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系外惑星の大気観測とHZ地球型惑星への展開
特に、中間赤外領域の可能性について

藤井友香@NAOJ
CPS seminar / 2020年10月22日(木)目次‣ 系外惑星大気へのさまざまなアプローチ：ホットジュピターHD209458bを例に‣ ホットジュピター大気データの蓄積と見え始めた統計的な性質‣ より小さな惑星、よりcoolな惑星への展開‣ 中間赤外線高分散分光を用いたHZ地球型惑星大気観測の検討

系外惑星大気へのさまざまなアプローチ(ホットジュピターHD209458bを例に)• 初めて見つかったトランジット・ホットジュピター• 0.7 MJ, 1.9 RJ, 0.05 AU, Teﬀ~ 1130 K (主星: G0型, 49 pc)• 大気についてもっともよく調べられている系外惑星の一つ

designated these time series nNa(t), mNa(t), and wNa(t) above.
Image Credit:NASAhttps://exoplanets.nasa.gov/resources/297/a-planets-transmission-spectrum/
We wishto compare these results with models of the planetary atmosphere. Several steps are required to transform themodel predictions into the same quantity as the observable,
and we describe these steps here.
We first produce several model calculations of the changep), mNa (middle), and wNa (bottom)
in the effective planetary radius as a function of wavelength,
enter of transit. The means of theas prescribed by Brown (2001). For all the models, we use
significantly offset below 0.

系外惑星大気観測の幕開け

ot these results in Figure 4.
e have observed a deeper

y that the observed decrehe STIS CCD. Because the32 " 10!4 for a change insit to out-of-transit of
1% across this range wouldoij, & Kimble (1999) derive
magnitude lower than this.
y by selecting 18 strong stelbserved wavelength range.
fine a set of bandpasses withhose listed in Table 1, but
e. We derive Dn and Dm for. The results are shown inn of either Dn or Dm withlso confirm our evaluation
uation (5). We find that they distributed, with a meanrd deviation as we found

ORETICAL PREDICTIONS

ntity that we observe is theepth in a band centered on

Na

透過光分光

トランジット中心からの時間 [days]

Fig. 4.—Top: Unbinned time series nNa (Fig. 2, top panel ). Bottom: These
data binned in time (each point is the median value in each bin). There are
10 bins, with roughly equal numbers of observations per bin (42). The error
bars indicate the estimated standard deviation of the median. The solid
curve is a model for the difference of two transit curves (described in x 3),
scaled to the observed offset in the mean during transit,
DnNa = !2.32 " 10!4.

Charbonneau+ (2002)
惑星大気によるNaの吸収 & 雲 ?

Prediction by Seager & Sasselov (2000)

conditions.

hydrogenatomsatoms
are sensitive
to the fromstellar theradiationhydrogenshould escapeplanet. pressure
The spectral
absorption width shows independently that the atoms have large
velocities relative to the planet. Thus hydrogen atoms must be
escaping the planetary atmosphere.
We have built a particle simulation in which we assumed thathydrogen atoms are sensitive to the stellar radiation pressure

Vidal-Madjar+ (2003) HST/STIS, Lyman-α

The Astrophysical Journal, 774:95 (17pp), 2013 September 10

透過光分光 (1)Deming et al.

1200 K isothermal temperature structure, as used above. The
grid utilizes the methodologies described by Burrows et al.
Knutson+ (2007) (紫外-近赤外, HST/STIS)(2001) and Burrows et al. (2010) and Howe & Burrows (2012),
but it incorporates
different
amountsof extra
The AstrophysicalJournal,
774:95(17pp),
2013graySeptember
10
Deming+(2013)(近赤外,
H2opacity.
O, HST/WFC3)
We interpolate in this grid to find that an extra opacity of
0.012 cm2 g−1 matches the 1.4 µm water absorption at the
bandhead, and provides suitably low absorption at 1.15 µm.
The lowering of the 1.15 µm absorption occurs because that
H2O
intrinsically weaker band requires a longer path length toproduce significant absorption, and long path lengths are masked
Figure 2 The HD209458 Lyman a profile observed with the G140M grating. The
by the extra opacity. A Burrows isothermal model having no
geocoronal emission has been subtracted; the propagated errors are consequently larger
Figure 3 Relative flux of Lyman a as a function of the HD209458’s system phase.
extra opacity (not illustrated) shows a much more prominent
in the central part of the profile, particularly in the Geo domain (see text). Dl represents The averaged ratio of the flux is measured in the In (1,215.15–1,215.50 Å and 1,215.80–
peak at 1.15 µm.
the spectral resolution. a, The thin line shows the average of the three observations
1,216.10 Å) and the Out (1,214.40–1,215.15 Å and 1,216.10–1,216.80 Å) domains in
performed before the transits (exposures A1, B1 and C1); the thick line shows the average individual exposures of the three observed transits of HD209458b. The central time of
The profile of the 1.4 µm band is not matched optimallyof the three observations recorded entirely within the transits (exposures A2, B3 and C3). each exposure is plotted relative to the transit time. The vertical dashed lines indicate theby the isothermal models, not as well as on Figure 10 for
Variations
seenHD209458
in the In domain
over theblueof thegrating.
line andThe first and the second contact at the beginning and the end of the transit; the exposures A1,
Figureare
2 The
Lymanasa absorption
profile observed
withthepeak
G140M
example. The real absorption line of sight passes through21
21
partially
over theemission
red peakhas(between
2130 km thes propagatedB1 andFigure
C1 wereperformed
the transits,
and the exposures
A2, B3 andC3 wereand 100 kmerrorss ). are
Quotedvelocities larger
geocoronalbeen subtracted;
consequently
3 Relative
fluxbeforeof Lymana as a functionof the HD209458’s
systemphase.
different temperatures on day and night hemispheres of the21
entirely
withinthetransits.
The
ratio
is
normalized
totheaverage
value
of
the
three
in
the
heliocentricreferenceare in the central
stellar reference
frame,
centredon213
kmspart of the profile, particularly in the Geo domain (see text). Dl represents The averaged ratio of the flux is measured in the In (1,215.15–1,215.50 Å and 1,215.80–
planet. Figure 10 accounts (crudely) for different temperatures
observations
completed
beforebeginning of the transits.
The ^1j error bars Å)
aredomains inframe.
^1j error
bars. c,a,The
two spectrain the ofIn thedomain,spectrathe b,
spectralresolution.
Theratiothin oflinetheshows
the average
threetheobservations
1,216.10
Å) and the
Out the(1,214.40–1,215.15
Å and 1,216.10–1,216.80
along the line of sight, not included in the isothermal model
statistical;
they arecomputed(see text). TheThe
In/Outratio time of
being
normalized
suchthethat
the ratiois 1 in theratio line
is plotted
performed
beforetransits(exposures
A1,OutB1 domain.
and C1);Thisthe thickshowsasthea average
individual
exposuresof throughthe threeboot-strap
observedestimationstransits of HD209458b.
central! (triangles), and as a function of l2 using
for Figure 14. Including that line-of-sight temperature variationsmoothlydecreasesby
15% during
thetransit
transit.
TheThethickline represents
theindicate thefunctionl1 using
l2 ¼ 1; 216:10
of theofthreeobservations
recordedA entirely
within the transits (exposures A2, B3 and C3).
each
exposure
is aroundplottedspectrum
relativeto thetime.
verticaldashedlines
Figure14.
TransmissionofHD
209458b
derived
from
HST
spec!
absorption
ratiothemodelled
throughaare
particlesimulation
which
includeshydrogenl1 ¼
1; 215:15
TheInratio
is always
significantly
1, witha minimum
may be essential to matching the band profile.
Variations
areA (circles).
seen in thedomainas absorption
overbelow
the bluepeakof the lineatandtroscopy.
first Ourandsecondresults
contact at
the thebeginningandthe endTheof theopentransit;the atoms
exposures
A1,
WFC3
solid
points.
squaresareour! (2130 km s21) and l2 ¼ 1; 216:10! (100 km s21).21In the domain
escapingtheplanet.
In bands
this simulation,hydrogenatoms
aresensitive
to A2,
the radiation
l1 ¼
1; 215:15
The Burrows model on Figure 14 is sufficiently highpartially
overAthered peak (between 2130 km s21AandB1 from
andofC1
were
performed
beforethetransits,
and theB3
andtheC3 diawere
100 km s ). Quoted velocities
re-analysis
theSTIS
definedbyKnutsonetexposuresal. (2007b),
andpressure
abovenarrowan altitudeof 0.5 times
Rocheradius,fromwhere
theaveragedensity value
is assumed
todefinedbythe
thesevalues,the Lymana intensity
decreasesby reference
mond
is
the
band
absorption
Charbonneau
etof al.
(2002).
entirely
withinthesodiumtransits.
The theratio
is normalized
to thethethreethetransits
heliocentric
are instellarreferenceframe,
centred on
213 kmduring
s21 inthe
resolution in wavelength to permit meaningful comparison with
5
23
10
21
be
2
£
10
15 ^
4%.
The
detectiondoes
notstrongly
dependonaparticular
selection
of
the
cm
;
these
two
parameters
correspondtoanescape
flux
of
,10
g
s
.
The redobservations
line is thecompleted
transmittancespectrumfromisothermal
Burrows
model,
before thebeginningof theantransits.
The ^1jerror bars
areframe. b, ^1j error bars. c, The ratio of the two spectra in the In domain, the spectrascale
the sodium absorption
measured by Charbonneau et al. (2002)
2−1
domain.
Whilethedecrease
of
the
Lyman
The
stellarradiationpressure
is
taken
to
be
0.7
times
the
stellargravitation.
The
mean
a
intensityisnotsensitive
tothepositionofl,havingan
extra
opacityofgraycharacterandmagnitude
0.012
cm
g
.
The
2
statistical; they are computed through boot-strap estimations (see text). The In/Out ratio
being normalized such that the ratio is 1 in the Out domain. This ratio is plotted as aand Snellen et al.height(2009). We plot the “narrow” band absorption
lifetimeofescapinghydrogen
atoms
is
taken
to
be
4
h.
The
model
yieldsanatoml,showing
that
most
of
the
absorption
occursinthe
it is function
more sensitivetothepositionofreddiamondintegrates
the
red
model
over
the
sodiumbandpass.
The
blue
line
1
!
smoothly decreases by around 15% during the transit. The thick line represents the−4of l1 using l2 ¼ 1; 216:10 A (triangles), and as a function of l2 usingfrom Charbonneau et al. (2002) on Figure 14 (triangle point
ispopulationa Dobbs-Dixon
for tail,
HDexplaining
209458b,
no grayabsorption
opacity, lasts
but wellwith λin a curvedmodel
comet-likewhywiththe computed
blue part of the line. !Using the whole domain where the absorption is detected, theincludeshydrogenatoms
l1 ¼ 1; 215:15 A (circles). The ratio is always significantly
below 1, with a minimum at absorption ratio modelled through a particle simulation which
2 than!3j.
(Rayleigh)normalized to magnitude 0.001 cm2 g−1 at 0.8 µm. Bewith error bar). The Snellen et al. (2009) results (not plotted)
after the endopacity,of the transit.
exoplanetary atmospherichydrogen21
is detected at more
21!
escaping from the planet. In this simulation, hydrogen atoms are sensitive to the radiationl1 ¼ 1; 215:15 A (2130 km s ) and l2 ¼ 1; 216:10 A (100 km s ). In the domain
cause the blue model has no gray opacity, we scale-down the modulation in thisare consistent with Charbonneau et al. (2002), considering thepressureabove
analtitude
of
0.5
times
the
Roche
radius,wherethedensityisassumedtodefinedbythesevalues,theLymanaintensity
decreases
duringthetransitsby© 2003 Nature Publishing
spectrumGroupby a factor of three for this comparison (see text).
NATURE | VOL 422 | 13 MARCH 2003 | www.nature.com/nature
145
different bandpasses. Integrating the Burrows model having
be 2 £ 105 cm23; these two parameters correspond to an escape flux of ,1010 g s21.
15 ^ 4%. The detection does not strongly depend on a particular selection of the(A color version of this figure is available in the online journal.)
0.012 cm2 g−1 extra opacity over the band used by Charbonneau
domain. While the decrease of the Lyman a intensity is not sensitive to the position of l2,
The stellar radiation pressure is taken to be 0.7 times the stellar gravitation. The mean
et al. (2002; red diamond point) produces agreement within theit is more sensitive to the position of l1, showing that most of the absorption occurs in the lifetime of escaping hydrogen atoms is taken to be 4 h. The model yields an atom
population in a curved comet-like tail, explaining why the computed absorption lasts well Figure
blue part of the line. Using the whole domain where the absorption is detected, theerror
14.bar.
Transmission spectrum of HD 209458b derived from HST spec-

