrum was typically 50-
100 pixel−1. The I2a setup covers the spectral region be-
tween 4900-7600 A˚, within which there is a large number
of iron lines available for estimation of the photospheric
parameters.
3. ANALYSES AND RESULTS
3.1. Estimate for Rotation Periods
We determined the rotation periods of the stars us-
ing the photometric observations provided by the Kepler
telescope (Borucki et al. 2010). In particular, we used
the Long Cadence data (30 minute integrations) avail-
able from the MAST archive from quarters 2 through
16, for up to a total of approximately 4 years of data.
Previously, Hirano et al. (2012a) used the simple aper-
ture flux data to obtain the rotation periods, but those
data needed to be treated carefully to remove systematic
and instrumental effects on timescales similar to the rota-
tion periods. In this paper we used the PDC-MAP final
data product, since it is designed to remove the unphys-
ical trends leaving the signal of stellar spots unaltered
(Smith et al. 2012; Stumpe et al. 2012).
Measurements of Stellar Inclinations for KOI Systems II 3
5 10 15 20
McQuillan rotation period [days]
5
10
15
20
O
ur
 ro
ta
tio
n 
pe
rio
d 
[da
ys
]
0 5 10 15 20
0.6
0.8
1.0
1.2
1.4
Fig. 2.— Comparison between the rotation periods listed in
Table 1 and periods estimated by McQuillan et al. (2013b).
To excise the data obtained during transits, we identi-
fied the transit intervals using the publicly available tran-
sit ephemerides (Batalha et al. 2013) downloaded from
the NASA exoplanet archive (Akeson et al. 2013), which
are based on the assumption of constant orbital peri-
ods. We also removed gross outliers, and normalized
the data from each quarter by dividing by the quarterly
median flux. We then computed the Lomb-Scargle peri-
odogram adopting the definition and algorithm described
by Press & Rybicki (1989). In general each periodogram
showed several peaks, the strongest of which can be at-
tributed to stellar variability. We selected the strongest
peak of the periodogram as the first candidate for the
rotation period, and adopted the full width at half max-
imum (FWHM) of the peak as the 1σ uncertainty. We
also performed a visual inspection of each light curve to
make sure that the stellar flux appeared to be varying
quasi-periodically with the candidate rotation period, as
opposed to a more regular periodic signal that would
be caused by orbital effects or pulsation. In particular,
we looked for quasi-sinusoidal variations with slow am-
plitude and phase modulation on a timescale of a few
rotation periods, as would be expected of starspots. We
also checked that there was not additional power at twice
the candidate rotation period, as it sometimes happens
when a star has two similar size starspots in opposite lon-
gitudes. Such a configuration causes the flux variations
to peak twice per rotation period, inducing a substitute
for a subharmonic peak at half the rotation period, which
in some occasions could be more significant than the real
rotation period peak, making our code identify the wrong
rotation period. In two cases, KOI-180 and KOI-2636,
the strongest peak corresponded to half the rotation pe-
riod, so we matched the correct peak with the rotation
period, and assigned the right uncertainty neglecting all
the power at half the rotation period. Table 1 summa-
rizes our rotation period measurements, including the
peak value of the periodogram power and the variability
amplitude, defined as the full range of flux after elimi-
nating the lowest 10% and the highest 10% of the flux
values (Hirano et al. 2012a).
McQuillan et al. (2013a) advocated the autocorre-
lation function, rather than the Lomb-Scargle peri-
TABLE 1
Rotation periods estimated from the Kepler photometry.
Also given are the peak periodogram power and the
variability amplitude (defined as in the text).
System Ps (days) Peak Power Variability Amplitude (%)
KOI-180 15.728 ± 0.726 3266.21 0.348
KOI-285 16.829 ± 0.588 1474.14 0.013
KOI-304 15.814 ± 2.606 1470.75 0.070
KOI-323 7.674± 0.143 4121.72 0.566
KOI-635 9.328± 0.558 1814.23 0.142
KOI-678 13.871 ± 0.058 14941.34 0.924
KOI-718 16.603 ± 0.828 1225.87 0.045
KOI-720 9.378± 0.027 7967.45 0.670
KOI-988 12.363 ± 0.064 15540.17 0.685
KOI-1615 7.797± 0.043 5516.86 0.256
KOI-1628 5.756± 0.378 2540.79 0.179
KOI-1779 7.154± 0.014 9393.85 0.548
KOI-1781 10.474 ± 0.084 7000.01 0.733
KOI-1797 10.826 ± 0.033 14184.33 0.679
KOI-1835 9.644± 0.028 11825.34 0.531
KOI-1839 6.252± 0.027 8156.72 0.810
KOI-1890 6.420± 0.039 2716.78 0.020
KOI-1916 10.318 ± 0.097 3128.06 0.144
KOI-2001 16.385 ± 0.081 14017.50 0.843
KOI-2002 10.708 ± 0.364 1989.29 0.138
KOI-2026 10.051 ± 0.555 2530.09 0.191
KOI-2035 7.127± 0.104 7347.71 0.741
KOI-2087 13.816 ± 1.155 1902.26 0.086
KOI-2261 11.366 ± 0.042 17009.04 0.515
KOI-2636 16.330 ± 1.317 1774.31 0.077
odogram, for measuring rotation periods with Ke-
pler data. We checked our measured rotation pe-
riods against a published table of rotation periods
that were determined using the autocorrelation function
(McQuillan et al. 2013b), and found good agreement be-
tween the results of both techniques (Figure 2), although
our quoted uncertainties are always larger.
