66.

Turbulent vortex with moderate dust settling probed by scattering-induced polarization
66. in the IRS 48 system
Haifeng Yang, Manuel Fernández-López, Zhi-Yun Li, Ian W. Stephens, Leslie W. Looney, Zhe-Yu Daniel Lin,
Rachel Harrison F We investigate the crescent-shaped dust trap in the transition disk, Oph IRS 48, using well-resolved (sub)millimeter polarimetric observations at ALMA Band 7 (870 µm). The dust polarization map reveals
patterns consistent with dust scattering-induced polarization. There is a relative displacement between the polarized
flux and the total flux, which holds the key to understanding the dust scale heights in this system. We model
the polarization observations, focusing on the eﬀects of dust scale heights. We find that the interplay between the
inclination-induced polarization and the polarization arising from radiation anisotropy in the crescent determines theobserved polarization; the anisotropy is controlled by the dust optical depth along the midplane, which is, in turn,
determined by the dust scale height in the vertical direction. We find that the dust grains can neither be completely
settled nor well mixed with the gas. The completely settled case produces little radial displacement between the total
and polarized flux, while the well-mixed case produces an azimuthal pattern in the outer (radial) edge of the crescentthat is not observed. Our best model has a gas-to-dust scale height ratio of 2, and can reproduce both the radialdisplacement and the azimuthal displacement between the total and polarized flux. We infer an eﬀective turbulence↵ parameter of approximately 0.0001 0.005. The scattering-induced polarization provides insight into a turbulent
vortex with a moderate level of dust settling in the IRS 48 system, which is hard to achieve otherwise.

IRS 48 の偏光観測によるダスト空間分布の制限と、乱流強度の議論
Gas content and evolution of a sample of YSO associations at d . 3.5 kpc from the Sun
Guang-Xing Li, Bing-Qiu Chen F Young Stellar Objects (YSO) are newly formed stars from・Oph
IRS Zhou,
48：非軸対象な円盤
67. Ji-Xuan
molecular clouds. They stay close to where they were born and serve as good tracers to study gas and star formation.
・偏光観測によりダストの空間、サイズを推定
During cloud evolution, young massive stars can disrupt the surrounding gas through stellar feedback, changing thegas distribution. We study the distribution of the gas around a sample of YSO associations located at d . 3.5 kpc
・乱流強度の推定もfrom the Sun by comparing the location and morphology between 12 CO (J = 1 0) emission, Planck 870 µm maps and

YSO associations. Based on the spatial distribution of the gas compared to that of the YSOs, we classify the YSO
associations into six types: direct, close, bubble, complex, diﬀuse, and clean associations. The complex associations
are large structures consisting of both gas-rich and gas-poor segments. We study the velocity dispersion-size relation
toward diﬀerent association types. From the ratio between diﬀerent types, we estimate a feedback time of ⇡ 1.7 Myrin the solar neighborhood. The sample sets a solid foundation to explore the relationship between interstellar medium
evolution, star formation, and Galaxy structure.

イントロ・円盤ダストのサイズや乱流強度は、惑星形成における議論において重要観測とその結果
・ターゲット：IRS 48：非軸対称天体、非常に明るい・観測設定：ALMA Band
polarization
Turbulent7、fullvortex probedby scattering-induced polarization in IRS 48

3

Figure 1. Polarized dust continuum images. Left: The colormap and the gray contours represent the polarized intensity
in mJy/beam. Gray contour levels are plotted at the levels of (4, 8, 16) ⇥ PI , with PI = 12 µJy/beam, which is the sameacross all three panels. The Stokes I image is shown with white contours plotted at the levels of ( 3, 3, 64, 1024) ⇥ I , with00
00
I = 14 µJy/beam. The cyan line segments represent polarization. The yellow rectangle with 1 ⇥ 0. 48 is the region of the rightpanels. The synthesized beam is 0.00 11 ⇥ 0.00 072 with a position angle of 73 . Upper right: The polarization fraction, with redline segments of uniform length representing the polarization orientation. Only regions with > 3 detection are shown. Lowerright: The position angle of polarization orientation in degree.