Na, H, CO, H O,
Fe, He…

12
gr
(2
bu
W
0.
ba
Th
in
pr
by
ex
pe

by
ex
di
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alo
fo

2 –4

Correlation c

Correlation coe

2–2

NATURE | Vol 465

透過光分光 (2): 分子吸収線の高分散分光
–2σ

0
–100
0 Geocentric–100
100radial velocity
200 (km s–1)
–1)
Geocentric
radialvelocity(kmsb Snellen+ (2010) VLT/CRIRES
c
0.03 Observed

c

+3σ
+3σ

Artificial signal
3× nominal

Artificial signal
+2σ
3×+2σ
nominal200

Correlation coefficient × 106

–3σ
100

–200

Observed–3σ–2σ

Correlation coefficient × 106

Or
–200

–1σ

–0.01
–0.02

–1σ

–2

NATURE | Vol 465 | 24 June 201–4–150 –100 –50
050
100 150–150 –100 –50
0veloci
Barycentric radial–1
Barycentric radial velocity + 14.8 (km s )Orbital phase

Figure3|
Thecarbonmonoxidesignalinte
Figure3|
Thecarbonmonoxidesignalintegrated
over
the
tran
0
+1σ
cross-correlation
signalfromallspectratakcross-correlation
signalfromallspectrataken
duringthetransit
0
0.01
21
combined
assuming
a planetary
orbitalvelo, incombined
assumingaplanetary
orbitalvelocity
of
140
km
s
–1σ
weightedby the signal
depth atofthe
theobserved
transit signal
0.00
weighted by the depthof the transit
epoch
–2σ
–1σ
integrated signal
is statisticallyintegrated –2
signal is statisticallysignificant
at the 5.6ssignificant
confidencea–0.01
23ratio of (1–3) 3
deriveaCO
volumemixing–3σderive a CO volume mixing ratio of (1–3) 3 10 for the upper a–2σ–2 of HD 209458b, with the precision governe
0
0
–100
100
200
of HD 209458b, with the precision governed by the uncertainty–0.02–1
Geocentric radial velocity (km s )–3σpressure–temperature
profileand inofthe
levepressure–temperature
profile
and
in
the
level
of
masking
the
C
–4
c
–60
–40
–20
0
20
40
60
–60
–40
–20
0
20
40
60
21
–60
–40
–20
0
20
40
60
–60
–40
–20
0
20
40
60
0
00
100
200
CH
.
The
CO
signalisblueshiftedby ,2tokm
50 km100
150 respect
–150
–100
0 by ,2
CH4. The COsignalis blueshifted
s withth
4–50–1)
Barycentric+ 14.8 (km s–1)
Barycentric
radial14.8radial(km svelocity
servedArtificial
ocentric
radialvelocity(kmvelocitys–1) +signal
ofvelocity
the
host
which
suggestsav+ 14.8(kms–1) flow
velocity of theBarycentric
host velocity
star,radialwhichsuggests
astar,
velocity
from the
3× nominal–4 thenightside
driven by
theflux
large
incidenththenightside
drivenbythelarge
incident
heat
on
the
daysidc signal in spectrum
Figure1 | transmissionCOthe transmission
spectrum of exoplanet
CO signalin theof exoplanet
Figure 3 | The carbon –150
monoxidesignal–50
integrated
the transit.
100 The150–100
0 over50
The iscross-correlation
shown between
a template spectrum
HD209458b.Artificial
cross-correlation
shown
spectrum58b. The
cross-correlationfromall spectraradial
takenvelocity
during thetransit(kmwere–1)
signalbetween aistemplate-signal
高分散分光による軌道速度決定+ 14.8sBarycentric
21
lines and
Large
spectraofHD
209458takenof Very
56 COlinesTelescope
and
Very
Large
Telescope
spectraofHD
209458taken, individual
combineda planetary
orbitalvelocityof average
140thekm sCO
Thisassumingmeans thatonThisaverageperspectrumsignalas
3× nominal
means
that
on
per
spectru
a planet between
orbital phaseof 20.025
h , 0.035.
The beginning
andweighted
a planet
orbital,phaseof 20.025, h , 0.035.
The beginning
and
by
the
depth
ofisthetransitsignalatlevel.
thepresent
observed
epoch.
The-and
2 km/sのblueshift(~
0.01–0.1
mbar)
in
Fig.
1
2
present
at
a
1s
in
Fig.
1
and
2
is
at
a
1s
level.
Figure3|
Thecarbonmonoxidesignalintegrated
over
the
21
21
e transitend
are ofat hthe
6 0.018.
velocityof the hoststarintegrated
signalisstatisticallysignificant
at
the
5.6s
confidence
level.
W
transitTheare systemic
at h 6 0.018.
The systemic
velocity
of
the
host
star
Our
directdetectionofabsorption
lines
leads
to
an
una
Our direct
detection
ofduring
absorptionli→ 昼面から夜面への風？
cross-correlation
signalfromallspectrathetra21
23 taken21
derive aof
COvolumemixing ratio of (1–3) 3 10 for the upper atmospher
the velocity of and
the Paranal
58a is 214.77
km s (blueshifted),
the velocity
theParanal
HD 209458a
is 214.77 km ands (blueshifted),2
identificationofcarbonmonoxideintheatmosphereofHD
identificationofcarbonmonoxideinth21
combined
assumingaplanetary
orbitalvelocity
of
140
km
s
21
of
HD
209458b,withtheprecision
governedbytheuncertainty
intheat theof s at the time ofory in theobservatory
direction ofinthe
is 11.0 km
s starthestardirectionof theis time11.0 kmand allows us to determine the CO abundance. Although tran
0.02

