SFN #330 (36-42本目)
百瀬宗武(茨城大・理)詳細解説は以下の3本• #37: Validating Scattering-Induced (Sub)millimeter Disk Polarization

through the Spectral Index, Wavelength-Dependent Polarization
Pattern, and Polarization Spectrum: The Case of HD 163296• #43: The big sibling of AU Mic: a cold dust-rich debris disk aroundCP-72 2713 in the β Pic moving group• #40: X-Shooter survey of disk accretion in Upper Scorpius I. Very highaccretion rates at age>5 Myr

from the Fortran 77 version by Bruce Draine 3 that is based
Dust
Self-ScatteringinHD
1632963the original code Disk
publishedby Bohren & Hu↵man (1983).
eﬀ vs. amax】
37. Validating Scattering-Inducedon(Sub)millimeter
Polarization【偏光効率P,
有効アルベドωA useful quantity to assess the dust scattering-produced
2.2 Disk Modelthrough the Spectral Index, Wavelength-Dependent
Polarization
2.1 Dust Opacity
polarization at a given wavelength, , is the product of thePattern, and Polarization Spectrum:
The
Case
of
HD
163296
We consider a dust density distribution that is determinedWe
follow
the
DSHARP
dust
model
in Birnstiel
single scattering albedo, !⌫ , and prescribed
single scatteringpolarizaZYD Lin et al. MNRAS, 496, 169-181;et https://arxiv.org/pdf/1912.10012.pdf
by distributions of surface density and scale height. We folal.fraction(2018). Opacity
calculated
Mie Theory whichtiontaken atis90, P = zusing
12 /z11 , where z12 and z11 areassumes compact(non-porous)
spherical
grains andlow the DSHARP notation in Huang et al. (2018) to describe
elementsof the scatteringmatrix(Kataokaet requiresal. 2015). !⌫【内容】ALMAによるHD163296ダスト連続波熱放射をRADMC-3Dを用
as input parameters the grain size and refractive indices. The
the substructures of HD 163296. For example, “B100” repreincreases
rapidly as the grain size parameter x ⌘ 2⇡a/ apdust model is a mixture of water ice, astronomical silicates,
いてモデリング。Band 7偏光観測イメージやBand6-7でのスペクトル指
sents the bright annular feature with a radius of 100 au. To
proaches unity, but P drops quickly for large (x > 1) grains.
troilite, and refractory organic material with mass fractions
model HD 163296-like disks with prominent substructures
These
two
competinge↵ects
mean
that
the
polarization
ata
数と一致するのは最大サイズamax =90μmのダストによる散乱がある場合
of ⇠ 0.2, 0.33, 0.07, and 0.4, respectively. The respective re(rings andThegaps), Gaussian-likewe adopt a compositeprofile
surface density
disisprescribgiven
wavelengthisdeterminedbygrainsofacertainrange
fractive indices of the di↵erent components are taken from
なことを確認。将来の多波長観測についても予言。方位角方向での偏光
tribution ⌃d that is composed of two di↵erent types of proof
sizes&(Kataoka
et al.Draine2015).(2003),
For aand,particle
WarrenBrandt (2008),
for thewith
lattera size◆✓files:
power-law-likeprofilesandGaussian-likeprofiles.
The
2largerthanthewavelength,
forwardscattering
dominates
two, Henning & Stognienko (1996). Furthermore, we average
角の変化がリングとギャップで異なる傾向。βとのtension（より大きな
8
1
R
R
>
formerconsistsofapower-lawdistribution
withinaregion
0
over
a
grain
size
distribution
thatfollowsa
simplepower-law,!⌫ even though forward scattering is e↵ectively not scat>exp,>
ダストサイズが必要）は，詳細構造を十分空間分解能しておらず，光学
3.5 , with minimum and maximum grain size cut-o↵s (thebetween
radii
of
R
and
R
and
an
exponential
decrease
in<
>
a
a
b
tering. A simple replacement for the apparent scattering ef2
ce↵
standard
MRN
grain
size
distribution;
Mathisetal.
1977),
teriorandexterior
totheregion:
的に厚いリングと薄いギャップを平均化したためではないか。
⌃G (R) = ⌃G,0 ⇥ficiency is the e↵ective albedo, !⌫ = (1 g⌫ )!⌫ where g⌫
◆2

✓
where
is the grain size. The
minimum (Ishimaru
grain size is1978;
fixed Birn>【柱密度モデル】
1
R
R
isthe aforward-scatteringparameter>0>
at 0.01µm, while the maximum grain size amax is left as,
exp>stiel et al. 2018; Tazaki et al. 2019; Zhu et al. 2019). Fig.Figure 1. The calculated degree of linear
polarization for single
a free parameter. The precise value for the minimum graine↵ective◆discussed✓ :22
scattering(P;
reddottedline),
albedoin Sec- d
8
1size
shows
the
calculated
P
and
!
for
ALMA
Band
7
(
=
【冪乗】
>
e↵
1
R
R
a
makes little di↵erence at the considered wavelengths>e↵ ; blue dashede↵ (black【用いたモデル】>line),
andtheproduct
P!
tion
2.1
(!
exp
,
R
< Ra
>
0.87mm)asafunctionofthemaximum
grain
size.
The
prod>
(Kataoka et al. 2015; Birnstiel et al. 2018). Mie opacity calasolid line) at = 0.87mm >for various2 grain sizes.
>e↵<0 Ris the radius of the center of
>
where
R
P!⌫is implemented, gives the grainthat contributes
most• ダストについてはBirnstiel et uct,al. (2018)のDSHARPで使われたもの
culationusingsize
the python
version of the
Mieto thep , R  R< R
(
)
⌃
(R)
=
⌃
⇥
,
a
p
b
p,0
polarization
attheobservingwavelength.
Fig.
1
shows
that
code available in RADMC-3D. The Mie code was ported>
Raetdensity
>
4
Z.
D.
Lin
al.
◆ 2, and✓ at R0
thesurfaceand>（右上）。本論文はamax =90μm固定(λ=0.87mmで偏光効率最大)。
c
3
d
>
the
maximumgrain
size forproducing
fromoptimal
the Fortran77 version
by Bruce
Drainethat ispolarization
based
>
R
1
R
R
b p
b
>
>
, R Rb()
exp> R
on thecode
published
Bohren& Hu↵manat=original
0.87mmis 90µm.
Webywilltherefore
hold a(1983).
• 円盤柱密度分布モデルは冪乗的分布とガウシアン的分布の重ね合わ
the Gaussian
distributions
interior and ex
max fixed at2: ab
Gaussian-like profile is prescribed as
usefulto assessdust scattering-produced
this Avalue
forquantity
our models
andtheassumethe same dust popula-2.2 DiskThe
Model
(1)
せで表現（右下）。
【ガウシアン】
tively.
This
type
of
profileisparticularlyspolarization
at
a
given
wavelength,,istheproduct
of
the