trend manifests itself as a displacement between the po-

deprojected as 0.00 019 or 2.5 au in the disk plane. The

解釈のための前提知識：自己散乱による偏光

2.5. Anisotropic Radiation Field

See also Kataoka et al. 2016aof the disc closer to the observer (the right half) brighter. The polarization fraction is, however, higher on the far side (especiallytowards the outer part of the disc) because the polarization degreeof the scattered light is higher for backward scattering than forforward scattering (see Fig. 6). The most striking difference between this case and the Rayleigh scattering case shown in Fig. 5lies in the polarization direction. The difference comes from thepolarization reversal in the large grain case, which yields an intrinsic (or face-on) polarization direction in the radial (as opposed toazimuthal) direction and an inclination-induced polarization along
the major (rather than minor) axis. The interplay between the intrinsic and inclination-induced polarization leads to polarization
directions in the region of high polarized intensity (the most easily
observable part) completely different from those observed in HL
Tau (see Fig. 1).
We are therefore left with an interesting conundrum. The polarization pattern in the HL Tau disc is suggestive of Rayleigh(disk, edge-on view)

H. Yang et al.
Yang, Li, et al. 2016

The second condition for producing the polarization due toscattering is the asymmetry in the distribution of the lightsource. In the case of sub-millimeter emission from protoplanetary disks, the light source is the thermal emission of dust
particles. Therefore, its spatial distribution should be anisotropic to produce polarized scattered light. For example, if theincident radiation from one direction is stronger than that fromthe direction different by 90, the ﬁnal scattering is partiallypolarized. Such asymmetry or anisotropy may occur inprotoplanetary disks showing ring-like or lopsided surface
brightness (e.g., Casassus et al. 2013; Fukagawa et al. 2013;
Isella et al. 2013; van der Marel et al. 2013; Pérez et al. 2014).
Akimasa Kataoka (NAOJ)

i=45°

Figure 7. Scattering-induced polarization by large grains. As in Fig. 5,plotted are the polarized intensity (colour map) and polarization vectors(line segments, with length proportional to the polarization fraction). Notethe strong asymmetry with respect to the major axis in both the polarized
intensity and the polarization vectors. The polarization along the major axis
in the central region is due to polarization reversal, which may be a robust
indicator of scattering by large, mm/cm-sized, grains. The near side of thedisc is on the right.

Figure 3. Polarization P at the scattering angle of 90° and the albedo w = ksca (kabs + ksca ) as a function of maximum grain size. The size distribution is assumed tohave a power law of n (a ) µ a-3.5 . The wavelengths are assumed to be 3.1 mm for the left panel and 870 m m for the right panel. The arrows indicate the maximum
grain size that has the most efﬁcient polarization by 90° scattering.

based mostly on CARMA observations at 1.3 and 2.7 mm. It is inagreement with the spatially averaged value obtained from ALMA
observations from 0.87 to 2.9 mm (ALMA Partnership: Brogan
et al. 2015). This value is significantly lower than the typical ISM
value of β ∼ 1.5–2. The difference is usually taken as evidence
for grain growth to millimetre size or larger (Testi et al. 2014), although other interpretations are possible. For example, Ricci et al.
(2012) showed that a value of β ∼ 1 or lower can be obtained without mm/cm-sized grains if part of the disc is optically thick. Somesupport for this possibility is provided by the spatially resolved distribution of β derived from the ALMA data, which shows β ∼ 0
indicative of optically thick emission at the central continuum peak
and two rings (B1 and B6, ALMA Partnership: Brogan et al. 2015,
see their fig. 3). Another possibility is that the index β is sensitive to not only the size but also the shape of the grains. Indeed,
Verhoeff et al. (2011) were able to reproduce the spectral energy
distribution (SED) of the disc of HD 142527 (with β ∼ 1 in the
millimetre regime) with irregular grains of sizes up to only 2.5 µm;
the grain shape was treated with the distribution of hollow spheres(Min, Hovenier & de Koter 2005). The grains inferred in our modelof dust scattering-induced polarization for the HL Tau disc have asignificantly larger maximum size (of the order of tens of microns).
They may still be able to reproduce the observed (averaged) opacity
spectral index of β ∼ 0.73 if the grains are irregular and/or part ofthe disc is optically thick. Detailed exploration of this possibility isbeyond the scope of the present work.
If large, mm/cm-sized, grains are responsible for the relatively
low value of β observed in the HL Tau disc, it is natural to ask
whether they can produce a polarization pattern that matches theobserved one through scattering. It is unlikely, because the key toproducing the observed pattern is the polarization degree of thescattered light peaking near 90◦ (as in the Rayleigh limit), and thisrequirement is not satisfied for mm/cm-sized grains. For example,
for the grain model adopted by Kataoka et al. (2015a), the polarization degree (defined as the ratio of the two elements in the scatteringmatrix, −Z12 /Z11 , which is essentially the polarization fraction but
can be either positive or negative) is nearly zero at 0.87 mm for all
scattering angles except around 135◦ , where it reaches a (negative)‘peak’ value of ∼−0.2 for amax = 1 mm and 1 cm (see the rightpanel of their fig. 2). The negative value is known as the polarization
reversal (e.g. Murakawa 2010; Kirchschlager & Wolf 2014) which,
together with the shift of the polarization ‘peak’ away from 90◦ , isexpected to produce a polarization pattern very different from the・ミリ波偏光の要因の一つ：自己散乱・ダストサイズが波長と同程度のとき(a=λ/2pi)
・非等方複写場の非等方性に対して垂直に偏光・傾いた円盤では短軸方向に偏光・face on 円盤ではに偏光・スケールハイトに依存