+1σ

透過光分光 (2): 分子吸収線の高分散分光

NATURE | Vol 465 | 24 June 201

HOWMAN ET Showman+
AL.
Vol. 682(2008) with Newtonian cooling+3σ

served

+1σ

–1σ
–2σ
–3σ
0
–100
100
Geocentric radial velocity (km s–1)c
Artificial signal
3× nominal

200

Correlation coefficient × 106

0

+2σ

2

0

–2

–4
50
100 150–150 –100 –50
0Barycentric radial velocity + 14.8 (km s–1)

Figure 3 | The carbon monoxide signal integrated over the transit. The
cross-correlation-signal
from all spectra taken during the transit were
高分散分光による軌道速度決定
21
combined assuming a planetary orbital velocity of 140 km s , individual
weighted by the depthof theのblueshifttransit signal(~
at 0.01–0.1
the observed
epoch. The- 2 km/s
mbar)
integrated signal →is statisticallysignificant
at
the
5.6s
confidence
level.
W
昼面から夜面への風？23
derive a CO volume mixing ratio of (1–3) 3 10 for the upper atmospher
of HD 209458b, with the precision governed by the uncertainty in thecf) spatially-resolved transmission spectroscopy of HD 189733b
Louden, Wheatley (2015) HARPS, Na lines
The Astrophysical Journal Letters, 814:L24 (5pp), 2015 December 1
See also Brogi et al. (2016)

Louden & Wh

signal-to-noise data it is possible to model variations
planetary absorption line proﬁle to spatially resolve velin the planetary atmosphere.
The star–planet system was simulated on a 2D pixe
which implicitly accounts for both the atmospheric abso
and the Rossiter–McLaughlin effect. A quadratic limb d
ing proﬁle was used to give each pixel on the star a rintensity (Claret & Bloeman 2011). We adopted arotational velocity of 3.10±0.03 km s−1, which was me
The Astrophysical Journal Letters, 814:L24 (5pp), 2015 December 1 Average
Louden& Wheatleyusing
Doppler
tomography
the sameset of H
Trailing
Leading
Model of from
Flowers+(2019)
observations (Collier Cameron et al. 2010). Solid body roof the star was assumed in order to derive the rotation
but if this assumption is not correct the rotation velocity
transit chord will not be affected.
The absorption by the planet is modeled as an oFigure 1. Top panel: a scale diagram of the star–planet geometry, with theocculting disk with the white light radius ratio of HD 189
thickness of the planetary atmosphere doubled for clarity. The planet moves
The atmospheric absorption is represented as an add
from left to right as it transits the star, with the leading limb colored red and thelayer with absorption of the sodium doublet modeled bytrailing limb colored blue. Bottom panel: our model illustrates how thebackground illumination of the planetary limb changes during transit, allowing
proﬁles. The atmospheric dynamics are parameterized
absorption by different parts of the planetary atmosphere to be separated. The
equatorial
velocitiesontheleading
andtrailing
limbs
red/blue lines represent the stellar ﬂux available for sodium absorption at theplanet together with the assumption of constant aplanetary equator on the leading/trailing limbs (the model accounts for stellar
1.3 (rigid
velocity
rotation velocities
and limbfrom
darkening,
andanalysis.
the planetary
orbit).limb
Theasymmetry
Figure 3. Posterior distributions of atmospheric
our bootstrap
On the leadingof theplanet (left)ina red shiftof 2.31.5
km s−1 is bodyfound. rotation). This allows continuous

Note: 見ている高度
atomic/ion lines

透過光
molecular lines

観測者
惑星光~10-2–1 bar

惑星

Transit, Phase Curves, Eclipse

Thermal phase curves
The Astrophysical Journal, 790:53 (9pp), 2014 July 20
Zellem+ (2014) Spitzer, 4.5µm
1.00251.0015
Relative Flux

models. Overall, our observed dayside emission (phto 0.75 in Figure 6) matches the predictions of the
et al. (2009) models well, especially for models includ
absorbers and a dayside inversion.
There exist a variety of factors that can influenceof the hot spot for a given hot Jupiter. Generally slarger atmospheric opacity would move the photohigher altitude (lower pressure), in which the radi
constant is shorter, leading to a smaller hot-spot oconversely for smaller atmospheric opacity (Dobbs
Lin 2008; Showman et al. 2009; Lewis et al. 20
et al. 2011a). Such opacity variations could result
食の前でピーク= 温度が最大となる地点
gas
metallicity (higher
metallicity generally impli
opacityandviceversa),orpotentiallyfromhazes,
thが主星直下点の東側にずれている
is currently no strong evidence that hazes significa→emission
super-rotationの予言と一致
theon HD 209458b (Deming et al. 2013).
for specified opacities, the existence of atmosphericdrag could lead to slower wind speeds, decreasing thTemperature [K]1.00201.00101.00051.0000Best-fit Model
1x Solar no Inversion
1x Solar w/ Inversion
3x Solar w/ Inversion

0.9995-0.6

ZeShowman+ (2009)

-0.4

-0.2

0.0
Orbital Phase

0.2

0.4

0.6

Figure 6. Comparison of our reduced data (black points) and its best-fit phasecurve model (red line) to the Showman et al. (2009) GCMs with 1× solar
abundance and no thermal inversion (green asterisks), an inversion with 1×
solar abundance (blue diamonds), and an inversion with 3× solar abundance高分散分光による惑星光の検出
観測
HCN
2 O + CO90 H
60

radiancetaken at different nodding positions may be combined to accurately remove the background4 and improve the signal. The
3
observations target bright, daysideemission from the com2
1
panion planet HD 209458b, nearsecondary eclipse. Each set
0
2.29
2.3
2.31
2.32
2.33
2.34
2.35
of spectra are phase resolved in the rangeswavelength⇠ 0.5- 0.6 and[um]
⇠ 0.4 - 0.5 for the 2.3µm and 3.2µm bands, respectively.
The planet has an orbital period T = 3.52 days and reference
crossing time t =0 = 52854.325456 JD (Butler et al. 2006).