◆2
✓
tion throughout the whole disk; note that this value is some8
We consider a dust density
distribution
thatis determined
1
R
R>
singlescattering
albedo,!
,andsinglescatteringpolariza0⌫>• 冪乗的分布（Σp) はinner diskやouterregionsに好適
what smaller
than thefiducial grain size of 140µm adoptedby distributions
exp,We
R fol R0profile is ob>rings.
The
full
surfacedensityofsurfacedensityandscaleheight.
<>tion fraction taken at 90 , P = z12 /z11 , where z12 and z11 are
R is the cylindrical2 radius,⌃ p,0 is the surface densitycby Ohashi & Kataoka (2019) for their independent mod-low the⌃where(R)=inner⌃G,0⇥ inR Huang,(2)◆ 2and of the✓et al.
DSHARP
notation(2018)
to describeGelements of the scattering matrix (Kataoka et al. 2015). !⌫• ガウシアン（ΣG) はringを表すのに好適
at
the
edge
range
p
is
the
power-law
a
>
all
the
various
1 power-law-like
R R0 “B100” repre- and Gauss
>
eling
of
the
HD
163296disk
polarization.
The
correspondthe
substructuresofHD
163296.
For
example,>increases rapidly as the grain size parameter x ⌘ 2⇡a/ ap, R > toR0 the char> exp
index.
The
quantities
andcorrespondab• 温度分布は冪乗則を，鉛直方向ガス密度分布は静水圧平衡を仮定，
ing
absorption
opacity,albedo,forward-sents the bright annular :feature with
2 a radiusof 100 au. Todproaches
unity, but
P dropse↵ective
quickly forlarge (xand> 1)thegrains.
distanceoverwhich
the surface
density decreases
2 g 1 per gram ofmodel acteristic
HD163296-like
disks
with
prominent
substructures
scatteringparameter
are

=
0.97
cm
abs that the polarization at a
These two competing e↵ects mean
ダストスケールハイトは独立に指定
’
’
is
the
radiusofthecenterof thedis-profile,
⌃G,0
where
R
interior
toR
andexterior
toR, respectively.
This type
of is
e↵
a
0
b
(ringsandgaps),
weadopt
acomposite
surfacedensity
dust,
!
=
0.64,
and
g
=
0.15
at
=
0.87mm,whereasatgiven wavelengthis determinedby grains of a certain range⌫⌫
profileisparticularly
suitableformodeling
the
inner
disk
the
surfacedensity
at
R
,
and
and
are
the
widthsof21e↵ctribution⌃thatiscomposedoftwodi↵erent
types
of
pro⌃
⌃
⌃
=
+
,
0
d
d
of
sizes
(Kataokaetal.
2015).
Foraparticle
with
a
size
p
=1.25mm,

=
0.46
cm
g
,
!
=
0.47,
and
g
=
0.044.
Gd⌫
abs⌫
MCMCでデータと合うパラメータを探すのは大変。
i
i
(R
<
50AU)
and
the
outer
B100
regionofHD
163296.
files: power-law-likeand Gaussian-likeprofiles.
The
the Gaussianprofiles
distributions
interior andexterior
to R0 Since
respeclargeropacity
than thewavelength,
forwardscattering
dominates
The
index
index
betweenthetwowavelengths
dust
ibe ofi material,
formerconsistsofapower-lawdistribution
within awe
regiongaps
may
not
entirely
free
of
adopt
surfacetively.
This
type
profileisparticularly
suitedfora modeling!
eventhoughforwardscatteringise↵ectively
notscat1D-Slab(i=46.7°)を使った簡単なケースであたりをつける(次ページ)。
is⌫ 2.06.
betweenradii offloor
Ra andRb andan exponential
decreaseprofile
indensityusing
the same
power-law-likebut withtering. A simple replacement for the apparent scattering ef-

4

Z. D. Lin et al.

slab(1D, Plane-parallel)を使ったテスト計算（柱密度・温度平面で）Dust Self-Scattering in HD 163296

D5

破線が100au付近のリングB100で取得された観測値（Band 6/7）。
(c)の点線は，Sp index(α)= 2を示す。
α =4 : 光学的に薄い極限

2. The plane-parallel calculation of an isothermal scattering slab in the two
αFigure=2：放射のみで光学的に厚いレイリー
temperature. The panels from the left to right are Band 7 Stokes I, Band 6 Stokes

The dashed lines bracket the observed range for ring B100 of HD 163296. The dott

ジーンズ領域α < 2 : レイリージーンズからのずれ＋Figure 2. The plane-parallel calculation of an isothermal scattering slab in the two-dimensional parameter space of surface density andtemperature. The panels from the left to right are Band 7 Stokes I , Band 6 Stokes I, and spectral index between the two ALMA bands.
The dashed lines bracket the observed range for ring B100 of HD 163296. The dotted line in panel c marks where ↵ = 2.