3.2. Spectroscopic Parameters
3.2.1. Photospheric Parameters, and Stellar Radius
Based on Takeda et al. (2002, 2005), we estimated the
basic photospheric parameters (the effective temperature
Teff , surface gravity log g, microturbulent velocity ξ, and
metallicity [Fe/H]) by measuring the equivalent widths
of the available iron absorption lines. That is, these pa-
rameters are established by requiring that the following
three conditions are simultaneously fulfilled: (a) exci-
tation equilibrium (Fe abundances show no systematic
dependence on the excitation potential), (b) ionization
equilibrium (mean Fe abundance from Fe I lines and that
from Fe II lines agree with each other), and (c) curve-of-
growth matching (Fe abundances do not systematically
depend on line strengths). We used typically ∼150–200
and ∼10–15 lines for Fe I and Fe II, respectively.
We next convert the photospheric parameters into stel-
lar masses and radii employing the Yonsei-Yale (Y2)
stellar-evolutionary models (Yi et al. 2001). Since an
accurate estimation of the stellar radius is essential in
our methodology, it is important to take account of the
accuracy of the photospheric parameters. Bruntt et al.
(2010) spectroscopically analyzed 23 solar-type stars, ar-
guing that the “true” effective temperature of a star de-
fined from the stellar luminosity and radius might have
a systematic offset of −40± 20 K from the spectroscopic
model parameter Teff , while spectroscopic measurements
4 Hirano et al.
 4600
 4800
 5000
 5200
 5400
 5600
 5800
 6000
 6200
 6400
 4600  4800  5000  5200  5400  5600  5800  6000  6200
T e
ff,
sp
ec
 
[K
]
Teff,KIC [K]
Fig. 3.— Comparison between our measurement of Teff , and the
value of Teff reported in the Kepler Input Catalog (KIC).
 3.8
 3.9
 4
 4.1
 4.2
 4.3
 4.4
 4.5
 4.6
 4.7
 3.8  3.9  4  4.1  4.2  4.3  4.4  4.5  4.6  4.7
(lo
g g
) sp
ec
 
[de
x]
(log g)KIC [dex]
Fig. 4.— Comparison between our measurement of log g and the
KIC value of log g. While both quantities are in good agreement
for log g & 4.3 dex, large discrepancies are seen for smaller value of
log g.
of the surface gravity log g did not show a significant
offset. Since our spectroscopic measurement of Teff is
similar to that of Bruntt et al. (2010), we assume that
the systematic error in Teff is 40 K, which is quadrat-
ically added to the internal statistical error listed in
Table 2 when we estimate the stellar radii and masses
based with the Y2 isochrones. To account for these un-
certainties (both statistic and systematic) in the photo-
spheric parameters, we randomly generated many sets
of (Teff , log g, and [Fe/H]) assuming Gaussian distribu-
tions for their uncertainties. Each set of (Teff , log g, and
[Fe/H]) was then converted to the mass and radius on
the Y2 isochrones. The resultant distributions give the
estimates (and errors) for the mass and radius of each
system. Table 2 summarizes our measurements of the
photospheric parameters together with the stellar radius.
To check whether our spectroscopically-derived pho-
tospheric parameters are compatible with the param-
eters that were determined from broadband photome-
try, we compared our effective temperatures and sur-
face gravities with the values reported in the Kepler In-
put Catalog (KIC). Figures 3 and 4 show these com-
parisons. The root-mean-squared residual between the
spectroscopic and photometric Teff and log g are 124 K
and 0.26 dex, respectively. This level of agreement seems
reasonable given the relatively large uncertainties in the
KIC parameters (∼ 200 K for Teff and ∼ 0.4 dex for log g,
Brown et al. 2011).
3.2.2. Projected Rotational Velocity
We measured the projected rotational velocity V sin Is
by fitting a model to the observed spectrum for each sys-
tem. Theoretically, an observed stellar spectrum Iobs(λ)
can be considered as the convolution of several func-
tions:
Iobs(λ) = S(λ) ∗M(λ) ∗ IP, (1)
where S(λ) is the intrinsic stellar spectrum taking into
account only thermal and natural broadening (including
microturbulence), M(λ) is the broadening kernel repre-
senting rotation and macroturbulence (Gray 2005), and
IP represents the instrumental line profile (see Figure
1). The IP was determined by deconvolving the spec-
trum of the flat-field lamp through the iodine cell. For
each target star, we generated the intrinsic spectrum
S(λ) based on the ATLAS9 model (a plane-parallel stel-
lar atmosphere model in LTE, Kurucz 1993) with the
input photospheric parameters being the best-fit values
derived above, and fitted the observed spectrum Iobs(λ),
allowing V sin Is to be a free parameter (which affects
M(λ)). As for the macroturbulence, we adopted the
radial-tangential model of Gray (2005) and assumed that
the macroturbulent velocity ζRT is expressed by the fol-
lowing empirical formula (Valenti & Fischer 2005):
ζRT =
(
3.98 +
Teff − 5770 K
650 K
)
km s−1. (2)
This empirical formula was derived based on the sta-
tistical distribution of the upper limit of ζRT, in which
V sin Is = 0 km s
−1 was assumed in fitting the spectral
lines for a large number of stars in the controlled sam-
ple (the SPOCS catalog). Taking the “lower” bound-
ary of the upper limit of ζRT as a function of Teff ,
Valenti & Fischer (2005) derived Equation (2) (see Fig-
ure 3 in Valenti & Fischer 2005). In the subsequent anal-
ysis we assumed that the uncertainty in ζRT is ±15% for
cool stars (Teff ≤ 6100 K) based on the observed disper-
sion of the upper limit of ζRT around Equation (2). But
for hot stars (>6100 K) for which the SPOCS catalog
has a relatively small number of stars, we conservatively
adopted ±25% for the systematic uncertainty in ζRT.
3.2.3. Correction for the Impact of Differential Rotation
The Sun’s rotation period varies with surface latitude;
the rotation rate at the Sun’s equator is faster than that
of the polar region by about 20%. It is natural to assume
that differential rotation is a feature of all our program
stars, and therefore that differential rotation needs to be
taken into account in our analysis.
