THE STAR FORMATIONNEWSLETTER
No. 315 (#34-43)
塚越崇(NAOJ)
24 May. 2019

1. The nitrogen carrier in protoplanetary disks
2. Observing the gas component of circumplanetary disks around wideorbit planet-mass companions in the (sub)mm regime
3. A high resolution mid-infrared survey of water emission fromprotoplanetary disks
4. ALMA High Angular Resolution Polarization Study: An Extremely YoungClass 0 Source, OMC-3/MMS 6
5. Polarization of stars with debris disks: comparing observations with models
6. The survivability of planetary systems in young and dense star clusters
7. Spatial segregation of dust grains in transition disks. SPHERE
observations of 2MASS J16083070-3828268 and RXJ1852.3-3700
8. Ejection History of the IRAS 04166+2706 Molecular Jet
9. Luminosity outburst chemistry in protoplanetary discs: going beyond
standard tracers
10.External photoevaporation of protoplanetary discs in Cygnus OB2: linking
discs to star formation dynamical history

Observing the gas component of circumplanetary disks around wide-orbit planet-masscompanions in the (sub)mm regime
Rab, Ch. et al.

• Wide-orbitのplanetary mass companionで周惑星円盤が受からないのはなぜか？

• モデルと放射輸送計算を行い、ALMAの感度と比較• サイズが大きい周惑星円盤だと分子輝線観測で受かるかもA&A proofs: manuscript no. Rab_

• Wide-orbitのPMCを想定してモデル化• 原始惑星系円盤に”埋もれていない”状態• power-law density prof., disk/planet mass ratio=1%,
Rout=1/3 Hill radius• CSChaのALMA non-detectionを念頭に[Wu+2017a,Ginski+2018]• PRODIMO codeによる化学&放射輸送計算• Synthetic observables• ALMA Band6-10 & Band6 of ngVLAを想定

モデルパラメータ

Table 1. Main parameters for our reference model.

Quantity
companion masscompanion e↵ective temp.
companion luminosity
strength of interst. FUV
disk gas mass
dust/gas mass ratio
inner disk radius
tapering-o↵ radius
column density power ind.
reference scale heightflaring power index
min. dust particle radiusmax. dust particle radius
dust size dist. power index
max. hollow volume ratiob
dust compositionc(volume fractions)
inclination
distance

Symbol
Mp
Tp
Lp
ISM

Md
d/g
Rin
Rtap
✏
H(1 au)
amin
amax
apow

Vhollow,max

Mg0.7 Fe0.3 SiO3
amorph. carbon
porosity

i
d

Value
20 MJ
2500 K
10 2 L1a
0.2 MJ
0.01
0.007 au
1 au
1.0
0.1 au1.15
0.05 µm
3 mm
3.5
0.8
60%
15%
25%
45
150 pcdv=1km/s• ALMA/CASA simulation, 6hours integration,
Ch. Rab et al.: Observing the gas component
CPDs
the (sub)mm
regime definitions, see Woitke et al.
Notes.ofFor
moreindetails
on the parameter

放射領域と放射サイズ

Ch. Rab et al.: Observing the gas component of CPDs in the (sub)mm regime(2009, 2011, 2016). (a) ISM is given in units of the Draine field (Draine& Bertoldi 1996; Woitke et al. 2009). (b) We use distributed hollow
spheres for the dust opacity calculations (Min et al. 2005, 2016). (c) The
optical constants are from Dorschner et al. (1995) and Zubko et al.
(1996).