30
0

Frame Number

2.2. Initial Reduction

3
モデル (x2)

90 H2 O
60
30
0
90 CO
60
30

0
Data from each detector on each observation night are
90 HCNH2O + COtreated separately throughout the following procedures.
60
We use Esorex from the ESO CRIRES reduction toolkit
30(v2.3.4) to reduce the two-dimensional spectral images, per0
20 0
20
40
80
60
40
20 0
20
40
60
80 40
form flat-fielding and bad pixel corrections, and optimally
11 )
1
V
(km
s
V
(km
s
)
V
(km
s
)
rest
rest
rest
extract (Horne 1986) one-dimensional spectra. We perform
主星と惑星の光を同時に観測し、Hawker+ (2018) VLT/CRIRES 2.3µm, 3.2µm
additional calibrations with our custom pipeline, X-COR,
その中の惑星光由来の分光特性を
Figure 2. Frame-by-frame cross correlation values as a function of
written in Python 2.7. Remaining bad pixels and regions
velocity. The CCFs are shifted into the planetary rest frame by the
高分散分光で探す
are replaced with the linear interpolation of their neighbors.
→ 3Dモデルと合わせるとsignificance(e.g.,
2020)
peak significance Kp , for each model 向上
template(HCN,Beltz+
H2 O, CO,and
Next, all of spectra are aligned to the highest signal-to-noise
their combination). The right panel shows the CCF with a model

3. High-resolution model spectra of HD 209458b showing line features of HCN, H2 O, and CO in the observed bands at 2.28-2.35µm
8-3.27µm.

e strongest absorption lines in the atmospheric model,
sed in Section 3.
Doppler shift the template by velocities ranging fromm s-1 to 250 km s-1 in steps of 1.0 km s-1 . At eachy, we linearly interpolate the template onto the dewavelength grid, and take the dot product with eachded spectrum. We then combine the cross-correlation
ns (CCFs) from all four detectors and both nights. The
nt matrices are shown in Figure 2. The dark trail at theary Vsys represents cross correlation from the alignment
el and intrinsic features.
sample Kp , the semi-amplitude of the orbital velocer a grid ranging from 30 km s-1 to 250 km s-1 in
f 1.0 km s-1 , and shift each row of the CCF matrix
orbital velocity Kp sin(2⇡ (t)). For the correct Kp thisrmation shifts the CCF matrix into the rest frame ofr–planet system, which itself has the systemic veloc. We then sum the pixels within an N ⇥ 5 sliding
w, centered on velocities ranging from -80 km s-1 tos-1 ; this step samples potential values for Vsys . If thete absorption lines are present in the data, then we exe sum to be maximized at the planetary Vsys . Finallyfor the data (left) and with a model injection at two times
nominal strength (right).

3. MODELING

In order to cross correlate with the data we generate highresolution model atmospheric spectra for HD 209458b using
the GENESIS code (Gandhi & Madhusudhan 2017). The
models involve line-by-line radiative transfer computed via
the Feautrier method to obtain the emergent spectra in the observed bands at a spectral resolution of R &300,000 (Figure3). The models assume a plane-parallel atmosphere in hydrostatic equilibrium and local thermodynamic equilibrium.
The atmosphere comprises 150 model layers evenly spaced
in log(P) for pressures from 10-8 -100 bar. We adopt the P–
T profile derived from recent retrievals of dayside emission
spectra both at low resolution (Line et al. 2016) as well as
at high resolution (Brogi et al. 2017). We also test a number of P–T profiles to confirm the robustness of our results,
including a radiative-convective equilibrium profile (Gandhi& Madhusudhan 2017). We consider molecular opacity
due to the prominent sources of absorption in the observedホットジュピター大気データの蓄積と見え始めた統計的な性質
From individual characterization
to statistical samples

The Astrophysical Journal, 828:22 (20pp), 2016 September 1

Parmentier et al.

Clouds ?

( Phase curve@Keplerからの示唆 )Figure 1. Pressure–temperature proﬁles from our ﬁducial model for a planetwith Teq=1900 K (gray: proﬁles at all latitudes and longitudes; light green:
proﬁles at the limb west of the substellar point; dark green: proﬁles at the limb
east of the substellar point; blue: proﬁles at less than 20◦ of the antistellar point;
red: proﬁles at less than 20◦ of the substellar point). The condensation curves ofseveral important species are plotted as dashed lines. At a given pressure, if thetemperature is cooler than the condensation temperature, we consider theatmosphere cloudy. For pressures lower than Ptop, we assume that clouds are
not present, with this cloud top pressure used as a free parameter.

Parmentier+ (2016)

Figure 1. Pressure–temperature proﬁles from our ﬁducial model for a planetwith Teq=1900 K (gray: proﬁles at all latitudes and longitudes; light green:
proﬁles at the limb west of the substellar point; dark green: proﬁles at the limb
east of the substellar point; blue: proﬁles at less than 20◦ of the antistellar point;
red: proﬁles at less than 20◦ of the substellar point). The condensation curves ofseveral important species are plotted as dashed lines. At a given pressure, if the
Figuretemperature is cooler than the condensation temperature, we consider
the 3. Phase shift of the maximum of the Kepler lightcurve relative to thesecondaryeclipse (top) and apparent geometric albedo of the planet in theareatmosphere cloudy. For pressures lower than Ptop, we assume that clouds
Kepler bandpass (bottom) as a function of the equilibrium temperature of thenot present, with this cloud top pressure used as a free parameter.
planet. The apparent geometric albedo includes the contribution from both thethermal and the reﬂected light. Different colors are analyses from different
authors. Purple circles are the data from Esteves et al. (2015)—used in thisstudy—green squares are from Angerhausen et al. (2015), blue triangles arefrom Shporer & Hu (2015),and orange diamonds are from Heng & Demory(2013). Kepler-43b was removed from the analysis as it is doubtful whether thesignal is of planetary origin or due to stellar activity(Esteves et al. 2015).

Figure 2. Phase curves from our ﬁducial model for a planet with Teq=1900 K

not form, and a particle size, both considered as freeFigure 3. Phase shift of the maximum of the Kepler lightcurve relative to theThe Astrophysical Journal, 828:22 (20pp), 2016 September 1

Clouds ?

( Phase curve@Keplerからの示唆 )Figure 1. Pressure–temperature proﬁles from our ﬁducial model for a planetwith Teq=1900 K (gray: proﬁles at all latitudes and longitudes; light green:
proﬁles at the limb west of the substellar point; dark green: proﬁles at the limb
east of the substellar point; blue: proﬁles at less than 20◦ of the antistellar point;
red: proﬁles at less than 20◦ of the substellar point). The condensation curves ofseveral important species are plotted as dashed lines. At a given pressure, if thetemperature is cooler than the condensation temperature, we consider theatmosphere cloudy. For pressures lower than Ptop, we assume that clouds are
not present, with this cloud top pressure used as a free parameter.

Parmentier+ (2016)

Figure 3. Phase shift of the maximum of the Kepler lightcurve relative to thesecondary eclipse (top) and apparent geometric albedo of the planet in the
Kepler bandpass (bottom) as a function of the equilibrium temperature of theplanet. The apparent geometric albedo includes the contribution from both thethermal and the reﬂected light. Different colors are analyses from different
authors. Purple circles are the data from Esteves et al. (2015)—used in thisstudy—green squares are from Angerhausen et al. (2015), blue triangles arefrom Shporer & Hu (2015),and orange diamonds are from Heng & Demory(2013). Kepler-43b was removed from the analysis as it is doubtful whether thesignal is of planetary origin or due to stellar activity(Esteves et al. 2015).

Figure 2. Phase curves from our ﬁducial model for a planet with Teq=1900 K

not form, and a particle size, both considered as freeFigure 6. Offset of the maximum of the Kepler light curve compared to the

parti
there
Th
albed
“app
seco
the p
plan
by t
emis
plan
brium
Kepl
plan
plan
days
temp
in th
hom
albed
deter
the c
B
Figu
curv
narro
of M
their
with
mod

NATURE
Keating+ ASTRONOMY(2019) but see also Beatty+ (2019)
dayside temperature

2,500

Teq

WASP-33b

WASP-18b

3,000

Tday (K)KELT-1b
HAT-P-7b

WASP-12b

WASP-14b
WASP-19b

2,000

Analytic model
Analytic model, including hydrogen dissociation( Phase curve@Spitzerからの示唆 )2,000

WASP-103b

2,500

Clouds ?

HD 149026b
The Astronomical
Journal, 158:166 (18pp), 2019 October

1,500
HD 149026b
HD 189733b

1,000

WASP-43b

However, it seems more likely that KELT-1b’s observed
HD 189733b
thermal HD
emissionis being heavily modulated by clouds. 500
209458b
Furthermore, when we consider the ensemble phase curve
properties of KELT-1b and the hot Jupiters, it appears that 0
essentially temperature
all of the objects for which we have 3.6 andnightside
4.5 μmSpitzer phase curves also show evidence for clouds2,500
affecting their thermal emission properties.
We see this in four different ways:
2,000
CoRoT-2b

1,500
1,000
2,500

flat

WASP-33b

WASP-43b

HAT-P-7b

WASP-12b

WASP-19b

Beatty et al.