As pointed out by Hirano et al. (2012a), there are two
main issues that arise because of differential rotation.
The first issue is that we do not know the latitude of the
spots that are producing the detectable photometric vari-
ations. Starspots are probably not randomly distributed;
Measurements of Stellar Inclinations for KOI Systems II 5
TABLE 2
Spectroscopic Parameters. Starred systems are multiple transiting systems. We show sin Is ≡ V sin Is/Veq in the rightmost
column based on the values of Veq and V sin Is. The listed errors in Teff represent the internal statistical error and do
not include the systematic error (see Section 3.2.1).
System Teff (K) log g [Fe/H] Ms (M⊙) Rs (R⊙) V sin Is (km s
−1) Veq (km s−1) sin Is
KOI-180 5592 ± 40 4.389 ± 0.100 0.12± 0.05 0.992+0.027
−0.022 1.029
+0.077
−0.088 3.15± 0.81 3.30
+0.45
−0.31 0.941
+0.276
−0.255
KOI-285 5962 ± 27 3.997 ± 0.060 0.16± 0.03 1.324+0.047
−0.051 1.914
+0.134
−0.145 4.21± 0.77 5.74
+1.28
−0.59 0.708
+0.180
−0.162
KOI-304⋆ 5777 ± 42 4.399 ± 0.095 −0.14± 0.05 0.962+0.026
−0.022 1.009
+0.068
−0.086 1.62± 1.28 3.22
+0.84
−0.53 0.484
+0.425
−0.383
KOI-323 5418 ± 30 4.558 ± 0.075 0.01± 0.04 0.927+0.021
−0.026 0.841
+0.054
−0.030 4.70± 0.30 5.56
+0.39
−0.24 0.838
+0.072
−0.071
KOI-635 6194 ± 52 4.493 ± 0.100 0.25± 0.07 1.260+0.026
−0.031 1.173
+0.050
−0.040 8.82± 0.52 6.39
+1.08
−0.54 1.360
+0.172
−0.194
KOI-678⋆ 5129 ± 32 4.532 ± 0.085 0.19± 0.04 0.886+0.023
−0.018 0.833
+0.031
−0.046 3.21± 0.45 3.04
+0.16
−0.18 1.058
+0.164
−0.157
KOI-718⋆ 6002 ± 42 4.577 ± 0.080 0.58± 0.04 1.215+0.028
−0.022 1.133
+0.049
−0.044 2.53± 1.26 3.45
+0.80
−0.31 0.697
+0.382
−0.351
KOI-720⋆ 5198 ± 40 4.580 ± 0.100 0.01± 0.05 0.862+0.023
−0.020 0.789
+0.040
−0.038 4.18± 0.30 4.25
+0.25
−0.20 0.980
+0.091
−0.087
KOI-988⋆ 5114 ± 45 4.544 ± 0.115 0.10± 0.04 0.861+0.022
−0.019 0.802
+0.032
−0.045 2.64± 0.57 3.28
+0.17
−0.19 0.808
+0.181
−0.177
KOI-1615 5934 ± 35 4.266 ± 0.080 0.21± 0.04 1.181+0.055
−0.038 1.321
+0.161
−0.135 8.74± 0.21 8.57
+1.24
−0.91 1.017
+0.127
−0.128
KOI-1628 6125 ± 40 4.274 ± 0.085 0.13± 0.04 1.218+0.050
−0.033 1.328
+0.168
−0.137 11.24 ± 0.27 11.71
+1.90
−1.42 0.957
+0.137
−0.133
KOI-1779⋆ 5781 ± 47 4.442 ± 0.105 0.33± 0.07 1.124+0.027
−0.024 1.058
+0.133
−0.049 7.41± 0.24 7.48
+1.03
−0.38 0.979
+0.073
−0.112
KOI-1781⋆ 4864 ± 55 4.478 ± 0.145 0.19± 0.06 0.815+0.020
−0.018 0.760
+0.028
−0.034 3.64± 0.22 3.67
+0.15
−0.17 0.994
+0.077
−0.072
KOI-1797 4934 ± 42 4.430 ± 0.115 0.16± 0.06 0.824+0.017
−0.015 0.781
+0.025
−0.028 3.69± 0.23 3.65
+0.14
−0.13 1.011
+0.075
−0.072
KOI-1835⋆ 5046 ± 70 4.313 ± 0.190 0.16± 0.07 0.862+0.264
−0.026 0.832
+1.216
−0.041 4.66± 0.20 4.36
+6.39
−0.21 1.012
+0.126
−0.580
KOI-1839 5465 ± 40 4.485 ± 0.100 0.05± 0.06 0.938+0.028
−0.022 0.908
+0.051
−0.062 7.41± 0.15 7.35
+0.46
−0.51 1.010
+0.077
−0.064
KOI-1890 6107 ± 40 3.971 ± 0.085 0.22± 0.05 1.477+0.094
−0.084 2.078
+0.285
−0.247 7.44± 0.49 16.38
+2.61
−2.00 0.452
+0.073
−0.066
KOI-1916⋆ 5945 ± 25 4.308 ± 0.060 0.31± 0.04 1.193+0.034
−0.025 1.265
+0.109
−0.094 6.38± 0.39 6.20
+0.86
−0.52 1.017
+0.125
−0.128
KOI-2001 5144 ± 30 4.484 ± 0.085 0.01± 0.04 0.839+0.016
−0.014 0.804
+0.021
−0.029 2.44± 0.66 2.48
+0.13
−0.10 0.980
+0.272
−0.268
KOI-2002 5963 ± 50 4.071 ± 0.110 0.15± 0.05 1.342+0.083
−0.080 1.776
+0.273
−0.263 5.67± 0.42 8.39
+1.64
−1.30 0.671
+0.140
−0.116
KOI-2026 5919 ± 47 4.178 ± 0.100 0.01± 0.05 1.106+0.060
−0.041 1.419
+0.200
−0.175 4.76± 0.50 7.15
+1.32
−0.98 0.659
+0.136
−0.116
KOI-2035 5484 ± 25 4.544 ± 0.060 0.13± 0.04 0.984+0.021
−0.023 0.885
+0.047
−0.024 6.35± 0.21 6.30
+0.39
−0.22 1.002
+0.055
−0.061
KOI-2087 5955 ± 25 4.364 ± 0.060 0.02± 0.03 1.084+0.024
−0.020 1.129
+0.087
−0.075 4.46± 0.73 4.15
+0.83
−0.47 1.048
+0.250
−0.221
KOI-2261⋆ 5154 ± 32 4.515 ± 0.080 0.12± 0.05 0.873+0.021
−0.016 0.830
+0.027
−0.042 2.81± 0.55 3.69
+0.17
−0.19 0.764
+0.156
−0.153
KOI-2636 5876 ± 35 4.337 ± 0.085 0.16± 0.04 1.118+0.037
−0.027 1.181
+0.145
−0.114 2.40± 1.26 3.67
+0.84
−0.47 0.631
+0.372
−0.336
they are likely to be concentrated around particular lati-
tudes. On the Sun, the “active latitudes” gradually vary
from about ±40◦ down to the equator, over the 11-year
solar cycle. Therefore, we need to take account the sys-
tematic errors due to the imperfect knowledge of the
spots’ locations. The second issue is the distortion in
the spectral line shape caused by differential rotation.