planet and is therefore not part of the disk. However, analytical and numerical models suggest that CPDs might be truncatedat 0.3 0.4 rHill owing to tidal truncation e↵ects and/or the low
specific angular momentum of infalling material during their formation (e.g. Quillen & Trilling 1998; Martin & Lubow 2011;
Shabram & Boley 2013). Although this more strict requirement
does not necessarily apply to wide-orbit PMCs, i.e. which likely
depends on their formation mechanism, we used rHill /3 as a reference quantity for the outer radius of our CPD models.
For the reference model we chose a tapering-o↵ radius ofFig. 1. Spectrum for our reference PMC with Mp = 20 MJ . The solid Rtap = 1 au, resulting in an outer radius of rout = 10.9 au. This
black line indicates the photospheric (intrinsic) spectrum and the blue value is equal to r /3 for a PMC with M = 20 M orbiting aHillpJline shows the spectrum with the added accretion luminosity.
solar-mass star at an orbital distance of a = 178 au. We present
models with smaller and larger disk sizes and also discuss thedi↵erent apparent (observed) disk sizes for the dust and gas in
12
Fig. 3. Left panel: Location of the main emitting region for the near-infraredand(sub)mmdust
emission
and
the
CO
J
=
3
2
line
in
the
reference
2.1.1. Planet-mass companion
detail in Sect. 3.1. We determined the inner radius of our disk
model. The vertical lines of the boxes denote the radii where the cumulative flux in radial directions reaches 15% and 85%, respectively. The
models by the dust sublimation radius, where the dust temperahorizontal lines denote the heights where the cumulative flux integrated
verticallyreaches 15%of
andthe85%
at eachradialposition.
The coloured
To modelthe irradiation
disk
by thePMC
we needto ture reaches ⇡ 1500 K.
contours show the CO number density and the dotted grey line shows
the visualextinction
AV reaches
Right panel:
Normalized
knowwhere
the mass
Mp , theluminosity
Lp , andunity.
the e↵ective
temperaThe density structure and surface density profile for our refradial intensity profiles (azimuthally averaged) for the dust continuumShown
are profiles
derivedfrom images
tureandT p .line
Forimages.
the known
wide-orbit
PMCs
typicalvalues convolvedof about withrepresentative beam sizes for SPHERE (blue, FWHM = 0.03100 ), ALMA (black,
FWHM = 0.01200 ), and ngVLA (red, FWHM = 0.00400 ). erence model are shown in Fig. 2; the gas and dust temperature
Mp ⇡ 10 20 MJ , Lp ⇡ 10 2 10 3 L , and T p ⇡ 1000 2500 K structures are shown in Fig. B.1. All relevant parameters for our

基準モデルでは全てoptically thick

円盤密度構造PRODIMO psuch as spe
butions (SEFor thechemical sp
phase chem
(freeze-out;
chemical ne
Kamp et al
did not incl
the planet.
UMIST 201
et al. (2017ready stabl
databases. T
curate treat
cal heating/

2.3. Synthe

We produce(sub)mm rewavelength
Bands 6 toallowed us
the observa
on the spec
strongest em
calculations
Atomic and2010), He (akrishnan e
To testthem to theservations i
pendix C).
for the synth
simulations
Most of150 pc (e.g.
presented o
for our moduse i = 80

3. Results

3.1. Obser

First we dismodel, whi
subsequentnear-infrare• 各Bandでの放射強度と感度• COもダストも円盤サイズ↑で強度↑• ALMA Bandsでは円盤質量への依存性弱• 光学的に厚いため
Ch. Rab et al.: Observing the gas component of CPDs in the (sub)mm regime
CSCha観測上限値[Wu+2017]

5. Comparison of CO line and continuum peak fluxes from CPD
PMC光度と放射場に対する変化• Fig.
models with varying structure properties (radius and mass) to 5 detec12

tion limits
for one band of ngVLA and various ALMA bands (see ApダストはPMC光度にsensitive
pendix C for details). The dust emission is reported at the same wavelength放射場の変化はCOのみ↑
as the corresponding line. A beam of 0.100 ⇥ 0.100 is assumed.
For the lines a channel width of 1 km/s was used. The peak flux is the
maximum flux value in the synthetic images; for the lines all channels
are considered to evaluate the peak flux (see Appendix D for details).
The diamond symbols show the fluxes for the lines and the star symbols
Fig. 6. Same as Fig. 5 but for models with varying dust properties. The Fig.
Same as Fig. 5continuum
but for models
withvaryingthe7.corresponding
fluxes.
The
black radiation
solid andproperties
dashed lines
underlying model is always the reference model with Lp = 10 2 L , of indicate
the PMCtheandtheenvironment(backgroundfield).
The underlying
sensitivity limits for the line and continuum,
respectively.
Md = 0.2 Mamaxはあまり影響がない
disk
structureis alwaysthat fromthe reference
model. levels as reported
J , and rout ⇡ 11 au.
Thelight greybox shows
the rangeof 3 sensitivity
submm放射は光学的に厚く、内側の by Wu et al. (2017a) for their ALMA Band 6 continuum non-detectionsof a sample of wide-orbit PMCs.