Beatty+ (2019) but see also Zhang+ (2018)

HD 209458b

WASP-14b

CoRoT-2b

WASP-18b

Energy balance model
Energy balance model, including hydrogen dissociation
WASP-33b

1. The most recent set of publishedWASP-103b
phase offset measureWASP-103b
1,500
ments—including KELT-1b—are
consistent with the1,500
WASP-12b
WASP-12b
HAT-P-7b
HAT-P-7b
HD 149026b
HD 149026b
observed planets WASP-19b
having aKELT-1b
constant phase offset of 14°
WASP-19b
WASP-43b
HD 189733b
WASP-14b
WASP-43b
forallplanetaryequilibrium temperatures, though with a
HD 189733b
1,000
WASP-14b
1,000
HD 209458b
HD 209458b
highscatter. This conﬂicts withthermal-only GCM
WASP-18b
CoRoT-2b
WASP-18b
CoRoT-2b
predictions
that cooler planets should show large
500
500(∼70°) offsets (Zhang et al. 2018).
2. The relatively low nightside ﬂux that we measure for
KELT-1b requires that the underlying atmospheric 0 Figure 6. 3.6 and 4.5 μmphase curve offsets as a function of equilibrium
0
1,500
2,000
3,000
4,000
intensity2,500
map (as opposed
to 3,500
the disk-integrated
ﬂux we 1,500temperature.
The light 2,500
blue rectangle3,000
shows our measured
offsetsfor KELT-1b.
2,000
3,500
4,000
For our analysis we used the results of Mendonça et al. (2018) for WASP-43b,(K)actually observe)T0display
a sharp transition between the
T0 (K)
rather
than
those
of
Stevenson
etal. (2017), but as a comparison we also showday- and nightsides, and be at a roughly constant intensity
Tnight (K)Tnight (K)

2,000

WASP-33b

WASP-103b

Tnight (K)3,500

LETTERSthe points from Stevenson et al. (2017) as the two grayed-out points offset just

Clouds ? ( 透過光スペクトルからの示唆 )A continuum from clear to cloudy hot-Jupiter exoplanets without primordial water depletion 4

Sing+ (2016)

LETTERS

Gao+ (2020)

80

Na

75

K
H2O

WASP-17b

WASP-39b

60
HD209458b

55

WASP-19b

50
45

HAT-P-1b

40
35

WASP-31b

30
WASP-12b

25
20

HAT-P-12b

15

HD189733b

WASP-6b5.3.4.5 .6 .7 .8 .9 1
1.5
2
Wavelength (µm)

2.5 3 3.5 4

5

Figure 1 | HST/Spitzer transmission spectral sequence of hot-Jupiter survey targets. Solidcoloured lines show fitted atmospheric models with prominent spectral features indicated. The spectra
have been offset, ordered by values of ΔZUB-LM (the altitude difference between the blue-optical andmid-infrared, Table 1). Horizontal and vertical error bars indicate the wavelength spectral bin and 1σmeasurement uncertainties, respectively. Planets with predominantly clear atmospheres (top) showprominent alkali and H2O absorption, with infrared radii values commensurate or higher than theMuted H2O absorption features insome transmission spectra

4 m s–2
10 m s–2–2
25 m sClear
Cloudy

5

3.3
3.2

4

3.1

3

3.0

2

2.9

1

2.8

0

2.7
500

10

0

6

900

1,300

1,700

2,100

Equilibrium temperature (K)

2,500

a

log[g (cm s–2)]

Relative altitude in scale heights z(λ)/Heq

65

1.4 µm water band amplitude (H)

70

3.4

2,900

Fig. 1 | Exoplanet cloudiness as a function of equilibrium temperature,
Hydrocarbons ?
Silicate clouds ?
gravity and atmospheric metallicity. Comparison between the observed15(points) and model (shaded regions) difference in planetary radii between

b

0
2
0
2

Discovery
Mass andRadius0302Spectroscopy

0
5
0
2

Detection

Insights
System architecture
Occurrence rate
Correlation with stellar properties

Details0402太陽型星周りの温暖な岩石惑星

Observation Type

低質量星周りの温暖な岩石惑星

ガス惑星・高温の惑星

0
1
0
2

Bulk Composition(gas/water/mantle/core fraction)
Atmospheric Profiles
Characterization(Surface Composition)of golden targets

Fig.1 Development of Exoplanet Exploration

Photometric/
Spectroscopic
Survey0202Discovery
Mass andRadius0302Spectroscopy

0
5
0
2

Detection

Insights
System architecture
Occurrence rate
Correlation with stellar properties

Details0402太陽型星周りの温暖な岩石惑星

Observation Type

低質量星周りの温暖な岩石惑星

ガス惑星・高温の惑星

0
1
0
2

Bulk Composition(gas/water/mantle/core fraction)
Atmospheric Profiles
Characterization(Surface Composition)of golden targets

Fig.1 Development of Exoplanet Exploration

Photometric/
Spectroscopic
Surveyより小さな惑星、よりcoolな惑星への展開

主星光と惑星光のコントラスト
透過光の
シグナル

惑星光の
シグナル

観測機器の系統誤差の影響
透過光の
シグナル

惑星光の
シグナル

HZ惑星（mini-Neptune? super-Earth?）の透過光スペクトル8

Benneke et al.

K2-18b (2.6 Re, 8.63 Me)

Benneke+ (2019); see also Tsiaras+ (2019)

Figure 3. Transmission spectrum of K2-18b computed from our global spectroscopic and broadband transit light-curve analysis

WFC3’s transit depth estimates with those of the Spitzer Space
elescope’s Infrared Array Camera at 4.5 μm (ref. 1) implies a lackf significant absorption features between the two different specal ranges covered24, thereby further indicating the absence of clear
ydrogen-dominated atmospheres.
Hydrogen is a powerful greenhouse gas, and its presence inubstantial amounts in an atmosphere therefore affects a planet’s

high surface temperatures and pressures for the inner planets thatare incompatible with liquid water. To be habitable, the inner planets must therefore have lost most of their atmospheric hydrogen, or
never accreted or outgassed important amounts of hydrogen in thefirst place6.
Given the irradiation levels experienced by the Earth-sizedplanets in TRAPPIST-1’s habitable zone, theory suggests that the

TRAPPIST-1 planets の透過光観測(RP/R*)2 – (RP/R*)2WFC3 (ppm)

H2 rich (H2O) (8σ)
H2 rich (CH4) (13σ)
H2O richN2 rich
CO2 rich

TRAPPIST-1 d

1,000
500
0
–500
–1,0001.1bc(RP/R*)2 – (RP/R*)2WFC3 (ppm)

1,500

1,500

1.2

1.3

1.4
1.5
Wavelength (μm)

1.6

1.7

d

H2 rich (H2O) (6σ)
H2 rich (CH4) (10σ)
H2O richN2 rich
CO2 rich

TRAPPIST-1 e

1,000
500
0
–500
–1,000
1.1

1.2

1.3

1.4
1.5
Wavelength (μm)

1.6

1.7

1,500

H2 rich (H2O) (4σ)
H2 rich (CH4) (5σ)
H2O richN2 rich
CO2 rich

TRAPPIST-1 f

1,000
500

雲のないH2-richな大気は
棄却

0
–500
–1,0001.1(RP/R*)2 – (RP/R*)2WFC3 (ppm)(RP/R*)2 – (RP/R*)2WFC3 (ppm)ade Wit+ 2018 (see also e.g., Zhang+ 2018)

1,500

1.2

1.3

1.4
1.5
Wavelength (μm)

1.6

1.7

ている。2次大気は保持できるかも。

H2 rich (H2O)
H2 rich (CH4) (3σ)
H2O richN2 rich
CO2 rich

TRAPPIST-1 g

1,000
500

(Stökl+ 2015,
Unterborn+ 2018, Hori &Ogihara 2019, Dorn+
2018, Airapetian+ 2017)

0–500–1,000
1.1

1.2

1.3

1.4
1.5
Wavelength (μm)