散乱 (効いている波長では暗くなる効果)
r=20au, 67au, 100auの観測値(Band 6/7)。
光学的に厚い<20au以外，あたりがついた

Figure 3. The range of surface density and temperature constrained by observat
regions correspond to ranges constrained by the Band 7 Stokes I , the Band 6 Stokes
6.

Radius [au]
20
67
100

Figure 3. The range of surface density and temperature constrained by observations at 20, 67 and 100 au. The red, blue, and green
regions correspond to ranges constrained by the Band 7 Stokes I , the Band 6 Stokes I, and the spectral index between Band 7 and Band
6.

Radius [au]Band 7 [mJy beam 1 ]Band 6 [mJy beam 1 ]ues for the FWHM are the geometric mean of the major andBand 6 [mJy beam 1 ]
2.08
0.76
0.44

Table 1. Values from observations used for 1D slab estimationin Section 2.3. The first column is the radius in au. The second
to fourth columns are the observed values for the fluxes in Bands
7 and 6 and the spectral index ↵ between the two bands.

3↵

Band 7 [mJy beam 1 ]
78.1
26.2
12.0

3.1

NUMERICAL RESULTS
Polarization Image↵ ues for the
1.4 minor axes
2.2 ages are on
1.9

mainly inte
(Dent et al
One caqualitativel
the rings pr
thermore, tthose of th
in the inne
are mainly

モデル計算結果
円盤モデル・パラメーター
6

Dust Self-Scattering in HD 163296

Z. D. Lin et al.
Parameters
Rt [au]Tt [K]q
Ht [au]

100
12.5
0.65
3
Power-law-like Surface Density

Feature⌃ p,0 [g cm 2 ]p
R a [au]R b [au]a [au]b [au]

Inner Disk

D48

D86

B100

8
0
0.5
26
0
6

0.02
0
38
58
5
4.5

0.1
0
76
96
5
5

0.5
7
105
135
5
10

Gaussian-like Surface Density
Feature

B67

B100

B155

⌃G,0 [g cm 2 ]R0 [au]c [au]d [au]

2.5
67
5
5

2.1
100
3
4

0.15
155
5
25

Table 2. A chosen set of parameters for the model. The firstcolumn displays the name of the parameter from Section 2 andits corresponding units in squared brackets. The parameters are
separated into three main groups: the general parameters, parameters used for power-law-like surface density components, and parameters used for Gaussian-like surface density components. For
the inner disk, the inner exponential decrease is ignored since it
is not resolved and thus a = 0.

※Dustheightはどう決めたか，
disk(Yangscale
et al. 2016,2017). The distribution of the deviation reflects the radiation anisotropy previously discussed
あまりよく分からなかった
in Kataoka
et al. (2015); Yang et al. (2016, 2017). For example, in an optically thin gap, most of the radiation comes7(⇠ 1.5) because of preferential scattering at shorter wavelengths, although deviation of the black body radiation fromthe Rayleigh-Jeans limit also plays a role, especially for theouter, colder ring. The observed rings have ↵ ⇠ 2, which is
not as low as that in the inner disk since the emission fromthe rings is blended with emission from the gaps that havemuch higher spectral indices.
The polarization fraction after convolution coincides
with the observed polarization within the first gap, but issomewhat lower than the observed values at outer radii given
by Dent et al. (2019). However, our model matches better thepolarization fraction presented in Ohashi & Kataoka (2019)at outer radii (see, e.g., the top panel of their Fig. 18). We
should caution the reader that the polarization fraction may
be significantly a↵ected by the di↵erential filtering of Stokes
I vs. Q and U by interferometers such as ALMA. This e↵ect
has yet to be thoroughly quantified, and will be exploredelsewhere.

4

DISCUSSION

4.1

Consistency Check and Discrepancies

4.1.1

Optical Depth

Figure4. Theimagesof 100µm
the disk model. The near-side of the disk is the bottom half of the image. Panel a: the BandWe find thata grainsizetwoof dimensional
90µm, similar
to the
1 . Panel b: the spectral index between Band 7 and Band 6. Panel c: the Band 7 polarization fraction in
6 Stokes
image reproduces
in Jy beamtheirfound in Dentet al.I(2019),
observed level of
1 , with the polarization vectors overplotted. Panel e: the color scale depicts
percent.
Panel
d:
the
Band
7
polarized
intensity
in
Jy
beam
polarization in the inner disk, although it is lower beyond theo↵set ofpolarizationfromfirst gap. theHowever,thethespectralindex orientation,
they obtained✓,was⇠ the vertical direction (parallel to the disk minor axis). ✓ is plotted east from north(counter-clockwise)= 0 meansthelikelypolarization
orientation is parallel to the disk minor axis. The vectors in Panel e are oriented
3.5 even in
the inner disk and andrings,✓which
is mostdue
the depth
verticaladopted
direction
3 ⇥model✓ to comparedmake the o↵sets obvious. The black ellipses for all six panels denote the rings. The Gaussian
to a lowerfrom
optical
in by
their
used
for convolution
areRplottedas black
filled circles at the lower left.
with ours.beams
For theinnerdisk at radius. 30 au,
we findthat optical depths of ⇠ 10 at Band 6 and higher for Band
7 are needed to produce the low spectral index observed
the midplane, it would not probe the midplane temperature.
indeeddepth
optically
its temperature
there. Thering
highisopticalof thethick,inner disk
implies that would natuIt would probe, instead, the temperature of the so-called
rally
to seethe therelatively
low brightness
one would
not bebe closeable toline emission
from the temperature“CO snow surface” above (or below) the midplane where theobserveinThisFig.seems
5b astoshadedback side thatof thewedisk(Zhu et(shownal. 2019).
be the line), which