••(Rtap =
but cawith thmodelsr ⇡ 1
disk isrout .
is smaet al. (results
the anaradiald/g =disks (
To
lines crout ⇡
is best
might
becom

Ch. Rab et al.: Observing the gas component of CPDs in the (sub)mm regime• ダスト進化の依存性••

暖かい領域で決まるinner
regionsion is not significantlya↵ected although the hot
emission
at most a↵ected line. The continuum fluxes drop by factors of 1.9
those wavelengths should be dominated by large dust grains to 3 compared to the reference model, with the lowest change
ダスト質量を減らすと円盤質量&サイズ↓• 100(agrain &µm). In the models presented in this work, the in but
the arengVLAband and thehighestchangein models
ALMA theB10.best
TheS/Nnot detectablein thelines.
For allに伴いフラックス↓(sub)mm dust
emission is dominated by optically thick emission dust
moresensitive
the opticalemissionof the theis temperaturereached withis thengVLA
andto ALMA
Band
7, whereas
and the sensitivity to the dust grain size is therefore lost. In the companion
and is thereforebands
more (Bandsa↵ected8byof the for
ALMA high-frequency
to the10)decrease
are not suited
COにはほとんど影響なし
amax =•1 µm model, the 860 µm emission is optically thick up to photosphericluminosity.
For
the
gas,
the
far-ultraviolet(far-UV)
detections of CPDs such as those modelled in this work.
r ⇡ 1.7 au and the average dust temperature in the emitting re-3.3. D

Similalikelying inand cobrown
of Zhucess iseral) thlower
Zhu etare sigthat atdust poWe

A&A proofs: manuscript no.

• CSCha(c)との比較••

solvedstraint
Dust disk model [Ginski+2018]disk. W-3
as our
• Mp=20MJ, Lp=2.5x10 L⦿,whetheMcpd=2x10-4MJ, amax=1um
As
light
im
Ginskiモデルだと6時間積分でも検出はsize. W
難しい
using aaxis a• tapered edgeにすればCOなら検出
mass o
できるかもしれない
ski etestima• ダスト検出は難しいが、ダストサイ
tent w
ズが大きければ受かる可能性がある
tance e179±1
A&A proofs: manuscript no. Rab_wocpd(GAIA
from tFig. tion
8. SameFig.disk
5 butaroundfor the CS
Cha companion
model.
of aasgas
a wide-orbit
PMC.
TheShown
flexibleCollab
mode
are theoriginal
Ginskietal.
(2018)
model
andseveraltapered-edge
Cha ispresented in this work allows us to make accurate prediction
models with di↵erent dust-to-gas mass ratios (d/g) and dust composi- have afor both the dust and gas observables of the potential CPD cantions/populations (large grains).
tancesdidates. As observing such CPDs will be time consuming
eve
model
with ALMA or ngVLA, such models are crucial to optimize observing
dependingbackgroundon the known
properties
ofInthTo
study programmes
the impact of enhanced
radiation,we
CPDm
軌道半径が大きければ大きい円盤が期• Fig.
showcompanionin
7 models
10 and 100 times stronger ISM
and itswith
hostastar.
ters ofradiation field ( ISM = 10, ISM = 100) and a model where we
model
待できるので検出される可能性がある
assume that the additional stellar radiation heats the dust disk to a
4.2.
What
can we
learn
from
gas
observations
minimum
temperature
of T
25 K.
Relativeto the reference we als
d,min =
Rout>10auくらいで、円盤全体で光学•
model the line fluxes increase in all three cases whereas the dust dust pr
used b
4.2.1.
Companion
andwith
diskTproperties
is onlya↵ectedin the model
的に厚ければ受かる
d,min = 25 K. We note thatfor the ISM = 100 model CO is efficiently photodissociated in origina
Current continuum
observationsof Jwide-orbit
PMCs indicat
-4M
12
COならM
くらい
the outer disk and
CO J = 3 d~10
2 emission
becomes optically the lin• theregimesmall
dust
disks,
eitherbecausetheyareunresolved(Ginskiet athin at r ⇡ 9 au compared to r ⇡ 10 au in the reference model.
Ginski2018)theor higher
not detected
at all (Wueasily
et al. 2017a).
However,
However,gas temperaturescompensate
for the in a com(sub)mpactemitting
CPD thedust might evolve on much shorter timescales
thasmallerregion.
The re
in
T
Tauri
disks
and
rapid
radialmigration
can
lead
to
stronglGenerally speaking, the size of the CPD has more impact
clude a• CPD検出のfeasibility

A high resolution mid-infrared survey of water emission from protoplanetary disks
Salyk, C. et al.