2802

Figure 2. Degree of polarization -Z12 Z11 of dust grains against scattering angles. The left panel is for the case of l = 3.1 mm and the right panel for l = 870 m m .
The size distribution is assumed to have power law of n (a ) µ a-3.5. The maximum grain sizes are 10 μm, 100 μm, 1 mm, and 1 cm.

結果・モデリング・Stokes I と PI でピーク位置がずれている
6・scaleheightをf̲set=
2にすれば再現可能4
Yang et al.
Parameterrc

RA (asec)
Dec (asec)

0.0

c

⌃0
amin
fsettle

0.2
0.4
0.2

0.0

0.3
0.4
0.5
0.6
Deprojected distance (asec)

0.2

0.4

Yang et al.

Fiducial value
61 au
109
1.0 g/cm2
0.1 µm2

Parameter
rw
w

T60
amax

Fiducial value
10 au
30
60 K
140 µmTable 1. Parameters and their fiducial value in our model.
See the text for their definitions.

0.6

et al. (2015) and Ohashi et al. (2020):✓◆
4
4(r rc )(c)⌃d (r, ) = ⌃0 exp
,
4
4
2rw
2 w

60
90
120
150
Deprojected angle (deg)Figure 2. Top: The total flux (black contours) and thepolarized flux (red contours). The peaks for the total flux
and the polarized flux are shown as red and black plus signs,
respectively. The grey fan-like region is enclosed by constant radius curves (r = 0.00 25, 0.00 6) and constantlines( = 60 , 150 ) in the deprojected disk plane. Bottom Left:
Radial profiles (azimuthally averaged) of total flux (black)and polarized flux (red). Their peaks are labeled with vertical straight lines. Bottom Right: Azimuthal profiles (radiallyaveraged) of total flux (black) and polarized flux (red). Theirpeaks are labeled with vertical straight lines. Note that the
total flux and polarized flux have di↵erent vertical scales.

(1)

where rc = 61 au and c = 109 are radial and azimuthal
locations of the peak, rw = 10 au and w = 30 are
the radial and azimuthal extent, and ⌃0 is the column
density at the peak. Note that = 0 corresponds to themajor axis of the projected disk, which has a position
angle of 100 in the sky plane.
For the temperature structure of the disk, we adopt a
simple power-law profile:⇣ r ⌘ 0.5
Td = T60
,
60 K