観測との比較
長軸方向のプロファイル

8

Z. D. Lin et al.

β~1(α≈3)というのは，α≈4と≈2の平均では？
↓

上3枚のパネル
モデル(赤実線)は観測(グレーやカラー)を概ね再現‥(Ohashi & Kataoka
2019の解析と比較すると，
Polarization fraction はより良く一致)
下2枚のパネルリング(中心面)での温度は，COで求められている温度より低く光学的に厚い可能性がある（外側でのSp indexやPolarization
fractionが高いことをよりよく説明）

Figure 5. Panel a: the 0.00 043 resolution Band 6 Stokes I image in mJy per beam. The units of the non-convolved image is scaled to thecorresponding mJy per beam. The convolved profile is plotted with a red solid line, the non-convolved profile with a black dashed line,
and the grey region represents the ±10% calibration uncertainty of the data. Panel b: the spectral index between Band 7 and Band 6.
The grey region is ↵ ± 0.3, where ↵ is the spectral index averaged between the Northwest and Southeast major axes. The ↵ = 2 is plotted
as a horizontal solid line. Panel c: we plot the convolved Band 7 polarization with a red solid line and the non-convolved polarization
with a black dashed line. The green and blue colored regions are the observed Band 7 polarization with uncertainty from random noise
and image fidelity along the Northwest and Southeast major axes of the disk (Dent et al. 2019). Panel d: the model optical depth along
the major axis for Bands 7 and 6. The lower limits for Band 6 optical depths at 67 and 100 au (denoted by triangles) are derived fromthe level of 12 CO extinction through the dust rings (Isella et al. 2018). The optical depth of 1 is plotted as a horizontal solid line. Panele: a comparison of various dust temperature profiles as a function of radius from previous literature. The vertical dotted lines throughall the panels are the locations of the rings.

strong evidence that the disk midplane is colder than thewhere Hd is the dust scale height and Hg is the gas pressure

多波長でのモデル計算の結果(B7以外の予言)(1) 強い波長依存性
10

•
Pol.
degreeの目盛りの数値に注目
Z. D. Lin et al.
• 長い波長ではτ~1のr<20au以外はPIでは暗い(2) 空間的なパターン
散乱偏光のパターンは，2つの効果の重ね合わせ
1. Inclination による：短軸‖Pol. Direction

2. 非等方な輻射場：• gapでは動径方向に光源(ring)→azimuthal Pol.
• ringでは横方向に光源(ring)→radial Pol.
Figure 6. Images of the polarization fraction without finite telescope resolution for 5 ALMA bands at = 0.87, 1.25, 1.5, 2.1, and 3.1
mm (panels a-e) where the polarization capability is available. The dashed white ellipses mark the location of the rings. Panel f plots
the distribution of the optical depth as a function of radius along the major axis for each band.

1. と2. が平行になり強め合うのは：• Gap では円盤の長軸近辺 (a-cが顕著)• Ring(白い破線)では円盤の短軸近辺，ただし，波長が短いところはoptically thickなので偏光効率は大幅にダウン(d, eで顕著)

maximum value for the color scale is di↵erent for each band,
starting at ⇠ 6% for the shortest-wavelength Band 7 and decreasing to ⇠ 0.35% for the longest-wavelength Band 3. Thechange of polarization fraction with wavelength is discussed
in detail below. Here, we focus on the spatial distribution of
the polarization fraction, which has an interesting transition
from Band 7 to Band 3. For all images, the azimuthal variation along the rings versus that along the gaps are opposite.
Along the rings, the polarization fraction is higher along theminor axis (x = 0). Along the gaps, the polarization fraction
is higher along the major axis (y = 0). The former trend ismost clearly seen in the Band 3 image, while the latter trend
is most obvious at Band 7.
The azimuthal variation can be understood in termsof two competing e↵ects: the inclination-induced polarization and the radiation-anisotropy-induced polarization. For
both the rings and gaps, the inclination-induced polarization produces polarization parallel to the disk minor axis,• 各場所でのPol. Fractionのスペクトル