• CTTS&TDに対するH2O輝線の高分散中間赤外分光• CTTSの方が検出率高い• スペクトル形状から放射領域を推定• 円盤 + Wind• 将来観測への布石: JWST(励起を多くする)や30m望遠鏡(SNR向上)

a Radial velocity shift between source and Earth: the sum of intrinsic and Earth-induced radial velocities.
b Signal-to-noise ratio of divided (unbinned) spectrum in the wavelength range 12.252 − 12.261 µm, 12.38 − 12.392 µm or 12.46 − 12.48 µm,

depending on the grating setting.
c Setting extends from ∼12.370–12.462 µm.
d 12.356–12.416 µm
e 12.233–12.296 µm
f 12.424–12.482 µm• MIR spectrographs on Gemini North

7• TEXES (R~100000, dv=3km/s), Michelle (R~25000, dv=12.5km/s)• 8 CTTSs & 3 TDs

TABLE 2
Source Properties

TEXES spectra

M⋆
L⋆
R⋆
d
i
log (Ṁ )log(Lacc )a
vhel
Refs
◦
−1
[M⊙ ] [L⊙ ] [R⊙ ] [pc] [ ][M⊙ yr ][L⊙ ][km s−1 ]DK Tau0.52
0.54 1.53
129 41 ± 11−8.3/−6.4−1.2/0.72
16.3
13,14,17,20
DoAr 441.4
1.3
1.7
146 25−8.4/−8.0−0.88/−0.45 −5.9
2,3,13,14,18,26
DR Tau0.4
3.00 3.19
196 37 ± 3b , 9c
−6.68/−5.1 0.44/1.3
21.1,27.6 5,13,14,16,17,20,23
FZ Tau0.57
0.50 1.4
130 38 ± 15
−7.3
−1.3
18.0
6,12,13,14,21,24,25
GI Tau0.60
0.33 1.09
131 83−8.2/−6.9−1.2/−0.62
17.3
13,14,17,20,27
HL Tau1.3
3.0
3.0
140 46.72 ± 0.05 −8.8/−7.1−1.3/0.41
19
1,8IRAS 04303+2240 0.43
1.49 1.2
148 ...
−6.1
0.93
19.1
4,10,13,14
RW Aur
1.48
1.60 1.8
152 45 − 60
−6.7/−6.0
0.87/1.6
15.9
7,9,12,13,14,15,24
SR 24 N
1.1
2.5...
114 121 ± 17
−6.9
0.03
−4.6
11,13,14,19
SR 24 S
1.5
2.5...
114 70 ± 5
−7.2
−0.23
−5.8
11,13,14,19
T Tau N
2.11
7.42 3.30
144 25 ± 5
−7.5
0.22
19.2
17,13,14,20,22(1) ALMA partnership et al. 2015 (2) Andrews et al. 2009 (3) Andrews et al. 2011 (4) Andrews et al. 2013 (5) Ardila et al. 2002 (6) Banzatti
flux calibrated,
fluxes are& Pontoppidan 2015 (7) Basri & Bertout
1989 (8)et al. 2006 (11) line
Fernandez-Lopezet computed assum
TABLE
3 Beck et al. 2010 (9) Dai et al. 2015 (10) Edwards
al. 2017 (12) Gahm et al. 1999 (13) Gaia Collaboration et al. 2016 (14) Gaia Collaboration et al. 2018 (15) Hartigan et al. 1995 (16) Isellaµm continuum
valuesgiven in Table 4.
Detected
et al. 2009 (17) Johns-Krull & Gaﬀord
2002 (18)Lines
Melo et al. 2003 (19) Natta et al. 2006 (20) Nguyen
et al. 2012 (21)flux
Pascucci
et al. 2015(22) Podio et al. 2014 (23) Pontoppidan et al. 2011 (24) Rigliaco et al. 2015 (25) Salyk et al. 2013
Salyk et al. 2015(27)Tauri
Simon etal.
flux(26)
variability
in T
stars
of 20–50% (Kλ0 aEupper Spin
2016JKa Kc
2012)
in a corresponding uncertainty
GM⋆ Ṁ
a From provided references, when available, or calculated using Lfrom Gullbringet al.results(1998)
acc = 0.8 R(upper → lower) [µm][K]⋆b From 1.3mm continuum image. Value used in our line shape models.
this uncertainty is diﬃcult to quantify witho
175 near-IR
12.2481 5795
para
c From13 → 16
2 14
spectro-astrometry.
photometric study of individual targets.
187 12 → 174 13
12.2654 6954
ortho
Name