The absorption lines of a Sun-like star are narrower than
would be expected for a star with no differential rota-
tion, because differential rotation reduces the weight of
the extremes in rotation velocity. Therefore, an analy-
sis of spectral lines that neglects differential rotation will
give a value of V sin Is that is systematically smaller than
the true equatorial projected rotation velocity.
We corrected for the first of these two issues using the
procedure described by Hirano et al. (2012a). Employing
the empirical relation given by Collier Cameron (2007)
for the magnitude of differential rotation, we express the
rotation rate Ω as a function of the latitude l on the
stellar surface:
Ω(l) = Ωeq(1− α sin
2 l), (3)
where Ωeq is the angular rotation velocity at the equator,
and
αΩeq = 0.053
(
Teff
5130 K
)8.6
rad day−1. (4)
Assuming that the observed rotation rates are due to
spots located at the stellar latitude l = 20◦ ± 20◦ (as
is the case for the Sun), we re-estimated the equatorial
rotation velocity (Veq) for each of the targets as
Veq =
2piRs
Ps
1
1− α sin2 20◦
, (5)
and added in quadrature the following lower and upper
systematic errors in Veq:
(∆Veq)low,sys =Veq
(
1
1− α sin2 20◦
− 1
)
, (6)
(∆Veq)upp,sys=Veq
(
1
1− α sin2 40◦
−
1
1− α sin2 20◦
)
.(7)
Table 2 gives the resulting estimates of the equatorial ro-
tation velocities. For reference, the assumed magnitude
of differential rotation was on average α ≃ 0.23 for the
targets listed in Table 2, which is nearly the same as that
of the Sun.
Regarding the second issue, the bias in the V sin Is
measurement, we performed a correction using the fol-
lowing procedure. First, we computed sin Is for each
target based on the preliminary measurements of V sin Is
and Veq (before any correction to V sin Is for differen-
tial rotation). A simulated line profile was then gener-
ated, using the model of Equation (1). In this case M(λ)
corresponds to the macroturbulence-plus-rotation kernel
in the presence of differential rotation using sin Is, Veq,
and α as input parameters. We adopted plausible val-
ues for the other spectroscopic parameters (i.e., the in-
trinsic Gaussian and Lorentzian dispersions, macrotur-
bulence, limb-darkening, and IP) in making the mock
profile. This mock line was then fitted assuming zero
differential rotation, with V sin Is as the only free pa-
rameter. After computing the ratio f of the resultant
best-fitting V sin Is to the product of the input Veq and
sin Is, we divided the originally measured V sin Is by the
ratio f to obtain the final V sin Is corrected for the im-
pact of differential rotation. We note that f ≈ 1 − α/2
6 Hirano et al.
was in general obtained, indicating that the measured
V sin Is is always underestimated when a rigid rotation
is assumed in fitting the spectrum (also see Figure 11
in Hirano et al. 2012a). The resultant V sin Is after the
correction of differential rotation for each system is also
summarized in Table 2.
Some of our program stars were also studied by
Walkowicz & Basri (2013), giving us the opportunity to
check on the agreement. For the stars KOI-180, 323,
and 988, respectively, Walkowicz & Basri (2013) found
V sin Is = 2.7±0.5 km s
−1, 3.3±0.5 km s−1, and 2.7±0.5
km s−1. Comparing these with the values in Table 2,
KOI-180 and 988 show a good agreement between two
measurements, but KOI-323 shows a ∼ 3σ level disagree-
ment. Furthermore, for KOI-261, Walkowicz & Basri
(2013) found V sin Is = 2.3 ± 0.5 km s
−1, which is in
agreement with the 1σ upper limit of 2.57 km s−1 deter-
mined by Hirano et al. (2012a) using the same technique
as applied here.
4. DISCUSSION
4.1. Evidence of Spin-orbit Misalignment
Figure 5 plots V sin Is against Veq, after making the
corrections for differential rotation. Single transiting sys-
tems are shown in panel (a), and systems with multiple
transiting candidates are shown in panel (b). The black
solid line represents Is = 90
◦. Systems falling on this line
would have the stellar spin oriented perpendicular to the
line-of-sight, and therefore likely aligned with the plane-
tary orbital axes (although an unlikely possibility is that
they are misaligned with the line of nodes coincidentally
along the line of sight). The dashed lines show different
degrees of misalignment (Is = 45
◦ and Is = 30
◦).