Dust Self-Scatterin

multi-wavelength polarization dmation to date (Kataoka et al.
et al. 2019).
creases, the inner disk becomes more prominent in the Carrasco-González
pohas rings and gaps and there exi
larization fraction map: compared with the rings and gaps
between the Stokes I and spe
at large radii, its polarization fraction is lower at Bandvarying
7, optical depth as a functcase of HD 163296. The polari
but becomes higher at Band 3. The former is due to highoptical depth (see panel f), while the latter comes about along
be- the disk minor axis andscattering explanation at this shcause the inner disk has an optical depth of order unitypolarization
at
at the longest wave
Band 3, which is optimal for producing polarization throughazimuthal orientation, which coof scattering-induced polarizatio
scattering (Yang et al. 2017). In all cases, the polarization
mechanism(s) may
orientations are broadly similar to that displayed in Fig.polarization4e
2017; Stephens et al. 2017; Ya
for Band 7, with the deviation angles from the disk minor
spectral index between Band 7
axis increasing somewhat with wavelength.
at the rings, which is indicativthat may or may not be scatt
In Fig. 7, the predicted polarization spectra (polarizaGonzález et al. (2019) fitted thtion fraction as a function of wavelength) are plotted tions
for at di↵erent radii including
ing by spherical grains of powe
Figure7. Predicted polarization
the inneraxis:
disk atthree
representative
locations spectraon theformajor
the inner20
au(R=20;
bluesolidline),
the
D48
gap
(red
dashedline),
andthat the relatively lodisk at R = 20au, the D48 gap, and the B67 ring. In suggestedthethe B67 ring (orange dash-dotted line) along the major axis. The
small radii requires maximum g
gap,blackthe dottedpolarization
monotonicallywithline shows afraction
power-lawdecreases
that is proportional
to 3 .
range (see their Fig. 7). Such gr
increasing wavelength and approaches the steep power-law
to produce the observed polariz
3
in Band 7, (Yang et al. 2016, sefound no significant emission at this velocity (or anydust grains we used the CASA-based uvmultifit tool
【Imaging】
43. The big sibling of AU(Martí-Vidal
Mic: a cold etdust-richdisk toaroundwhereelse) in the CO data
cubes.
al. 2014)debris
that allows
fit models to the4
Moór
et al.
CP-72 2713
in the β visibilities.
Pic moving Priorgroupto this running, the data
In order to estimate an upper limit for the 12 CO (2–measured
Attila Moór et al. AJ, 159, 288; https://arxiv.org/pdf/2005.00861.pdf
1) integrated line flux we assumed that the possible COweights were recomputed using the statwt task. Simgas is co-located with the dust material in the disk.
ilarly to the case of PACS image analysis, we adopted
We defined a disk mask considering those regions ina Gaussian ring model witha fixed
FWHM width of【内容】βPic. moving group(24±3Myr)にある若いK7/M0
star CP-72
the millimeter continuum image that are adjacent with
0.2R. The obtained best-fit solution has R = 140±14 au
2713に付随するデブリ円盤のSpitzer/Herschel/ALMA
A cold dust-richdebris7m
disk array観測の解
5
!!
!!
◦
◦
◦
130◦CP−72± 162713
and yields
themaximum brightness and having SNR>2. Using the(3. 8±0. 4), i = 46 ± 12 , P A =aroundfound
noquotedsignificantring”
emissionat this velocity (or anydust grains
weused density
the CASA-based
uvmultifit
tool析。AU Micと主星の性質は似ているが遥かに半径が大きい。”Cold
a
flux
of
3.8±0.6mJy(theuncertainty
obtained mask we constructed the spectrum and calcuwhere else) in the CO data cubes.
(Martí-Vidal et al. 2014) that allows to fit models to the
In order to estimate
an upperlimit
for the 12 CO (2–measuredincludes
visibilities. Prior
to this calibration
running, the data uncertainty).
a
10%
We
used
the
lated its rms. To consider the correlation between theのみ。ringの場所がplanetesimalsのdynamical
excitationを表すとして起
1) integrated line flux we assumed that the possible COweights were recomputed using the statwt task. Simtaskimage
to analysis,compileimages ofbest-fit
and in the disk.
neighboring channels due to the√applied weighting funcgas theis co-locatedwith modelthe dust material
ilarly to tclean
the case of PACS
we adopted
源を議論(planetary stirring
We defined a disk mask considering those regions ina Gaussianvs.
ring self-stirring)。fractionalmodel with a fixed FWHM width of luminosityがtheresiduals(Figure1e-f).
No
significant
residual emistion, this rms was multiplied by 2.667 (see Matrà et al.
the millimeter continuum image that are adjacent with
0.2R. The obtained best-fit solution has R = 140±14 au◦◦◦
i = 46± 12◦ , P A =± 16disk
and yields
maximumbrightness
and having
SNR>2. Using the(3.!! 8±0.!! 4),
1.1x10-3と非常に大（d<40pcではこのレベルのFLは他に4つしかない）.
2019, and references therein). Since based on our resion
appears
in130theregion.theThe
fittedcenterof thea flux density of 3.8±0.6 mJy (the quoted uncertainty
obtained mask we constructed the spectrum and calcu!!
3 fromstellarposi- between thesults the bulk of the dust is located at radii !100 au andring uncertainty).
shows anWeoffset
includes model
a 10% calibration
used the of 1.
lated
its rms. theTo consider
the correlation
tclean task to compile images of the best-fit
model
neighboring channels!!due to the√appliedweighting func!!
!! and
!!
◦
3
±
0.
42,
∆DEC
=
−1.
22
±
0.
45).
(3σ confidence interval), by
tion
(∆RA
=
−0.
the
disk
inclinationis"82【用いた主なデータ】
the residuals (Figure 1e-f). No significant residual emistion, this rms was multiplied by 2.667 (see Matrà et al.
2019, and references
therein). Since
based on our resion appears
in thedisk region. Theaccuracy
fitted center ofof
the an ALMA
The
positional
observation
withadopting a stellar mass of 0.71 M% (see Sect. 3.1) we!!
sults the bulk of the dust is located at radii !100 au andmodelring shows
an byoffsetGaiaof 1. 3 from the stellar posi• Distance = 36.16± 0.03pc−1!!
◦
0.!! 42, ∆DEC = −1.
22 ± be0.!! 45).estimated(3σ ofconfidence
interval), expect
bytion (∆RA= −0.!! 3 ±calibration
the disk inclinationis "82. Using theamaximum
line
width
of
5
km
s
normal
can
as
the
ratio
the
• Spitzer/MIPS : 24&
70μm
The positional accuracy of an ALMA observation withadopting !!
a stellar mass of 0.71 M% (see Sect. 3.1) we
9 in ourcase,corrected rms value and this linewidth we derived a 3σachieved
resolution
and
i.e.
∼0.
expectamarginallymaximumline widthof 5thus
km s−1 . Using thecan beestimatedas theratio SNR,of atthe 100μm,• Herschel/PACS normal:achieved
100calibration
&
160μm
(FWHM
≈8,5”
corrected rms value and this linewidth we derived a 3σ
resolution and SNR, i.e. ∼0.!! 9 in our case, thus
−1
12
for
the
CO (2–1) intethe
observed
offsetisnotsignificant.
The
best-fitting
upper
limit
of
0.216
Jy
km
s
the observed offset is not significant. The best-fitting
upper limit of 0.216 Jy km s−1 for the 12 CO (2–1) interesolved)
gratedline flux.
disk parameters
derivedfrom the analysis
of PACSfrom
andgrated line flux.
disk
parameters
derivedtheanalysisof PACS and• ALMA(ACA 7m) ALMAat 1.33mm
&
ATCA
at
8.8mm
data are summarized in Table 2.
Table2.2. Best-fitting dust disk parameALMA data are summarized in Table
DEC offset [arcsec]