d
d

d

• 検出: 6/8 CTTSs & 1/3 TDsddperience
types 4048
of datasets.
Withinest continuum fluxes, As
andshown
relatively
mid-IRwater
inweakTable
3, theemission lines at 1
128 5 working→ 115 6with these
12.2772
ortho
our chosen sample, water is detected in 6 out of 8 classiemission fluxes (Najitaetal.
2013).
Figure2shows
that
12.4070 µm, which are visible in the TEXE
151 and4948
orthodisks
13 →14 1 12.3757
cal 16
T 4Tauridisks
out of 3 (pre-)transitional
DK Tau’s non-detection is likely due to the low SNR of
have
excitation
temperatures.
They
174 category
12.3962ortho
13 → 163includes
14(whichDoAr 44, 5781
SR 24 N andSR 24
these data, while GI
Tau’ssimilar
non-detection
is likely due
to
16
→
15
12.40704945
para
similar
Einstein
A 2coeﬃcients,
S). For3 theits low line flux. As
described
in Section
and Figure but since the 1
13 classical2 14T Tauri disk RW Aur, emission was
detected
datafrom 2013,prior data1, all 8 of the classicaland onethe transitional
118 3 in→our
10TEXES
12.44483629and inortho
5 6
anddisks,12.4070µmof lines
are ortho and para wat
from
but
not
in
our
Michelle
data
from
2011;
this
disks
(DoAr
44)
were
known
water
emitters
132007,→1212.4535
4213
ortho
7 6
4 9
lines, respectively, theybased
haveona degeneracy rat
is discussed
further in Section 3.4. The classical disks
Spitzer-IRS spectra. Similarly, AS 205 and RNO 90, witha Theoreticalline
HITRAN
database
optically
therefore, the expecte
with non-detections,
DKcenters
Tau andfrom
GI Tau,have the molecular
lowVISIR-detectedwater
emission,thinwere emission,
known Spitzer-IRS(Rothman et al. 2013).
would be 3. Given the error estimates in Tablenn• ortho/para下限値~3dwater emitters with bright (7 Jy and 2.1 Jy, respectively)
mid-infrared continuum flux levels (Pontoppidan et al.

possible to confirm whether the line ratios arewith optically thin or optically thick emission

Rotation diagram

TABLE 5Model fits

High-Res
Spitzer
SourceT [K]log Nfixed−Ω a [cm−2 ] log Rfixed−N b [AU] T [K]log N [cm−2 ] log Rc [AU] Spitzer Ref
DoAr 44534 ± 36
18.23 ± 0.22
0.12 ± 0.11
450
18.3
-0.036
2
DR Tau (T)726 ± 13718.04 ± 0.37
0.025 ± 0.19
500
18.6
0.18
2
DR Tau (M)725 ± 4618.30 ± 0.17
0.15 ± 0.09
FZ Tau (T)748 ± 5417.83 ± 0.18
−0.09 ± 0.09
550
19.0
−0.15
3
FZ Tau (M)1014 ± 161 17.36 ± 0.27
−0.32 ± 0.14
HL Tau396 ± 89
20.41 ± 0.60
1.2 ± 0.3
750
19.0
−0.097
3
IRAS 04303+2240 425
19.8
0.9
250
20.9
0.43
3
RW Aur591 ± 300
18.7 ± 2.0
0.35 ± 1.0
600
18.3
0.20
1
T Tau N479 ± 156
19.8 ± 0.7
0.9 ± 0.4
700
19.6
0.0
3
AS 205520 ± 75
19.36 ± 0.54
0.68 ± 0.27
300
20.6
0.34
2
RNO 90595 ± 160
18.11 ± 0.99
0.06 ± 0.50
450
18.3
0.23
2
(1) Carr & Najita (2011) (2) Salyk et al. (2011a) (3) This work.
Ω0a Not a true column density. The true column density, N = N
22
fixed−Ω Ω where Ω0 = π(1 AU) /d and Ω is the true emitting solid angle.
For Nfixed−Ω ! 1019 cm−2 , N must be < Nfixed−Ω for lines to remain optically thin.
b Emitting radius assuming N = 1018 cm−2 and Area = πR2 .
c Assuming distance given in Table 2• Trot=400-1000K• Spitzerデータより高く出る=>波長分解能をあげるとコンパ
クトな高温領域をトレースする

スペクトルを波長補正後concatして形状を評価

CO回転振動励起と比較• H2OはCOより線幅が狭い• COはbroad+narrow成分を持っており、disk+wind成分に対応している[Banzatti&Pontoppidan2015]• COのnarrow componentに近いプロファイ
ルを持つ天体があり、これらはwind起源だと思われる• もしnarrow componentが円盤とすると、

R>100auが必要となるが、IRS/SpitzerとPACS/HerschelのH2Oの結果は~a few auの円盤で再現されておりinconsistentFig. 7.— Comparison of composite TEXES emission line profiles and CO fundamental emission line profiles obtained with
(Banzatti & Pontoppidan 2015). ‘NC’ refers to the narrow component extracted from the CO profiles (see Banzatti & Pontopp