Most of the data points in Figure 5 do indeed fall near
the Is = 90
◦ line, indicating a tendency toward spin-
orbit alignment. Four of the systems—KOI-323, 1890,
2002, and 2026—show evidence for significant spin-orbit
misalignments with more than 2σ confidence. All four of
these systems are single-transiting candidates. Some of
the multiple-transiting candidates also show evidence for
misalignment but only at the 1σ level; these are KOI-304,
988, and 2261.
As the multiple-transiting systems are of special im-
portance, it is worth focusing on those possible misalign-
ments and check if the results for the rotation period,
stellar radius, and V sin Is are robust. A spurious find-
ing of misalignment can result from an underestimate of
either V sin Is or P , or an overestimate of Rs.
First, we check on the rotation periods. The relevant
light curves and periodograms are shown in Figure 6.
Each light curve shown in Figure 6 shows an evident
pattern of quasi-periodic flux variation, and the peri-
odograms for KOI-988 and KOI-2261 exhibit a clear and
unambiguous peak that surpasses a power of 104. For the
case of KOI-304, on the other hand, there are multiple,
relatively weak peaks of comparable power. These mul-
tiple peaks could be ascribed to differential rotation or
rapid starspot evolution. Nevertheless, visual inspection
of the light curves does not reveal any problem with the
quoted rotation periods of 15.8± 2.6 days for KOI-304.
Next, we check on the determination of the stellar ra-
dius. We have already shown in Section 3.2.1 that the
photospheric parameters (Teff and log g) are in reason-
ably good agreement with the KIC values. Here we fo-
0
5
10
15
20
0 5 10 15 20
V
s
in
I s
[k
m
s
-1
]
Veq [km s
-1
]
(a) single
Is=90
o
Is=45
o
Is=30
o 18902002
2026
285
323
0
2
4
6
8
10
0 2 4 6 8 10
V
s
in
I s
[k
m
s
-1
]
Veq [km s
-1
]
(b) multi Is=90
o
Is=45
o
Is=30
o
2261
988
304
Fig. 5.— Projected rotational velocity (V sin Is) as a function
of the stellar rotation velocity at the equator (Veq), for (a) single
and (b) multiple KOI systems. We plot here the newly observed
25 KOI systems. The solid lines indicates Is = 90◦ while the
dashed lines represent different degrees of misalignment (Is = 30◦
and Is = 45◦). In the lower panel, the data point with the very
large upper uncertainty in Veq is KOI-1835, which has a poorly
determined surface gravity (and thus a poorly determined stellar
radius). Note that the panels (a) and (b) show different ranges of
Veq.
cus on the estimate of stellar mass and radius, based on
the Y2 isochrones. Figure 7 shows the placement of the
measured values of Teff and log g (red crosses) on the
theoretical isochrones (blue dashed lines) and the loci of
equal stellar radius (black solid lines) of the Y2 theoret-
ical evolutionary models for main-sequence stars. The
measured values of Teff and log g for KOI-304, 988, and
2261 conform with the models.
Finally, we check on the measurements of V sin Is
based on the observed line broadening in the Subaru
spectra. Figure 8 shows part of the observed spectrum
(blue dots) along with the best-fitting model spectrum
(red line) for each of (a) KOI-304, (b) KOI-988, and (c)
KOI-2261. For reference, the green area shows the spec-
tral lines that would be expected for sin Is = 1 (i.e., spin-
orbit alignment). The breadth of the green area arises
from variation of the macroturbulent velocity ζRT by
±15% from the value computed by Equation (2). A mis-
aligned system will show narrower lines than the green
region. This figure illustrates the main difficulty of this
probe of spin-orbit alignment: one must isolate the very
Measurements of Stellar Inclinations for KOI Systems II 7
KOI 304
510 520 530 540 550
BJD−2454900 [days]
0.997
0.998
0.999
1.000
1.001
1.002
R
e
la
ti
v
e
fl
u
x
0 5 10 15 20
Period of rotation [days]
0
500
1000
1500
2000
P
o
w
e
r
KOI 988
210 220 230 240
BJD−2454900 [days]
0.985
0.990
0.995
1.000
1.005
1.010
1.015
R
e
la
ti
v
e
fl
u
x
0 5 10 15 20
Period of rotation [days]
0
5.0•103
1.0•104
1.5•104
2.0•104
P
o
w
e
r
KOI 2261
340 350 360 370 380
BJD−2454900 [days]
0.99
1.00
1.01
R
e
la
ti
v
e
fl
u
x
0 5 10 15 20
Period of rotation [days]
0
5.0•103
1.0•104
1.5•104
2.0•104
2.5•104
P
o
w
e
r
Fig. 6.— Light curves and Lomb-Scargle periodograms for multi-
ple systems showing a possible spin-orbit misalignment (KOI-304,
988, 2261). In the periodograms, the intervals surrounded by the
two blue lines correspond to the rotation periods and their uncer-
tainties (§3.1).