10

0

-10

10

ters derived from the analysis of PACS and
ALMA data

Table 1. Measured continuum fluxes and predicted photospheric fluxes
for CP−72 2713

Data setPA (◦ )i (◦ )R (au)

+60
−64

+26
−21

+22
−26

Meas. flux1. Instrument
Photosph.
flux
Reference
Table
Measuredcontinuum
fluxes and predicted
photosphericfluxes
PACS 100µm
153
29
99
(µm)for (mJy)
ALMA 1.33 mm 130±16 46±12 140±14
CP−72 2713 (mJy)
a

λ

9.00
11.56
22.09
23.68
71.42
100b
160
1330b
8820

【Photometry】

128.1±12.3
60.0±2.818.9±1.5
16.1±0.671.1±5.6
110.6±9.7
103.9±16.9
3.80±0.59
0.096±0.023

λ

(µm)

IRC
WISE
WISE
MIPS
MIPS
PACS
PACS
ALMA
ATCA

105.2
65.0
a 18.3
16.0
1.8
0.9
0.35
0.005
10.6×10−5

Meas. flux(mJy)

Ishihara et al. (2010)
Cutri & et al. (2013)
Cutri & et al. (2013)
SEIP
This work
This work
This work
This work
Norfolk et al. in prep.

Instrument

9.00
128.1±12.3
IRC
11.56
WISE
a The quoted
flux densities60.0±2.8
are not color corrected.
22.09
18.9±1.5
WISE
b The emission
is marginallyresolved at these wavelengths.
23.68
16.1±0.6
MIPS

Photosph. flux
(mJy)

Reference

3. RESULTS AND ANALYSIS

3.1. Basic stellar properties of CP−72 2713

105.2
al. (2010)
To estimate theIshihara
basic stellaretparametersof CP−72 2713
and provide photosphericat infrared
65.0
Cutri &flux
et predictions
al. (2013)and millimeter wavelengths, we modeled the stellar
18.3
Cutri & et al. (2013)
photosphere by fitting an ATLAS9 atmosphere model
16.0(Castelli & Kurucz 2004) to SEIP
the photometric data of the0-10
RA offset [arcsec]

Table 2. Best-fitting dust disk parameters derived from the analysis of PACS and
ALMA data
Data setPA (◦ )i (◦ )R (au)

PACS 100µm

153+60
−64

29+26
−21

99+22
−26

ALMA 1.33 mm

130±1646±12140±14

Moór et al.

Stellar Parameters

CP−72 2713 exhibits flares and rotational modulation
due to stellar spots in its TESS light curve (Fig. 7). Consistently with previous ASAS based results of 4.456 days
ATLAS9 atmosphere model (Catelli & Kurucz 2004による)で(Kiraga 2012) and 4.48 days (Messina et al. 2010) we ob星のSEDデータをfitting,
MESA& Stellar
Trackstained a rotation periodof PIsochrones= 4.437 d froma Fourieranalysisof 2016;
our light
curve.
In addition,
we identified 51(MIST,
Choi+
Dotter2016)
で年齢を見積もり。
flares or group of flares. We found that the star spent
14+5about
Pic Moving
group(24±3Myr)の一員と考え-3 Myrと，β936 minutes
in a flaring
state, which is 1.24% of
the time covered by TESS.
て良い結果を得た。
Table 3. Stellar parameters
Parameter

Data

References

Distance (pc)
Spectral type

Teff (K)L∗ (L" )R∗ (R" )M∗ (M" )36.62±0.03 pc
K7Ve
K7IVe
M0
3900±70 K
0.18±0.01
0.94±0.04
0.71+0.03−0.05

Bailer-Jones et al. (2018)
Torres et al. (2006)
Pecaut & Mamajek (2013)
Gaidos et al. (2014)
This work
This work
This work
This work

LX /Lbolv sin i∗ ( kms−1 )-2.92
7.1±0.9aProt (days)
vrad (kms−1 )4.437
8.0±0.4a

Kiraga (2012)
Torres et al. (2006)
Weise et al. (2010)
This work
Torres et al. (2006)
Lindegren et al. (2018)a Weighted average of the values reported by the cited papers.

estimate of 0.2
counts for thethe distance oftential systema
introduce even(e.g. Miyake &1994) and theized by a single
in our error ca

100
Flux density [Jy]& et al.
etry isspots
nt 3D
that iso et al.
ds our= 0 inog g ofempernferred
the oba stelr masstainedal preTracks
coursecribedM! andctively.
is sat3 Myr,r massy conmeters

fore changing the λ0 parameter has no effect on our results. We note that though most debris disks show considerably steeper millimeter slopes (Gáspár et al. 2012;
MacGregor et al. 2016; Marshall et al. 2017), interestingly, the excess SED of AU Mic can also be well reproduced by a blackbody (Matthews et al. 2015). For thefractional luminosity we obtained fd = 1.1 × 10−3 .

10-1
10-2
10-3
10-4
10-5
1

2MASS
IRC
WISE
MIPS
PACS
ALMA
ATCA
PHOTOSPHERE
MBB MODELACE MODEL

10

100
1000
Wavelength [µm]10000

Figure 2. Color-corrected SED of CP−72 2713. The green
curve shows the best fitting modified blackbody model. As
a comparison, the orange curve shows an Analysis of Collisional Evolution (ACE) model for the disk with a total
mass of 50 M⊕ was also overplotted (see Sect. 4.3). Thisshows the model SED after 24 Myr of evolution assumingṀ∗,wind = 300Ṁ#,wind .