ケプラー回転円盤でフィット
L(R) ∝ R p (Rin < R < Rmid)∝ R q (Rmid < R < Rou t)• Single peak profile• Rのcutoffなし(Rmidなしで決まる)

Bestfit, min, maxdS• inclinationのせいではない• Double peak profile• Rmid=0.5-30au• Tex>3500Kなので、基本的にはsnowline
ではなく温度勾配を見ているのだろう• ただ、HLTauの例ではRmidはsnowline(13au)に対応しており、しH2O
アバンダンスやガス密度のtruncationがあれば見えるのかもしれない

S
d
d
d

• RWTau,DRTauでは時間変動• 2007~2011~2013• 円盤構造(tidal arm), dust settlingなどによる温度変化 or chemical changeSSpatial segregation of dust grains in transition disks. SPHERE observations of 2MASSJ16083070-3828268 and RXJ1852.3-3700
Villenave, M. et al.

• 遷移段階円盤2つの高分解能近赤外線観測 with SPHERE• ALMAデータを含めてモデル化• mmダストの方がコンパクトな分布 (radial, verticalどちらも)• 文献も含めて近赤外の穴の位置とミリ波ピーク位置を比較• 惑星で作られた構造かもdeviation of 9 mas. The position of the cut is indicated by the dotted line in the left panel.

• IRDIS-SPHERE/VLT, J- and H-bands, DPI••

the vertical thickness of the disk. If the scale height of thelarge grains is too large, after convolution by the beam, thetwo sides of the ring would appear as connected, leading toa flat intensity
profile along
the minor axis. On the other
二つの遷移段階円盤(J1608,J1852)
hand, if the zone of large grains is too thin vertically for agiven disk mass, the fluxes at the peaks of the minor axis
ALMAアーカイブデータも使用
cut would become too large. The appropriate scale heightin our modeling is between 3 and 5 au at 100 au, which
is similar to the value obtained for HL Tau (1 au at 100 au,
Pinte et al. 2016). We note that a smaller scale height, suchas in HL Tau, can not be excluded by our model. We show
A&A proofs: manuscript no. Transition_disks_J1608_J1852_accepted
the non-convolved
model insimulation)
Appendix A. All the model’s
PSFによるconvolution(≠obs.
parameters are presented in Table 2.

• MCFOST[Pinte+2006,2009]によるRTC•

assuming the same spatial distribution for small grains andJ1608の結果
large grains does not provide a good match to the images.
Small grains. The synthetic scattered light image computed

A&A proofs: manuscript no. Transition_disks_J1608_J1852_accepted

To reproduce the observations, the radial extent of small
grains and their scale height need to be larger than those of
the large grains, from 50 to 150 au radially, and about 12 auvertically at 100 au. We note that this layer is already included in the millimeter predictions displayed in Fig. 4,• inc.~74º, <80deg• Large grain

Fig. 6. Spectral energy distribution of J1608 (blue circles)
and our model prediction (red line). The model is corrected by
the Av.

• minor axisを再現するためArticle number, page 6 of 18
Fig. 4. Left: Normalized ALMA Band 6 image of J1608, with a peak flux of 7.1 mJy. The beam is shown in the bottom left of theimage. Middle: Normalized model obtained after convolution by a beam of 0.2400 ⇥0.2300 . Right: Radial cuts along the major andminor axis, with each map normalized to its maximum. For clarity, the cut along the minor axis is shifted by -0.3 in normalized
intensity. The peaks are located at 0.400 (62 au) along the major axis and 0.1400 (22 au) along the minor axis.
A&A proofs: manuscript no. Transition_disks_J1608_J1852_accepted• Small grain

Fig. 4. Left: Normalized ALMA Band 6 image of J1608, with a peak flux of 7.1 mJy. The beam is shown in the bottom left of theimage. Middle: Normalized model obtained after convolution by a beam of 0.2400 ⇥0.2300 . Right: Radial cuts along the major andminor axis, with each map normalized to its maximum. For clarity, the cut along the minor axis is shifted by -0.3 in normalized
intensity. The peaks are located at 0.400 (62 au) along the major axis and 0.1400 (22 au) along the minor axis.
CavityRear facing

Fig. 5. Left: Observed polarized intensity of J1608. Middle: PI model map. Both maps are normalized to their maximum. Right:
Radial cuts along the major axis for the PI images of the data and model, smoothed with a Gaussian kernel with a standard
deviation of 9 mas. The position of the cut is indicated by the dotted line in the left panel.