small differences in line broadening due to rotation as
opposed to macroturbulence and instrumental broaden-
ing. For all the three systems shown here, V sin Is cannot
be much larger than the values listed in Table 2 (see in
particular the bottom of each absorption line), unless the
assumed macroturbulent velocity is in error by more than
15%. It is important to remember that Figure 8 shows
 4
 4.1
 4.2
 4.3
 4.4
 4.5
 4.6
 4.7
 5000  5500  6000  6500
lo
g 
g 
[de
x]
Teff [K]
(a) KOI-304
[Fe/H] = -0.14
0.8 MSun
0.9 MSun
1.0 MSun
1.1 MSun 1.2 MSun
0.8 RSun
0.9 RSun
1.0 RSun
1.1 RSun
1.2 RSun
 4
 4.1
 4.2
 4.3
 4.4
 4.5
 4.6
 4.7
 5000  5500  6000  6500
lo
g 
g 
[de
x]
Teff [K]
(b) KOI-988
[Fe/H] = +0.10
0.8 MSun
0.9 MSun
1.0 MSun
1.1 MSun 1.2 MSun
0.8 RSun
0.9 RSun
1.0 RSun
1.1 RSun
1.2 RSun
 4
 4.1
 4.2
 4.3
 4.4
 4.5
 4.6
 4.7
 5000  5500  6000  6500
lo
g 
g 
[de
x]
Teff [K]
(c) KOI-2261
[Fe/H] = +0.12
0.8 MSun
0.9 MSun
1.0 MSun
1.1 MSun 1.2 MSun
0.8 RSun
0.9 RSun
1.0 RSun
1.1 RSun
1.2 RSun
Fig. 7.— Placements of measured Teff and log g in the Y
2
isochrone for (a) KOI-304, (b) KOI-988, and (c) KOI-2261 (red
crosses with error bars). The blue dashed lines indicate the evolu-
tional tracks and the black solid lines are the “iso-radius”, based
on the Y2 isochrone. Note that the errorbars in Teff are enlarged
so that they contain the systematic error of 40 K.
only a part of the observed spectrum. The true statisti-
cal significance of the results is higher than it might seem
visually because V sin Is was determined from data over
a wider range of wavelengths.
In summary, the detailed visual inspection of the multi-
transiting systems with possible misalignments did not
raise any specific concerns for all of KOI-304, 988, and
2261, which remain viable candidates for multi-planet
8 Hirano et al.
 0.5
 0.6
 0.7
 0.8
 0.9
 1
 1.1
 1.2
 608.5  608.52  608.54  608.56  608.58  608.6  608.62  608.64  608.66  608.68
N
or
m
al
iz
ed
 F
lu
x
Wavelength [nm]
(a) KOI-304 VsinIs  = Veqbest-fit
 0.4
 0.5
 0.6
 0.7
 0.8
 0.9
 1
 1.1
 1.2
 608.5  608.52  608.54  608.56  608.58  608.6  608.62  608.64  608.66  608.68
N
or
m
al
iz
ed
 F
lu
x
Wavelength [nm]
(b) KOI-988 VsinIs  = Veqbest-fit
 0.5
 0.6
 0.7
 0.8
 0.9
 1
 1.1
 1.2
 608.5  608.52  608.54  608.56  608.58  608.6  608.62  608.64  608.66  608.68
N
or
m
al
iz
ed
 F
lu
x
Wavelength [nm]
(c) KOI-2261 VsinIs  = Veqbest-fit
Fig. 8.— Part of the spectrum used for fitting V sin Is for (a)
KOI-304, (b) KOI-988, and (c) KOI-2261. In each panel, the blue
dots show the observed spectrum and the red solid line indicates
the best-fitting model computed by Equation (1). The green area
indicates the range of model spectra satisfying sin Is = 1 (i.e., spin-
orbit alignment) with a macroturbulent velocity ζRT differing by
±15% from the assumed value.
systems with misaligned stars. Each individual detection
is statistically marginal, with less than 2σ confidence,
but if the uncertainties have been accurately determined,
then together it is likely that at least one system is mis-
aligned. To quantify this statement we can compute
the probability that all three multiple systems (KOI-304,
988, 2261) are well-aligned, defining this for convenience
to mean Is ≥ 75
◦ (sin Is ≥ 0.9659). Assuming that both
Veq and V sin Is have uncertainties drawn from indepen-
dent Gaussian distributions, with dispersions set equal
to our quoted uncertainties, we calculate the probabil-
ity for each system to have 0.9659 ≤ sin Is. If the lower
and upper observation errors are different, we adopt a
two-sided Gaussian with different upper and lower dis-
persions. We then compute the products of the resulting
probabilities to find the net probability pall aligned that all
three systems are aligned. We find pall aligned = 0.0025,
implying that at least one system among the three KOI’s
is very likely to have spin-orbit misalignment. This re-
sult cannot be definitive, though, given the possibility
of systematic effects, or uncertainties that are correlated
between different systems due to shared assumptions and
techniques. Specifically, all the three systems fall on the
regime where the measurement of V sin Is tends to suffer
from systematic effects (V sin Is . 3 km s
−1). It is bet-
ter to regard KOI-304, 988, and 2261 as candidate mis-
alignments that are good targets for additional follow-up
observations.
4.2. Distribution of Stellar Inclinations
In the previous subsection, we have seen that some of
the systems (both single and multiple) may have spin-
orbit misalignments. A natural question is “what is the
fraction of misaligned systems?” Although the number
of our samples is still small, a histogram of the observed
Is may be helpful to gain an insight into the underlying
true distribution of the spin-orbit angle, just as the his-
togram of sky-plane angles was useful in the case of RM
measurements (e.g., Pont et al. 2010). One issue con-
cerning the conversion from the observed Veq and V sin Is
to the distribution of Is is that sin Is ≡ V sin Is/Veq
could extend beyond unity due to measurement uncer-
tainties. Theoretical distributions of sin Is always satisfy
0 ≤ sin Is ≤ 1, which inhibits a direct comparison be-
tween the theoretical and observed distributions. Here,
we present a Bayesian method that avoids this problem
by placing a prior on Is.