3.3.1. The dust component

To derive anusing the mea
limit, we needthat populatetransition. Co
gas constituen(Matrà et al. 2of the rotation
is difficult to ju
cesses in the exthat local thercollisions domi
emission is op
between 10 anin the range of

3.4. Sourc

So far in ou
served excess estellar disk. Kn
at this point it
confusion as w

Flux density [Jy]Disk Properties

100
10-1
10-2
10-3
10-4
10-5
1

10

100
1000
Wavelength [µm]10000

Dust removal mechanisms·
強いStellar wind (M* = 300M⊙)によるmass lossを考慮すると，サイズ0.1μm以下のダストは吹き飛ばされうるものの，より大きなダストは吹き飛ばされない（右図）

Collisional evolution

12

Moór et al.
100

10-1

ただし，Poynting-Robertson Eﬀectや

10-2

Stellar windによるdragsがあると，Analysis of Collisional Evolution(ACE, Krirov+ 2006; Lohne+ 2008)を使いシミュレーション（右図）

10

-4

10-5
Normalized surface brightness

ダスト粒子の公転運動角運動量が失われ，中心星に比較的短い時間で落下する。
それを考慮した考察が必要

10-3

10-6
100
10-1
10-2
10-3
10-4
10-5
10-6

100µm - 1Myr
100µm - 10Myr
100µm - 24Myr
1.3mm - 1Myr
1.3mm - 10Myr
1.3mm - 24Myr

14

24Myr

xm =
10

xm =1
0

xm=1

xm=1

xm =
10

tonset (Myr)

xm =1

0

xm=1

xm=1

Stirring ofplanetesimal belts

tonset (Myr)

Moór et al.
Particle concentration models predict a significquicker formation of larger bodies
even farfromthe radiusof thep50
star(at r(classical)>100 au) through thegravitationalcolSelf-stirring
region of the
Self-stirring (pebble concentration)
30
Planetaryconcentrated
stirringof locally
pebbles of(Johansen
et al.
the system.
Location of the planetesimal belt
Krivov & Booth 2018, and references
Ac
Thetherein).
collisional40
ing to these scenarios large planetesimals
areforalr
feasible way
present in the disk by the timeParticleof the dissipati
concentr
20
14
Moór et al.
the gas-rich protoplanetary material.
to thequickerDue
formation
sibly smaller size of the largest starbodieshun(at(fewr >10
Planetesimal beltからcollisional cascadeを起こしdust小30
kilometers, Simon et al. 2016; Schäfer
et al.concen
2017)
of locally10respect to the classical model (where
the&1000
km2
Krivov
Booth
粒子を連続的に供給しringを作るには，
20
40
60
80
100
120
140
160
planetesimals can stir the neighbouring
ing to diskthesepromp
scenRadius (au)Planetesimal beltがdynamically excitedなことが必要。
the excitation of the surroundingpresent
smallerinbodies
the di20
more time than in the previous scenario.
AccordiFigure 6. Onset time of the stirring as a function of stelthe gas-rich
prot
その励起源を検討。円盤ガス散逸後の流れは‥
locentric radius for different stirring scenarios (Sect. 4.6).
Krivov & Booth (2018), in this scenario
the onset
sibly smallersize
The red dot corresponds to the observed property of theof effective self-stirring at a radiusr can beSimoncompkilometers,
CP−72 2713 system.
10
asrespect to the cla
20
40
60
80
100
120
140
160
"! "!
"−1 !
"4
!
planetesimals
canRadius (au)
ρ
1
vfrag
129
微惑星の衝突 →衝突カスケード → 小粒子落下0.18(M∗/M" )xm (r/30au)−3/2 g cm−2 , where xm is apebbletonset =−3the excitation
−1 of
x
γ
1gcm
30ms
m
scaling factor (for the minimum-mass
more time(Onset of stirring)による小粒子生成
Figure 6.solar
Onsetnebula
time of the stirring as!a function
of!stel- ""−3"3 than i−3/2 !
providedan different
anxm ≡ 1), Kenyon & Bromley (2008)M∗
r & Booth (Smax(Sect. 4.6).
locentric
radius forstirring scenarios
Krivov(Myr)
alytical formula for the formation
time
1000
km-size to the observed200km
The
redofdotcorresponds
property of M
the"of 100au
effective self-st
Onset of dynamically stirringの時間スケールを planetesimals (the onset time of CP−72
2713 system.
the stirring)at a stelas
20
40
60
80
100
120
140
160
where
γ
is
a
parameter
ineq.
20
in Krivov & B
locentric
radiusr:!
Radius (au)
いくつかの場合で検討
vf
129−3/2(2018),−2 ρ is the bulk density of planetesimals,
pebble/M)x(r/30au)
g
cm
,
where
x
is
a
"30.18(M!
"!
∗ 3/2" mmtonset =2Mrtheminimum
collisional
velocityneededfor fragm
"
xm
−1.15
scalingfactor(fortheminimum-masssolar
nebula
Myr.
(4)
=
145x
tclassical
24Myrという年齢で
onset
m
!
is
the
size
of
the
largestplanetesi
tion,
and
S
80
au
M
max
∗
xm ≡ 1), Kenyon & Bromley (2008) provided an an- Classical self-stirring（ダスト成長しつつ破片が出る）
Smax
Figure
6
shows
–
by
using
the
default
valuesprop
比較的遠方でダストリングを作るにはalyticalformulafortheformation
time
of
1000
km-size
200km
Realistically, the xm factor could not be higher than 10
byKrivov&
Booth(2018)
of
γ
=
1.5,
ρ
= 1gc
Self-stirring(Pebbleconcentration:
局所的な自己重力不安定
planetesimals
(the
onset
time
of
the
stirring)
atastel初期質量が大きい円盤で(Mustill & Wyatt 2009). Taking a disk with the above−1=
30ms, and Smax = 200kmv
fragwhere– γthatis atheparlocentric
radiusr:
により，ガス散逸直後から~100kmの微惑星がある場合)described radial surface density profile with xm = 微惑星より大きな天体による強めの摂動が必要
10,edge of the possibly stirred region(2018),(purple ρdashed-dis the b
"
!
"
!
3
3/2
and assuming that the location of the planetesimal belt
to the position of thetheminimum
planetesimal
2Mcloser line) is
col"classical−1.15- Planetary stirring (惑星によるexcitation)
Myr.
(4)
=
145x
t
(classicalself-stirringは不可)
at 140 au in the CP−72 2713 system onset
corresponds tomthe
CP−722713 system. Hereagain
istion,weandSmaxado
80 auin the M∗
outer edge of the initial protoplanetary disk, the totalof 6themode–
xm = 10. Considering the caveats
Figureshows
3.
initial mass of solids would havebeen ∼1.2×10
Realistically,
the xM
factorcould
not be higherofthan
10
m⊕
the
uncertainties
thefundamental
model &
param
by KrivovBoo
xm : minimum mass Solar Nebula を基準にした初期質量のスThis hugely exceeds the measured
dust massesof 2009).
pro- Taking(Mustill& Wyatt
a disktheabove−1(Krivov& withBooth2018),
it meansvfrag
that=the30ms
excitati,
toplanetary disks around similardescribed
mass starsradialin theLu- density
10,be consistentsurface
profile with
m =canthe CP−72
2713 xdisk
the p
ケーリングファクター(Kenyon & Bromley 2008)edge of withthe possib
pus and Chameleon I star formingand
regions(Ansdellal. location of the planetesimal belt
assuming
thatetthe【観測天体】円盤（星）距離をGaiaで導出，40. X-Shooter survey of disk accretion in Upper Scorpius I. Veryhigh accretion rates at age>5 Myr
C.F. Manara et al. A&A, 639, A58; https://arxiv.org/pdf/2004.14232.pdf