(2)

where we take T60 = 60 K as the fiducial number, similar
one comparisonsadopted bybetween
Ohashitheetdataal. (2020),
which wasFigure to
3. theDetailed
and our models.
From top to bottom, each row represents the data andetc. It is impossible to discuss the true nature behindfsettle = 2, but we can get a quantitative sense of themixing, in terms of the turbulent parameter ↵ (Shakura& Sunyaev 1973). In the most common scenario, the・乱流強度の制限
gas-to-dust scaleheightetratio
Yang
al.follows (Youdin & Lithwick 2007):self-scattering. Grains aligned with toroidal magnetic
fields would produce polarization vectors that all point
towards the central star. If the magnetic fields possess
a complicated eddy structure or if the grains are alignedwith eddy di↵erential motions between gas and dust,
we would expect much more complicated polarization
・ダストスケールハイトからα/Stを制限
patterns.
Radiativealignment,on the other hand, de✓◆ especially◆ 1/2
tions that makeantransition
disks
with
crescent
structures,
the
1/2 ✓
Hd
St
1 + 2St
pends onthe
anisotropy
in radiation flux at the peak
= 1+, than(3) 10 (Ohashi
ing in a significant
G/D
can
be
smalleretal.
2020),oreven
Hg↵
1 + Stof the spectrum energy distribution inside the crescentponsible for the azas low as 3 in the case of HD 142527
(Soon
et
al.
2019).
shaped structure. That is to say, the radiative alignwhere
St
is
the
Stokesnumber.
For
f
=
2,
we
have
settle um ダスト
・B7偏光が受かってる
=
100
=
2
model.
Fora
G/D
between 3 and 100,
we depends
have theturbulence
lement
mainlyon the light inside the crescent↵ ⇡ St/3. In the Stokes regime, we have St = ⇢s a/⌃g ,fsettle=⇢s2is model
parameterasgrain↵ =size,
0.0001
at a wavelength of about 10s of microns. Given that・ガス面密度等を仮定すると、
where
the dust solid density,
a is theand ⇠ 0.005.
smallergrains aremoreazimuthally extended (Geers
r observation.
estimate
spherical
dust
grains.
⌃g is ・不定性は主にガス面密度から
the gas This
column density.This
For theadopted assumes
grain size compact
et al.does
2007;notvana↵ectder Marelal. 2013), we would expectof 140µm, weshift
have:onableradialIt should be noted that porosity
theetina mostly radial radiation flux in the crescent, resulting✓◆ferred✓◆ 1sizes✓◆ 1much. In general, the peak waveux and total flux.
grain
too
in an azimuthal polarization pattern. This is completely
⇢s
⌃d
G/D
↵ = 0.0014⇥, opacity remains the same as longfor theazimuthallengthofthescattering
32
in contrast to our observation.
3 g/cm
1 g/cm10, where the east
as f ⇥ amax is fixed, wherefillingfactor(4) f is theWe cansee thatthe(Tazaki
theoretical expectations on grain
2
where
⌃
=
1
g/cm
adopteddust column
theircorresponding polarization patterns
mpared to dthe west is the etal. 2019).
So, thedensity,
maximum alignment
grain sizeand
thatis responsiare
very diverseroughly
and interesting.
Our observation at
G/D
is the gas-to-dust
ratio.
Thescattering-induced
greatest uncerromandtheinterplayble
for thepolarization
goes
tainty comes from the gas column density. The lower
870 µm shows no signs of the aforementioned patterns
tion and radiation
as 1/f . At the same time, the dust solid mass density
density of the gas column requires a higher turbulence tofrom grain alignment. Observations towards even longer
is withreducede↵ects cancel
each others . These twowavelengths
stir the dust grains to a state
fsettle to= 2.f ⇢
Nominally,
are needed
to see whether and how grains
out,Inwhichmeans thatthe inferredturbulent
the gas-to-dust ratio is 100.
protoplanetarydisks,
are alignedin thisparametertransition disk system. Even a spa-

議論

f turbulence↵ remains the same.

58. Yangfan Shi, Feng Long, Gregory J. Herczeg, Daniel Harsono, Yao Liu, Paola Pinilla, Enrico Ragusa, Doug
58. Small and Large Dust Cavities in Disks around mid-M Stars in Taurus