にhd~3-5au@100au• Rin~77au• R=50-150au• hg=12au@100au• Rin~50au• COガスとcomparable• 北東にbackward scatteringが見えないためには0.5µm
ダストが良いthe polarization degree and the total intensity. While thepolarization degree from backward scattering is similarly Notes. Each parameter was adjusted during the modeling, exlow for all grain size (see Figure 7 of Keppler et al. 2018), cept for the grain size (amin -amax ).
A&A proofs: manuscript
Transition_disks_J1608_J1852_accepted
theno.
phase
function of the intensity varies significantly withthe grain size. We reproduce this feature by selecting grains
smaller than 0.5 µm. Grains larger than 1.2 µm also have avery small eﬃciency for backward scattering. However, withsuch grains our models showed that the position of the maximum intensity along the major axis is independent of theinner radius position and located further out than seen in
the data. These grains are therefore not compatible withthe observations. For grains with an intermediate size (between 0.5 and 1.2 µm), backward scattering is very eﬃcient,Fig. 8. Left: Normalized ALMA Band 3 image of J1852(peak
flux leadis 1.5 tomJy).
The beam issignal
represented
in the bottomof thewhich
would
a significant
northeastof the leftdisk
image. Middle: Normalized model obtained after convolution by a beam of 0.3800 ⇥0.3100 . Right: Cuts along the major and minor
coming
from
its
far
side.
axis, with each map normalized to its maximum. For clarity, the cut along the minor axis is shifted by -0.5 in normalized intensity.

J1852の結果

The peaks are located at 0.300 (44 au) along the major axis.

We show the SED compared to that of the modelin Fig. 6. All the model parameters are presented in TaFig. 8. Left: Normalized ALMA Band 3 image of ble
J1852
is 1.5 mJy).representation
The beam is represented
the bottom
left of the2,(peakand aflux
schematicof our inmodel
is shown
image. Middle: Normalized model obtained after convolution by a beam of 0.3800 ⇥0.3100 . Right: Cuts along the major and minor
in
the
top
panel
of
Fig.
B.1.
axis, with each map normalized to its maximum. For clarity, the cut along the minor axis is shifted by -0.5 in normalized intensity.
The peaks are located at 0.300 (44 au) along the major axis.