1.6

理論的にも、H2を保持していないだろうと言われ

1.7

g. 2 | Transmission spectra of TRAPPIST-1 d, e, f and g compared with synthetic atmospheres dominated by hydrogen (H2), water (H2O), carbon

creasebufferSwain
but inlong r(e.g.,
4
Stevenson et al.
WFC3
Hawaii detector, we can anticipate some of the instrument Howesystematics that we are likely to encounter. For example, slit lo
at fluences > 40, 000 e- per pixel, WFC3 experiences an in- or drif
crease in intensity (similar to a ramp or hook) after each MIR
buffer dump (e.g., Berta et al. 2012; Deming et al. 2013; IRAC
Swain et al. 2013; Wilkins et al. 2014). For similar reasons, exhib
but independent of flux, WFC3 also exhibits an HST orbit- Harrin
long ramp that is readily apparent in most spatial scan data Steven(e.g., Kreidberg et al. 2014b; Wakeford et al. 2016). Because was si
Stevenson+
2016
WFC3
is
a
slitless
spectrograph,
it
is
not
subjecttoslitlosses. exper
F IG . 2.— Recommended spectroscopic
observingmodes
fortransitingex!!100 ppm
15 ppm <=
=>
30-50
ppm
=>
However,
NIRSpec’s
1.6
aperture
maywillbe besubject
to are lik
oplanets. For objects brighter than J ∼square
11, at least
three visits
required slit
to achieve
coveragefrom 0.6 to 5.0 µm.telescope
NIRSpec’s jitter Ultilossescompletein the spectral
event oflarger-than-expected
low-resolution
budge
or drift.prism mode may be a viable single-visit option for fainter targets. MIRI’s MRS mode uses a dichroic to simultaneously observe in all four
期待１：TRAPPIST-1惑星や同等のトランジットHZ惑星の透過光分光による
can m
usesacquirea Si:Asarray
is asub-band
successor
Spitzer’schannels, MIRI
but can only
spectraldata thatfrom one(A, B,to
or C)
at
IRAC
µm. three
IRAC’sredderchannel ment
any giventime.detectors
Therefore, atthe 5.8
MRSandmode8.0
requiresvisits forcomplete
大気組成の検出wavelength
coverage.
exhibiteda well-studied, time-dependent rising ramp (e.g., perfor
Harrington et al. 2007; Knutson et al. 2009; Agol et al. 2010; ments
ter. Although
itsarrayswere
later
replaced,the spacecraft
achiev
Stevenson
et
al.
2012)
while
the
bluer
channel’sfalling
ramp
期待２：中間赤外線領域で、HZ惑星の熱輻射検出?
(次スライド)still undergoes
thermalbreathing on∼ 30
min- past eratel
was similar
but less-studied.
Asaatimescale
result, weofcan
leverageutes.
Theprimarysourceofthesesystematics
is
the
large
therlength
experiencewith
Spitzertobetterunderstandandremovewhat
F IG . 2.— Recommended spectroscopic observing modes for transiting exenteringinandexiting Earth’s
host sare likelyexperienced
to be similarwhen
systematics
MIRI.
oplanets. For objects brighter than J ∼ 11, at least three visits will be re- mal gradients
JWST will there
not orbit
thetransit-spectroscopy-specific
Earth, so it will not undergo error capab
quired to achieve complete spectral coverage from 0.6 to 5.0 µm. NIRSpec’s shadow.Ultimately,is nolow-resolution prism mode may be a viable single-visit option for fainter tar- such extreme
temperature
but large spacecraft
budget for
JWST or variations,any of its instrument
modes. moUntil we betwe

James Webb Space Telescope への期待‣‣

主星光と惑星光のコントラスト
透過光の
シグナル

惑星光の
シグナル

中間赤外線では、ハビタブル惑星の熱輻射を
恒星光を打ち消すことなく検出できる可能性があるProxima Cen b: フェーズカーブによる大気特徴付け ?

trophysical Journal, 741:51 (9pp), 2011 November 1

RobinsonModeled thermal phase curves of Proxima Cen b
assuming habitable condition

Kreidberg, Leob (2016)10µm brightness temperature(~thermal emission)

bare rock(no heat re-distribution)

atmospherewith 30% heat re-distribution 3- inclinationとの縮退- JWSTのMIRにおける noise floor: 30-50 ppm ?
(2011)
1. True Robinson
color image of
the Earth–Moon system, taken as part of NASA’s EPOXI mission, compared to a simulated image using 10 µm brightness temperaturesur models. The spectra on the right show the corresponding flux at 10-pcさらに長期の観測だとトレンドが載る
from the Moon (gray), Earth (blue), and the combined (基準なしで補正難しい)
Earth–Moon flux (black), notFig. 2.— Thermal phase curves for a bare rock (left) and a planet with 35% heat redistribution. The models both assume an inclinationof 60 degrees and an albedo of 0.1.

physical and orbital parameters. We assumed a planet
radius of 1.1 R , based on predictions from the terrestrial
planet mass-radius relation (Chen & Kipping 2016). See§ 5 for a discussion of the uncertainty in the planet-tostar radius ratio.
Figure 2 shows the predicted thermal phase curves inng transit effects. The panel below the spectra shows the wavelength-dependent lunar fraction of the total signal. Images
and spectra
are5 for10a phaseof cases.
the wavelengthrangeµm forangle
the two
We calculate the phase curve directly from the tem-

28

The Astronomical Journal, 154:77 (8pp), 2017 AugustSnellen et al.

Proxima Cen b: 中分散分光を用いた吸収線検出 ?

The Astronomical Journal, 154:77 (8pp), 2017 AugustSnellen et al.

The Astronomical Journal, 154:77 (8pp), 2017 AugustSnellen et al.

Snellen+ 2017

igure 2. Planet model spectra (see Section 2.2) assuming a standard Earth atmospheric model for temperatures and gas mixing ratios, with, on the right, the assumed/p proﬁle. The upper panel shows the case where the stratospheric temperatures were artiﬁcially made isothermal and equal to the tropopause temperature
sothermal stratosphere model), the middle panel shows a model for clear sky conditions (clear atmosphere model), and the lower panel shows a spectrum for a caseith opaque high-altitude cirrus clouds (optically thick cirrus model).

Retrievalure 2. Planet model spectra (see Section 2.2) assuming a standard Earth atmospheric model for temperatures and gas mixing ratios, with, on the right, the assumed
p proﬁle. The upper panel shows the case where the −3
stratospheric−1 temperatures were artiﬁcially made isothermal and equal to the tropopause temperaturewasused
to simulate
spectraat 5×10cm
integration
time ofs, delivering
S/N ofthermalstratospheremodel), themiddle panel
shows a model
for clearresolution
sky conditions (clear
atmosphere model),and16.6
the lower
panel shows aan
spectrumfor 200
a casefor a total
5 (optically thick cirrus model).
h opaque high-altitude
cloudsat the simulated wavelengths).
corresponding
to cirrus
R>10
exposure time of 38.85 s. This extrapolates to an S/N ofHigh-pass filtered mock data-