Based on Bayes’ theorem, the posterior probability dis-
tribution of Veq and Is is
P (Veq, Is|D) ∝ P (D|Veq, Is) · pprior(Veq) · pprior(Is),(8)
where “D” represents the observed data for Veq and
V sin Is for each of the observed systems. We again as-
sume that observational data for Veq and V sin Is fol-
low the Gaussian distributions with their centers being
Veq = p
(i) and V sin Is = q
(i), and dispersions being σ
(i)
p
and σ
(i)
q , where i is the label of the system. In this case,
the conditional probability P (D|Veq, Is) is expressed as
P (D|Veq, Is) ∝
1
σ
(i)
p
exp
{
−
(p(i) − Veq)
2
2σ
(i)2
p
}
1
σ
(i)
q
exp
{
−
(q(i) − Veq sin Is)
2
2σ
(i)2
q
}
. (9)
Assuming a uniform distribution for pprior(Veq) (0 ≤
Veq), we marginalize Veq, so that we obtain the poste-
Measurements of Stellar Inclinations for KOI Systems II 9
 0
 0.5
 1
 1.5
 2
 2.5
 0  0.2  0.4  0.6  0.8  1  1.2  1.4
Pr
ob
ab
ilit
y 
De
ns
ity
Is [radian]
blue single
red multi
Fig. 9.— Posterior distributions of Is computed by Equation
(10). The dashed lines correspond to the result for each system in
the sample (light-blue for single and light-red for multiple systems).
The solid lines are averaged shapes of the posteriors for each cate-
gory. The black solid curve represents the isotropic distribution of
Is for reference.
rior distribution for Is:
Pi(Is|D) ∝
∫
∞
0
1
σ
(i)
p
exp
{
−
(p(i) − Veq)
2
2σ
(i)2
p
}
1
σ
(i)
q
exp
{
−
(q(i) − Veq sin Is)
2
2σ
(i)2
q
}
dVeq · pprior(Is). (10)
In case that the observed result for Veq or V sin Is has
different upper and lower errors, we adopt two-sided
Gaussian functions as in §4.1. When a prior defined
in 0 ≤ Is ≤ pi/2 (i.e., 0 ≤ sin Is ≤ 1) is applied, the
posterior Pi(Is|D) also could have a non-zero value in
0 ≤ Is ≤ pi/2.
For each of the observed KOI systems, we compute
the posterior distribution by Equation (10). We here as-
sume the isotropic distribution for the prior on Is (i.e.,
pprior(Is) = sin Is). This is physically unlikely consider-
ing the fact that many systems show a good spin-orbit
alignment from measurements of the RM effect. How-
ever, the prior distribution is not so important since we
do not attempt to quantitatively compare any distribu-
tions here (see the next subsection for a quantitative
comparison). Instead, in order to visualize the distri-
bution of Is, we take the average of the posterior distri-
butions by stacking Pi(Is|D) for observed systems. In
Figure 9, we plot the averaged posterior distribution for
either of single (blue) and multiple (red) KOI systems by
the solid lines. These plots correspond to a sort of his-
togram of Is considering that the peak of the posterior
for each system likely represents the most plausible value
of Is, and all the systems have an equal weight. The two
distributions (single and multiple) are similar, but single
systems show a slightly wider distribution than that of
multiple systems with small bumps at Is . 1.0 radian
(≃ 57◦). For reference, we show by the black solid line
the isotropic distribution of Is. Note that in this analy-
sis (and other statistical analyses below), we added the
seven KOI systems reported in our previous campaign
(KOI-257, 261, 262, 269, 280, 367, 974, Hirano et al.
2012a) to the list of targets subjected to the statistical
analysis.
It should be stressed that our observed systems (both
single and multiple) have no hot Jupiters and all the
planet candidates are Earth-sized or Neptune-sized ones.
Little is known about the spin-orbit angle for these
classes of planets, and the observed distribution of the
angle could be more or less different from that for close-in
giant planets. We also note that while hot Jupiters are in
general isolated single planets (Steffen et al. 2012), many
of single transiting systems in our sample may actually
be multiple systems (e.g., transits of outer planets are
unobservable due to geometry). This possibility makes
it difficult to interpret the comparison of Is for single and
multiple systems.
4.3. Statistical Tests
Figures 5 and 9 suggest that the observed distributions
of Is differ from an isotropic distribution and also differ
from perfect spin-orbit alignment. However, the degree
to which the observed distributions are different from or
similar to each other is quantitatively not clear. Also, we
are interested in whether single-transiting and multiple-
transiting systems have the same distribution for Is. We
test the following two hypotheses with the Kolmogorov-
Smirnov (KS) test:
(a) the observed values of Is (all systems) are drawn
from an isotropic distribution,
(b) the observed distributions of Is for single and mul-
tiple systems are the same.
We perform a Monte-Carlo simulation to implement
the KS tests. We take the following steps based on the
observed values of Veq and V sin Is.
1. First, we randomly generate (V
(i)
eq , V sin I
(i)
s ) for
system i assuming Gaussian distributions (two-
sided Gaussians if needed) with dispersions set
equal to the quoted measurement uncertainties.
2. We then compute I
(i)
s ≡ arcsin(V sin I
(i)
s /V
(i)
eq ) for
each system. Whenever V sin I
(i)
s > V
(i)
eq , we set
Is = 90
◦.
3. Based on the set of {I
(i)
s } with all the systems in
the sample, we implement the KS test and record
the value of D (the largest difference between the
two cumulative distributions).
4. We repeat the preceding steps (1 to 3) 106 times,
recording the values of D and finding the median
and standard deviation of the collection of D val-
ues, and the corresponding probability that the two
distributions may be the same, which we denote by
p(D > Dobs).
We first test the hypothesis (a). The two distribu-
tions tested are the observed distribution of {I
(i)
s } and
the theoretical isotropic distribution. As a result of im-
plementing the steps 1 - 4, we obtain Dobs = 0.344
+0.094
−0.063,
corresponding to p(D > Dobs) = 0.00071
+0.00912
−0.00071. There-
fore, the hypothesis (a) is highly unlikely, and this result
should indicate that the stellar equators in our sample
10 Hirano et al.
are preferentially edge-on, suggesting a tendency toward
spin-orbit alignment.