We note that we are able to recreate the same general structure as seen in J1608, with two lobes along the
major axis, no emission from the northeast region of the Fig. 7. Spectral energy distribution of J1852 (blue circles) anddisk (backward scattering) and a faint line south of the our model (red line). The model is corrected by the Av.
A&A proofs: manuscript no. Transition_disks
disk, but these characteristics are not perfectly reproduced.
For example the southern line is brighter in our model thanin the observations (see Fig A.1). To reproduce better this 4.4. Modeling J1852
part of the image, the layer of small grains scattering lightmust be thinner vertically, but this in turn has an impact The SED of J1852 shows a steep increase around 20 µm,
Fig. 9. Left: Observed PI image of J1852. Middle:on
PI themodelextent
map. Bothmapslobes.
are normalized
to theirof theMoreover,themaximum.
southernRight:
line Radial
is typical of transition disks, due to the outer edge of the cavcuts along the major and minor axis for the PI images of the data and the model. Each curve is convolved with a Gaussian kernel
ity. A clear silicate feature at 10 µm is present, characteristic
closer
to
the
star
than
in
the
data.
A
way
to
push
this
with a standard deviation of 11 mas.
line further would be for example to increase the mass of of small hot grains close to the star. While scattered light
Tenuous inner disk
the scattering grains, theToscaleheight or the inclination. image shows signal just outside of the coronagraph, there
outer ring, to account for the second brightness increase in Outer ring.
reproduce the peak of intensity
is no corresponding clear millimeter emission at the same00
case0.295
would
lead
toan
anincrease
extinction
the phothe near-IR polarized image and the millimeterHowever,ring. The eacharound(⇠43
au),
of theofsurface
density->silicatefeaturetosphere
that
is
not
observed
in
the
SED.
Thus
the
model
structureis illustrated
bottompanelMiddle:of Fig. PI
B.1.
The map.
in theis needed,
42 and Right:
65 au.Radial
Thelocation. The scattered light image shows a second brightFig. 9. Left:
Observed in
PIthe
image
of J1852.
model
Bothsmallmapsgrains
are normalized
to between
their maximum.
ring at 0.29500 , which does correspond to the ring detected
cuts along
and minor
for the PI
imagesthe data
and themodel.
is convolved
with a Gaussian
kernelthreezonesthearemajor
described
in theaxis
following
sections.
mass
of
this
iscurveadjustedtothebothrelative
thatofwepresent
hereisregionaEachcompromise
to reproduce
reproducewith a standard deviationmas.
Innerof 11ringbrightness
of
the
rings
and
the
SED.
The
scale
heightisthe ALMA and SPHERE images, together with the SED. in the millimeter image.
to be the samethat of theinnertoring.
This should fixed
be considered
as aasworking
model
derive the
To reproduce the scattered light and SED features, we->
内側の散乱光円盤
Inner
The
innerSEDのジャンプ,
region,responsible
for thesilicate
outerdisk.
ring, toaccountforthe secondbrightness
increase
in
Outer
ring.
To
reproducethepeakofTheseintensity
Our
model
reproduces
relatively
well
the position
andconsider three zones in our model: 1/ a tenuous inner disk
main structural characteristics
of
the
J1608
system.
emissionaround10 µm image
in theand
SED,
modeled bysmall
the near-IR
polarized
theismillimeter
ring.
The
around
0.29500 of2/
(⇠43
au),
an in
increase
ofconcentrated
theandsurfacethebrightness
the
peaks
the
major
minordensity
axisregion, solely required to reproduce the silicate feature,are:1/aishighinclination,
large
grainsmore
grains(0.01-5µm)
betweenand 5panelau, although
thisThestructureis illustrated
in the0.1
bottomof Fig. B.1.
in the small
needed,between
42toandau. Thedirections.
Thegrainssecondis ringshows
an oﬀsetthe65
center
as
Outer
ring
verticallyandradially
than
the
small
grains,and3/adis2/ an inner ring, responsible for the sharp jump in the midnot
well-constrained.
To reproduce
the shapeof the emis- seen
threezones are describedin the following
sections.
massinofthethis
regionis adjusted
to notreproduce
relative
data.
However,we wereable to thereproduce
tribution
of
size
in
small
grainsthat
produceslowpolarizedIRsion, we consider a mixture of silicate, composed at 65% thebrightness
of
the
rings
and
the
SED.
The
scale
heightis SED and for the emission seen in the SPHERE image
surfacebrightness
insidethegapbetween
the
inner
and
-> ミリ波of olivine (Dorschner et al. 1995) and 35% of astronomical
intensity in outerbackward
scattering.
justoutside of the coronagraph radius, and finally, 3/ anthesamesouthern
part
disk.
This region isfixed ring
to beintheas thatof ofthetheinnerring.

• 3つの円盤構造で再現•••

円盤構造まとめsilicates(DraineLee 1984).
brighterthe data thanin the model.
our model
Inner disk.
The&innerregion, responsible for the silicate twiceOurmodelinreproduces
relatively
well theInposition
andis empty
ring
are sharp,
while axis
inemission around 10 µm in the SED, is modeled by small thethegapbrightnessof andthe thepeaks
in edges
the major
and minor
practice,alow
surfacedensityofsmall
grainscouldbefillinggrains (0.01-5 µm) between 0.1 and 5 au, although this is directions. The second ring shows an oﬀset to the center as
Innerring. The position
the wall in
SEDoftraces
the it.seen in the data. However, we were not able to reproduce
not well-constrained.
Toofreproduce
thetheshapethe emisdust
is dependent
on both
the innerrasion,temperature,
we consider whicha mixtureof silicate,composed
at 65%
the surface brightness inside the gap between the inner anddius
and thesize of theetdust.
For simplicity,
to useof olivine(Dorschneral. 1995)
and 35% weof choseastronomical

Article number, page 7 of 18

he
15,
mto
illd
in
.

rien
al.
oar
lli
at
esill
nd
ht
al.

o94;
de-

HD142527

22天体で比較• NIRのcavityよりmm-ringの方が遠い場所にピーク

PDS70
J1852

J1608

[de Juan Ovelar+2013]
Fig. 10. Comparison of the cavity sizes as measured in the millimeter deand Juanin scattered
light for a sample of 22 transition disks.
Ovelar+2013のモデルと比較
The cyan points correspond to the two systems modeled in thispaper, while the yellow star symbols refer to HD 142527 that hasa stellar惑星質量を満たす天体は15個存在
companion in the cavity, and PDS 70 where a planet has
been detected. The upper limits, indicated with arrows, refer tosystems for which scattered light cavities could not be measured
down to the coronagraphic radius. Horizontal bars correspond to
one tenth of the millimeter beam, while vertical bars represent
the interval between the cavity radius and the peak in scattered
light (Table 3). The red region shows the ratio for which theplanet mass inferred with the prescription of de Juan Ovelar••

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