Our spectral simulations used a standard Earth atmospheric
2000 hr−1 assuming that calibration uncertainties do not
JWST(6.5m),
MIRI(R~2000)
−3
−1
model
and atgas
mixingcm
ratios(McClatchey integration
contributethe s,noise.
s usedfor
to temperatures
simulate spectra
5×10resolution
time ofto16.6
delivering an S/N of 200 for a total
5 of which are shown in Figure 2. To
trrespondingal. 1972),積分時間:
theR>10
spectraplanetspectraare smoothed
to theat 5thedayssimulated wavelengths).
toexposureThetimemodelof 38.85
s. Thisextrapolates
to an S/Nof spectral−1
oundspectral
certainsimulations
extremes inthermalemission,runs were 2000 hrresolving
powerthatof R=2200
mean of thedoMRS
Ouruseda standard
Earth model
atmosphericassuming
calibration(theuncertainties
not 3B band)
erformedfor both clear(the “clear
atmospheredel for temperatures
andskygasconditions
mixing ratios(McClatchey
tothetonoise.
andspectrum
subsequently
binnedthe wavelength
steps of the
3BThe top-left panel shows the model spectrum convolved
Figure 3. contribute
From model
simulated MRSMIRI toobservations
assuming the clear-atmospheremodel.
Precisely
known(blackbody)
2.andPlanetmodel
spectra(see
Sectionassuming
acloudstandard
Earthatmospheric
for temperatures
andgasnormalized
mixingratios,
with,
on
the
the
assumedto theresolution
ofmodelthe
MRS
3B
channelofMIRI,
relative
to
star
ﬂux.
The
middle-leftpaneltheshows
associatedhigh-pass spectral signal andal.Figure1972),
thespectraof whichare 2.2)
shown
in Figure2.
Tomodel”)
foranopaquehigh-altitude
cirrus(located
at
The
model
spectraare smoothedtoaverage
theright,
spectralFigure3. Frommodel
spectrum
toplanet
simulatedThe
MRS
MIRI
observations
assuming
thetheclear-atmospheremodel.
Thespectrum
top-leftpanel showsmodelits
spectrumconvolved
MRS
channel.
resultingclear-atmospheremodel
T/p proﬁle. The upper panel shows the case where the stratosphericthetemperatures
wereartiﬁciallymadeisothermal
and
equal
to 5×24
thestartropopause
temperaturetothe resolutionof theMRS
3B channel
ofasMIRI,
normalized
relativeto the average
ﬂux.
The
middle-left
panel
shows
its associatedhigh-passandlower-left
panel
with
noise
added
if
Proximawasobserved
for
hr.
The
top-right
panel
shows
the statistical
1, spectral
2, and signal
3σ conﬁdence
intervals when no
undcertainextremesinthermalemission,
model
runs
were
resolvingpowerofR=2200
(the
mean
of
the
MRS
3B
band)
star
spectrum(*).2(isothermalbar, nearstratosphere
the tropopause
Muinonenetal.
1989),
calledthenormalized
thewasandstellar(which,forisand 3σ conﬁdence intervals when nomodel), the middle panel shows a model for clear
conditions(clearatmospheremodel),
the lowerpanel
a spectrum
fortheaclarity,casethelower-leftpanel
with noise
added as byif Proxima
observed
forspectrum
5×24hr. shows
The top-rightpanel shows
statistical 1, 2,COsky
2 signal is present, pointing to a 3σ upper limit of 50 ppm for the planet/star contrast. The bottom-right panel shows the same, but with the simulated CO2 signal
signalis present,
pointingTheto a 3σ
upperindicates
limit ofto
50the
ppm
forwavelengththeplanet/star contrast
contrast.
Thebottom-rightpanel shows
the same,
but withspectrum
the simulated
CO signalCO
formedforthickboth
clear sky(the“clearatmosphereandsubsequently
binnedthestepsofthetop3Bwith opaquehigh-altitude
cirrusconditions
clouds(opticallythickcirrusmodel). present—detected
optically
cirrusmodel”.
Also,
to explore
a situation
with
at∼3.8σ.
y-axismean
planet/starof the(non-spectrally
ﬁltered)template
between13.2 and 13.5 μm. Note thatassumedtobefeatureless)
is
shown
in
the
panel
of
present—detected at ∼3.8σ. The y-axis indicates the mean planet/star contrast of the (non-spectrally ﬁltered) template spectrum between 13.2 and 13.5 μm. Note thatthesesimulations
thatbothboththeinstrumental
spectral
response
andmodelthe stellar
at a wavelength-to-wavelengthscale whichhave been
determined, which isdel”)thermal
and for an
opaquebetween
high-altitude
cirrus cloudat
MRS assume
channel.
The
resultingclear-atmosphere
spectrumthese
assumethat
theBecause
instrumental
spectral
response
and the stellar
spectrumata spectrum
wavelength-to-wavelengthisarge
contrast
the surfaceand (locatedstratosphere,
asimulations
Figure3.
thespectralﬁltering
techniqueis onlyscale have been determined,
2

2

29

flux [W/m2/µm]

中間赤外線領域の晩期型星のスペクトル
10

4

10

3

晩期型星の中間赤外線領域のスペクトルにはさまざまな吸収線が入り、（非トランジット系では）それらを6000 K
4000 K
2600 K

100 ppm以下の精度で決めるのは難しい（私たちはいつも惑星と主星を合わせたスペクトルを
観測する）

102

101

100

Based on BT-Settl model(Allard+ 2012)
7

8

9 10

20
wavelength [µm]

YF & Matsuo under review

主星-惑星をスペクトル上で分離するには、
3040

公転速度によるドップラーシフトが十分分解できる
R >~ 10000 でないと厳しいRef) 恒星スペクトルモデルの不定性

TRAPPIST-1の近赤外スペクトルを
いくつかの温度の理論スペクトルで説明 (黒点/白斑を考慮)=> poor fit

Zhang+2018, Wakeford+ 2019

思想
‣ JWSTでは、TRAPPIST-1のような、太陽近傍の超晩期型星周りのトランジットハビタブ
ル惑星の透過光分光観測が期待される。

• そのようなターゲットは今のところ1つ (今後も探査は続く e.g., exo-Jasmine )‣ 他の可能性：TMTなどの地上大型望遠鏡による高コントラスト＋高分散観測• Proxima Cen bが良いターゲット。それ以上はInner working angle次第。反射光を見る。

‣ それぞれに challenging。
‣ もう少し先へ向けてハビタブル惑星の観測をキーサイエンスとしたミッションが検討中‣ 幅広くいろいろな観測の可能性を考えておく (装置の設計にも理論モデルにも影響)中間赤外線高分散分光を用いた
HZ地球型惑星大気観測の検討

中間赤外線高分散分光
きっかけ考えられるメリット‣ （中間赤外）大気分子の吸収が強い‣ （中間赤外）温度・半径の情報‣ （高分散）吸収バンドがオーバーラップした分子でもロバストに検出できる‣ （高分散）たくさんのシャープな線によって、装置の系統誤差が大きくても検出できる可能性がある大気モデルと高分散スペクトル‣ 地球の混合率を仮定‣ 温度構造は右図のように単純化‣ 雲なし‣ 水平方向の依存性は考慮せず‣ ライン: HITRAN2016‣ 散乱無し‣ 計算にはGoogle colabのGPUを使用

4

Figure 2. The assumed vertical profiles of the35

Examples of high-resolution planetary spectra in MIR

radiance [W/m2/µm/sr]

ラインの数:

O(100)

ラインの数:

O(100)

N 2O

CO2

H 2O

1ppm NH3
wavelength [µm]O3+wavelength [µm]
YF & Matsuo, under review

—

—

36

Examples of high-resolution planetary spectra in MIR

radiance [W/m2/µm/sr]

ラインの数:

O(100)

ラインの数:

O(100)

N 2O

CO2

H 2O

1ppm NH3
wavelength [µm]O31ppm PH3+

PH3

wavelength [µm]
YF & Matsuo, under review

—

—

観測
惑星スペクトル

惑星視線速度-50 km/s

恒星スペクトル +公転周期~10日
波長 [µm]観測者
Star

惑星スペクトル( 主星の視線速度変動 ~ 1 m/s )恒星スペクトル +波長 [µm]+50 km/s

観測
惑星スペクトル

惑星視線速度-50 km/s

恒星スペクトル +公転周期~10日
波長 [µm]

平均からの差分を検出
幅の広い吸収線は、差分を取るとfeatureが一部削られる
波長 [µm]惑星スペクトルStar( 主星の視線速度変動 ~ 1 m/s )恒星スペクトル +波長 [µm]+50 km/s

観測シミュレーションの設定

Figure 5. Same as Figure 4 but with O3 .

‣ 恒星スペクトル：BT-Settl (Allard+ 2012)‣ ノイズ：恒星光ノイズ、Bose factor, 黄道光
ノイズ, readout noise, dark current‣ 波長分解能R=30,000‣ 口径 6.5 meter (スケーリング可能)‣ end-to-endの透過率: 0.2‣ 系統誤差なし (あとで議論)

Table 2. Assumptions for the host star and theplanetary
orbit.
ターゲット星のプロパティ
description
mid-M
late-M
star radius (R? )0.14 R
0.10R
star temperature (T? )3000 K
2500 K
planet/star flux ratioa ⇠ 70 ppm ⇠ 200 ppm
planet orbital radius
0.0485 au
0.0146 au
planet orbital period
11.26 days 2.27 days
planet orbital velocity 46.83 km/s 69.70 km/s
aEvaluated at 15µm, assuming a black body
spectrum with 288 K for the planet spectrum.

8観測条件

description
value of theularfeatures in thesymbol
combined spectrum
spectralR star and the
30,000star
and resolution
the planet. The
planet are
telescopediameter
6.5 msuppresnot
spatially
resolved,Dand no starlighttotal instrument
throughput is assumed.
⇠
0.2
sion
The assumptionsformocktoobservations
aredistance
targetd and noise estimate
5 pc
given
in Section
followed
planetaryradius 3.1, Rwhich
Ris(=
6.371 ⇥ by
106 them)p
descriptionof our analysis
procedures
exposure time⌧exp
1800 secin Section 3.2. The results of parameter constraints
Table 3. Fiducial values for observational paramare discussed in Section 3.3.
eters. The scaling of the required integration time

The hoSettl mosponding
ing log g
update ftectabilit
tauri b w
ation (Sn2016).
The di5 pc, althscaled by
The incla↵ects thsight vel
noted.