Johnstone, Xue-Ning Bai, Ilaria Pascucci, Carlo F. Manara, Gijs D. Mulders, Lucas A. Cieza F High-angularresolution imaging by ALMA has revealed the near-universality and diversity of substructures in protoplanetary disks.
However, disks around M-type pre-main-sequence stars are still poorly sampled, despite the prevalence of M-dwarfs
in the galaxy. Here we present high-resolution ( 50 mas, 8 au) ALMA Band 6 observations of six disks around mid-M
stars in Taurus. We detect dust continuum emission in all six disks, 12CO in five disks, and 13CO line in two disks.
The size ratios between gas and dust disks range from 1.6 to 5.1. The ratio of about 5 for 2M0436 and 2M0450
indicates eﬃcient dust radial drift. Four disks show rings and cavities and two disks are smooth. The cavity sizes
occupy a wide range: 60 au for 2M0412, and 10 au for 2M0434, 2M0436 and 2M0508. Detailed visibility modelingindicates that small cavities of 1.7 and 5.7 au may hide in the two smooth disks 2M0450 and CIDA 12. We perform
radiative transfer fitting of the infrared SEDs to constrain the cavity sizes, finding that micron-sized dust grains may
have smaller cavities than millimeter grains. Planet-disk interactions are the preferred explanation to produce thelarge 60 au cavity, while other physics could be responsible for the three 10 au cavities under current observations
and theories. Currently, disks around mid-to-late M stars in Taurus show a higher detection frequency of cavities thanearlier type stars, although a more complete sample is needed to evaluate any dependence of substructure on stellarmass.

M型星周りの円盤の高解像度
ALMA
観測 (JOYS): Linked accretion and ejection in a
JWST Observations of Young
protoStars
59. Class I protobinary systemŁukasz Tychoniec, Martijn L. van Gelder, Ewine F. van Dishoeck, Logan Francis, Will R. M. Rocha, Alessio
Caratti o Garatti, Henrik Beuther, Caroline Gieser, Kay Justtanont, Harold Linnartz, Valentin J. M. Le
Gouellec, Giulia Perotti, Dev Rangaswamy, Benoît Tabone, Thomas P. Ray, Nashanty G. C. Brunken, Yuan
Chen, Patrick J. Kavanagh, Pamela Klaassen, Katerina Slavicinska, Manuel Güdel, Goran Östlin F Accretionand ejection sets the outcome of the star and planet formation process. The mid-infrared wavelength range oﬀers
key tracers of those processes that were diﬃcult to detect and spatially resolve in protostars until now. We aim tocharacterize the interplay between accretion and ejection in the low-mass Class I protobinary system TMC1, comprising
two young stellar objects: TMC1-W and TMC1-E with 85 au separation. With the James Webb Space Telescope

21

イントロ
・M型星(軽い星)での惑星形成の重要性
・M型星周りでも惑星は見つかっている
・M型星周りでは惑星形成は難しい(小さい円盤、早いダストの消失)
・M型星円盤の高解像度観測の必要性
・円盤が小さいので構造を見るには高解像度観測が必要
・暗いのであまり調べられていない

4

urce

412
434
436
450
A 12
508

観測
・ALMA B6
・連続波
・CO、13CO、C18O

S HI ET AL .

Table 1. Properties of Host Stars
2MASS | Short Name

D

SpT

[pc]
J04124068+2438157 | 2M0412
J04343128+1722201 | 2M0434
J04360131+1726120 | 2M0436
J04504003+1619460 | 2M0450
J05075496+2500156 | CIDA 12
J05080709+2427123 | 2M0508

148.7
145.7
144.5
144.4
170.0
170.7

AV

J

[mag]
M4.3
M4.3
M2.7
M4.75
M3.7
M3.5

0.84
0.30
1.08
0.0
0.5
0.0

11.151
11.20511.10511.725
11.415
11.396

Lphot

M⇤

R⇤

Lacc

Macc

Refs

[L ]

[M ]

[R ]

[L ]

[M yr 1 ]

0.126
0.115
0.146
0.054
0.13
0.11

0.30
0.30
0.58
0.24
0.39
0.44

1.21
1.16
1.07
0.84
1.13
1.03

0.00204
0.00168
0.0062–0.0028–3.2e-102.6e-104.5e-10–3.2e-10
–

L23
–
–
L18
HH14
L18

References—L23: Long et al. (2023); L18: Luhman (2018); HH14: Herczeg & Hillenbrand (2014).