In the second test, the two tested distributions are
both observed distributions of Is, one for single and
the other for multiple systems. Implementing the KS
test, we find Dobs = 0.255
+0.091
−0.065, which corresponds to
p(D > Dobs) = 0.665
+0.265
−0.380. This result indicates that
the two observed distributions are not significantly dif-
ferent, and might be drawn from the same distribution.
To see this result is robust, we repeat the steps 1-4, imple-
menting instead of the KS test the K-sample Anderson-
Darling (AD) test (e.g., Hou et al. 2009), which has
more sensitivity around the tails of distributions. As
a consequence, the p-value of 0.519+0.123
−0.277 is obtained,
which also implies the two observed distributions are
not significantly different. The results of these statistical
tests cannot corroborate recent findings by RM measure-
ments, asteroseismology, and the spot-crossing method
that multiple-transiting systems preferentially show a
good spin-orbit alignment (Sanchis-Ojeda et al. 2012;
Hirano et al. 2012b; Albrecht et al. 2013; Chaplin et al.
2013). However, it is premature to conclude that our
result actually contradicts the previous findings; more
systems are needed (particularly multiple-transiting sys-
tems) for a more definitive conclusion. The exact sample
size that will be required depends on the true distribution
of the spin-orbit angle.
5. SUMMARY
In this paper, we investigated the stellar inclinations
for KOI systems by combining the rotation periods es-
timated from the Kepler photometry and projected ro-
tational velocities V sin Is determined from Subaru spec-
troscopy. We constrained the stellar inclination Is for 25
KOI systems, and discussed statistical properties using
all the systems observed so far by Subaru. There are
several implications that we list here.
1. Based on the KS test, the observed distribution of
Is is significantly different from an isotropic distri-
bution, suggesting that the direction of stellar spin
is correlated with the planetary orbital axis. Spin-
orbit alignment has been reported for many tran-
siting systems, but most of the systems with RM
measurements have hot (warm) Jupiters. Our mea-
surements pertain to Neptune-sized or Earth-sized
planets, which are likely to have a different history
of formation and migration than giant planets. In
particular the smaller planets are not as likely to
have strong tidal interactions with their host stars,
and therefore the orbital orientations may reflect
more primordial conditions.
2. A certain fraction of the systems show possible
spin-orbit misalignments (Is . 75
◦). We had a
closer look at the seemingly misaligned multiple
transiting systems (KOI-304, 988, and 2261), and
they all survived as candidates for misaligned stars.
3. The statistical tests indicate that the observed dis-
tributions of Is for single and multiple transiting
systems are not significantly different. The sensi-
tivity of this test is limited, however, by the small
number of multiple systems (only 11). The aver-
aged posterior distribution shown in Figure 9 sug-
gests that the single transiting systems might have
a larger fraction of spin-orbit misalignment. This
should be confirmed or refuted by further observa-
tions of transiting systems.
As Hirano et al. (2012a) noted, our present method
cannot discriminate the state of Is = −90
◦ (retrograde
orbit) from that of Is = +90
◦ (prograde orbit). The
degeneracy between Is = −90
◦ and Is = +90
◦ would
certainly make the fraction of misaligned systems look
smaller than the real fraction (systems with Is ≈ −90
◦
would appear to be aligned in Figure 5), but it does not
affect the statements 1. and 2. of the above summary.
In addition, given the fact that the measurements of the
RM effect so far have not revealed a strong evidence of a
“perfectly anti-aligned” system (i.e., λ ≈ ±180◦), it is ex-
pected to be a rare case to find a system with Is ≈ −90
◦.
All the other retrograde cases (e.g., −75◦ . Is . 0
◦) are
actually regarded as ”misaligned” in Figure 5 as in the
case of prograde orbits. In other words, our methodol-
ogy gives the lower limit on the fraction of misaligned
systems.
One task left is the confirmation of the planetary na-
ture for the KOI planet candidates on which we focused
in this paper. While the false positive rate for KOI mul-
tiple systems is proved to be negligible (Lissauer et al.
2012), any contamination from background/foreground
source(s) leads to a wrong determination of the rota-
tion period and/or spectroscopic parameters. A deep di-
rect imaging search for companions around the KOI stars
would be helpful both in terms of putting a constraint
on the magnitude of contamination and identifying the
possible cause of spin-orbit misalignment.
This paper is based on data collected at Subaru Tele-
scope, which is operated by the National Astronomical
Observatory of Japan. We acknowledge the support for
our Subaru HDS observations by Akito Tajitsu, a sup-
port scientist for the Subaru HDS. T.H. expresses spe-
cial thanks to Masayuki Kuzuhara, Yuka Fujii, Akihiko
Fukui, and Yasushi Suto for fruitful discussions on this
subject. The data analysis was in part carried out on
common use data analysis computer system at the As-
tronomy Data Center, ADC, of the National Astronom-
ical Observatory of Japan. T.H. and Y.H.T. are sup-
ported by Japan Society for Promotion of Science (JSPS)
Fellowship for Research (PD:25-3183, DC1: 23-3491).
J.N.W. and R.S.O. gratefully acknowledge support from
the NASA Origins program (NNX11AG85G) and Ke-
pler Participating Scientist program (NNX12AC76G).
N.N. acknowledges support by the NAOJ Fellowship, the
NINS Program for Cross-Disciplinary Study, and Grant-
in-Aid for Scientific Research (A) (No. 25247026) from
the Ministry of Education, Culture, Sports, Science and
Technology (MEXT) of Japan. We acknowledge the very
significant cultural role and reverence that the summit of
Mauna Kea has always had within the indigenous people
in Hawai’i.
Measurements of Stellar Inclinations for KOI Systems II 11
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