3.1

a fringe ofMock
fectly remo
high-reso
Gaussian r
ing powedard devia
power iswavelengthof SPICA2µm) (e.gj = Nstar,
pected r
220
et al. read
where
the factor
analy
Bose

mid-M s

mid-M

コントラストと軌道傾斜角の制限
15 µm CO2恒星スペクトルが不定の場合

積分時間: 1.2 days

1.2 days3.5 days
積分時間:

5 days

5 days 9 days

光子ノイズ: 600 ppm

late-M starlate-M star

恒星スペクトルが既知の場合

光子ノイズ: 350 ppm

有意性
8 days3.5 days
68.3%
12 days 9 days
95.5%
99.7%

より現実的な場合（右）では3倍程度時間がかかる

正解

Thermal inversioさまざまな分子種

mid-M star

14
CO2 (13-17.5 µm)

H2O (10-24 µm)

O3 (9.3-10.1 µm)

14 days
60 days

14 days
60 days

50 days
60 daysO3 (12.5-16 µm)
30 days
40 daysσ~110 ppm

3.5 days
9 dayslate-M star

68.3%
95.5%
99.7%
N2O (16-18 µm)
24 days
32 days

σ~80 ppm

3.5 days
9 days17 days
23 daysσ~350 ppm

10 days
14 days

10 days
14 daysσ~170 ppm

Thermal inversion入りモデル

Ref) 地球の熱輻射スペクトルの場合
Thermal inversion入りモデル
Figure 9. Same as Figure 8, but the mock data are now analyzed by analysis (B).
CO2のみ考慮, 主星が5pcのlate-Mの場合
with thermal inversion ("US standard”)

6 days9 daysFigure 10. Left: high-resolution thermal emission spectra assuming the vertical profile of temperature andCO2 sameas Earth’s “US standard” profile. Right: Prameter constraints based on the left panel.
- 成層圏温度が高くなるとfeatureの深さは減少。

panel ofFigure 10, assuming
a late-M host star.
bands may overlap. With high-resolution spec- ただし、Thermal
inversionはシャープな輝線を形成するため、検出に有利にWhile thetotal integration time required for detroscopy where the fine structures of the specはたらく。軌道傾斜角も決まりやすい。
tection is longer than our fiducial case due totral shapes are resolved, it is possible to detect
the reduced line depths (corresponding to theeach molecule separately, although the signals

mid-

9.8 µm v.s. 14.5 µm O314

CO2 (13-17.5 µm)

H2O (10-24 µm)

O3 (9.3-10.1 µm)

3.5daysdays
3.5daysdays
1414- O3は狭い波長域に多数のラインがあり高分散分光向き
days
days
60 9days609daysO3 (12.5-16 µm)

17 days50
23 days60

1030daysdays
14
40daysdays

mid-M starlate-M star

- CO2が少ない大気では、コントラストで得をする14.5 µm

N 2O

O3 のバンドの方が検出しやすくなる- そんな大気があるか ?
- cf) 海惑星の海底風化 (Nakayama+ 2019)
Thermal inversion入りモデル
8

O3
CO2 0.1ppm

7

4
3
2
1
0
12

12.5

13

13.5 14 14.5 15
wavelength [µm]

15.5

16

16.5

5
4
3

4
3
2

1

1
12

12.5

13

13.5 14 14.5 15
wavelength [µm]

15.5

16

16.5

6

5

2

0

O3
CO2 10ppm

6
radiance [W/m2]

radiance [W/m2]

late-M starradiance [W/m2]

6

5

8

O3
3.5
days
3.5
days
17
days
10
days
CO
ppm
Figure9.
Sameas
Figureanalyzedby(B)
2 330analysis
7 8, but the mock data are now
7
9 days9 days23 days
14 days

7

6

8

O3
CO2 1ppm

radiance [W/m2]

8

0

5
4
3

6 days9 days2
with thermal inversion ("US standard”)
1
12

12.5

13

13.5 14 14.5 15
wavelength [µm]

15.5

16

16.5

0
12

12.5

13

13.5 14 14.5 15
wavelength [µm]

15.5

16

16.5

Ref) 恒星スペクトルのvariability
K2 (0.43–0.89µm) と Spitzer (4.5µm) によるTRAPPIST-1の時間変動のデータ
The Astrophysical Journal, 857:39 (9pp), 2018 April 10

et al.
Morris+Morris(2018)< 100 ppm

Figure 5. Maximum-likelihood model ﬁt for bright spots on TRAPPIST-1 (red curve), ﬁtting the Kepler (left) and Spitzer (right) light curves (black points)
simultaneously.

太陽近傍の低質量星
HZの付近に地球サイズ惑星の存在が検出されている恒星※恒星パラメータはGaidos+ 2014に基づく

惑星

惑星質量
[Me]

距離

Sp

恒星
有効温度

Proxima Cen

1.3 pcM72883 Kb>1.17

Ross 128

3.4 pcM53145 Kb>1.35

Gliese 1061

3.7 pcM6

3000 Kd>1.7

Lyten’s star

3.8 pcM43317 K

b

>2.89

b

>1.05

c

>1.11

その他のものも含めると、
5 pc以内にlate-M (M6 or later) 7個
mid-M (M3-M5) 10個

c

>4.4

(Gaidos+ 2014)

Teegarden’s star

3.9 pcM72700 K

GJ 682 b

5.1 pcM53190 Kターゲットごとの見積もり
MIR HR for HZ planets
13‣

ざっくりしたスケーリング
dle panel of Figure 6) becomes
portional to:
✓
◆ 2✓
◆2 ✓
◆ 2✓ ◆ 1
he overall shape of the di↵erenRpdD⇠ains similar. The reduced di↵er⌧esp, 0 /
R
5
pc
6.5
m
0.2
an therefore be compensated by
(13)
積分時間
惑星半径
距離
望遠鏡口径
効率
to reproduce a similar signal.
In particular, the observations of the nearest
hand, integration time required
possible target Proxima Centauri b, at the disis less sensitive to the assumptance of 1.3 pc and with the estimated radiusar spectrum.
Furthermore,
ProximaCen b (1.3thepc, 主星M5, 1.1Re) の場合
Teegarden b (3.86 pc, 主星M7, 1Re) の場合
of ⇠ 1.1R , would require ⇠ less than 1 days
ween the inclinationand the con-days
6.5m望遠鏡では、<1dayseven without an assumption on 6.5m望遠鏡では、~3
the stellar specgnificant for O3 . This is because
trum. Adays
SPICA-like specification
with 2.5 m
2.5m望遠鏡で透過率0.1では、~102.5m望遠鏡透過率0.1では、~45
days
re well packed and the broad line
and throughput 0.1 implies ⇠10 days of total
ed by adjacent lines (Figure 7)
integration time.
shift moves beyond these narinclination 望遠鏡の系統誤差を<100ppmに抑えることができれば、
is well constrained
4. DISCUSSION
presence is inferred. Similarly,
4.1. Variety of Planet Thermal Emission
6.5 m望遠鏡では5pc以内程度の天体の中間赤外における最も強いラインの検出可能性あり。
also better at constraining theSpectrum（2.5weakm望遠鏡では少し厳しい。がんばってProxima
Cen bのみ。）
pite its relatively
features.
4.1.1. Stratospheric Thermal Inversion
und mid-M stars in particular,This section discusses the possible variety of
can in fact yield constraints on

0.1 µm

1.0 µm

10 µm

高層大気 分子の吸収

波長

分子の吸収
熱輻射 (温度構造)

反射スペクトル

JWST
透過光分光
惑星食分光

TMT
近傍のM型星周りのハビタブル惑星OSTLUVOIR, HabEx
太陽型星周りのハビタブル惑星

透過光分光
惑星食分光
近傍のM型星周りのハビタブル惑星

まとめ‣ 恒星と惑星を同時に観測した中から惑星光を取り出すさまざまな手法や多波長観測を用いて、系外惑星大気の多角的な特徴付けが進行中- トランジットや食の利用、フェーズカーブ、高分散分光- 可視 vs 中間赤外

‣ 木星型惑星の大気については、サンプルが蓄積し、一部で統計的な議論が可能に- 雲についての示唆など‣ HZの地球型惑星大気までもう一声‣ 中間赤外線を利用したHZ惑星の熱輻射とその中の分子の検出は、晩期型星の周りでは主星の光を打ち消す装置なく検出できる可能性がある