結果
Obs. Date

Table 2. Summary of ALMA observations

Nant

Baselines

On-source Time

Mean PWV

Peak I⌫ (taper)

RMS Noise (taper)Beam (taper)

[m]

[min]

[mm]

[mJy/beam][mJy/beam][mas⇥mas, deg]

2021-08-27
38
92.1-10803.3
16.80
0.3
0.19 (0.26)
0.032 (0.039)
74⇥35, 37 (76⇥73, 32)
・大きなcavityが見えた天体から何もない天体まで存在
2021-09-27
45
70.1-14361.8
15.29
1.0
0.47 (0.69)
0.039 (0.041)
45⇥24, -47 (43⇥41, -41)

・60 au ギャップは惑星によるものだろう

2021-09-27
2021-09-28
2021-09-13
2021-09-13

45
47
37
37

70.1-14361.8
70.1-14361.8
178.3-12594.5
178.3-12594.5

15.39
15.05
16.80
16.87

1.0
0.7
0.6
0.6

0.31 (0.42)
0.23 (0.27)
0.19 (0.30)
0.38 (0.63)

・10 au のものはいろんな説がある

0.037 (0.039)
0.029 (0.029)
0.031 (0.036)
0.035 (0.039)

44⇥24, -46 (43⇥42, 38)
Figure 2. Best-fit visibility fitting results versus observations. Columns from left to right: (1) the real part of deprojected and binned visibilities
36⇥26, -40 (36⇥36, -35)
from observations and best-fit models, (2) observational continuum images, (3) model images convolved with the same beam as observations,72⇥33, 4 (76⇥73, -9)(4) residual images with white contours at 3 and green contours at 3 , overlaid orange contours are from continuum images at 3 level (5)azimuthally averaged radial profiles for observations and convolved model images, light gray shaded regions represent the standard deviation
65⇥25, 4 (63⇥60, 36)at each radial bin divided by the square root of the number of beams in the radial bin.

結果・ガスとダストの広がりの比は、1.6から5.1まで
・5.1は、ダストのradial driftが起きていることを示唆・ダストトラップされて構造が見つかっているものでは比が小さい
10

S HI ET AL .

Table 4. Measurement of gas emission from CO isotopes
Molecule

Integrated Flux

Peak Intensity

Velocity Range

Chan. Width

Chan. RMS

Beam

R90%

[Jy km s 1 ][mJy beam 1 km s 1 ][km s 1 ][km s 1 ][mJy beam 1 ][mas⇥mas, deg]

[au]

5.5
5.8
4.2

215⇥209, -3
220⇥212, -2
222⇥211, 5

198.9
187.2
-

1.6
1.5
-

6.9
4.4
3.4

209⇥175, -67
170⇥159, -57
169⇥159, -55

153.9
155.9
-

1.9
1.9
-

3.5
3.8
2.9

112⇥99, -51
113⇥100, -55
113⇥101, -54

64.4
-

4.9
-

2.6
2.9
2.2

110⇥97, -77
112⇥98, -82
111⇥99, -85

25.8
-

5.1
-

3.0
3.3
2.5

101⇥94, -37
102⇥96, -40
101⇥97, -35

53.7
-

2.7
-

2.9
3.3
2.5

100⇥93, -39
101⇥96, -43
101⇥97, -38

56.3
-

2.8
-

12

CO
13
CO
C18 O

1.08 ± 0.06
0.54 ± 0.06
< 0.15

71.4
39.2
-

4.0-8.0
4.0-8.0
4.0-8.0

12

CO
13
CO
C18 O

1.21 ± 0.11
0.61 ± 0.09
< 0.21

130.5
74.2
-

2.0-10.0
2.0-10.0
2.0-10.0

12

CO
13
CO
C18 O

0.38 ± 0.07
< 0.24< 0.1756.1-

1.4-9.4
1.4-9.4
1.4-9.4

12

CO
13
CO
C18 O

0.10 ± 0.04
< 0.11< 0.1030.3-

2.0-12.0
2.0-12.0
2.0-12.0

12

CO
13
CO
C18 O

0.24 ± 0.05
< 0.15< 0.1257.6-

2.0-12.0
2.0-12.0
2.0-12.012

0.19 ± 0.06
< 0.20< 0.1939.6-

2.0-12.0
2.0-12.0
2.0-12.0

CO
13
CO
C18 O

2M04120.4
0.4
0.4
2M04340.4
1.0
1.0
2M04360.8
0.8
0.8
2M04501.0
1.0
1.0
CIDA 121.0
1.0
1.0
2M05081.0
1.0
1.0

Rgas /Rdust