過去のALMA論文のデータも再計算
Manara et al.: X-Shooter survey of disk accretion in Upper Scorpius
Table 2. Stellar, disk, and accretion properties for the targets in the Upper Scorpius region
Name【内容】年齢 >5Myrの範囲にあると言われているUpper Sco (~145pc)に含まれALMAによるダスト円盤検出があった(つまり比較的massiveな円盤)36個をλ=300-2500μmで高分散分光，質量降着有無を調査。
【背景】円盤質量の散逸タイムスケールや古典的粘性降着モデルが妥当かどうかの検証には，1-3Myrの天体だけでなく，年齢の異なる天体も調査対象に含め，時間軸上での傾向をとらえる必要あり【キーとなる観測量】質量降着率と，円盤質量や中心星質量との関係，
Manara et al.: X-Shooter survey of disk accretion in Upper Scorpius
およびその分散 Manara
et al.: X-Shooter survey of disk accretion in Upper Scorpius

2MASSJ15534211-2049282
2MASSJ15583692-2257153
2MASSJ16001844-2230114
2MASSJ16035767-2031055
2MASSJ16035793-1942108
2MASSJ16041740-1942287
2MASSJ16041893-2430392
2MASSJ16042165-2130284
2MASSJ15354856-2958551_E
2MASSJ15354856-2958551_W
2MASSJ15514032-2146103
2MASSJ15530132-2114135
2MASSJ15582981-2310077
2MASSJ16014086-2258103
2MASSJ16020757-2257467
2MASSJ16024152-2138245
2MASSJ16054540-2023088
2MASSJ16062196-1928445
2MASSJ16063539-2516510
2MASSJ16064385-1908056
2MASSJ16072625-2432079
2MASSJ16081566-2222199
2MASSJ16082324-1930009
2MASSJ16082751-1949047
2MASSJ16090002-1908368
2MASSJ16090075-1908526
2MASSJ16095361-1754474
2MASSJ16104636-1840598
2MASSJ16111330-2019029
2MASSJ16123916-1859284
2MASSJ16133650-2503473
2MASSJ16135434-2320342
2MASSJ16141107-2305362
2MASSJ16143367-1900133
2MASSJ16154416-1921171
2MASSJ16181904-2028479

disk
type
Full
Full
Full
Full
Full
Full
...
Transitional
Full (binary)... (binary)
Evolved
Full
Full
Full
Full
Full
Full
Transitional
EvolvedEvolved
Full
Full
Full
Evolved
Full
Full
Full
Full
Full
Full
Full
Full
Full
Full
Full
Evolveddist[pc]
136 ± 4166 ± 4138 ± 9143 ± 1
158 ± 2161 ± 2145
150 ± 1
145
145
142 ± 2146 ± 2147 ± 3145
140 ± 1142 ± 2145 ± 2145
139 ± 3144 ± 7143 ± 2140 ± 2138 ± 1145
139 ± 3138 ± 1
158 ± 5143 ± 3155 ± 1
139 ± 2145
145
145
142 ± 2132 ± 2138 ± 2

SpT
M4
K0
M4.5
K6
M2
M3
M2
K3
M4.5
M4.5
M4.5
M4.5
M4.5
M3
M2
M5.5
M4.5
M1
M4.5
K7
M3
M2
M0
M5.5
M4.5
M0
M4.5
M4.5
M3.5
M1
M3
M4.5
K4
M3
K7
M5

T e↵
[K]
3270
5110
3200
4205
3560
3415
3560
4730
3200
3200
3200
3200
3200
3415
3560
3060
3200
3705
3200
4060
3415
3560
3850
3060
3200
3850
3200
3200
3340
3705
3415
3200
4590
3415
4060
3125

AV
[mag]
1.2
0.0
0.8
0.7
0.3
0.7
0.3
1.4
0.0
0.0
0.3
0.8
1.0
1.2
0.4
0.6
0.6
0.8
0.0
0.4
0.7
0.5
1.1
0.6
0.3
1.0
0.5
1.2
0.6
0.6
1.0
0.3
0.3
1.9
2.8
1.6

L?
L
0.09
2.57
0.08
0.48
0.13
0.14
0.45
0.90
0.10
0.10
0.05
0.05
0.05
0.12
0.08
0.03
0.10
0.34
0.03
0.29
0.18
0.15
0.32
0.06
0.05
0.32
0.04
0.04
0.03
0.22
0.11
0.12
1.05
0.52
0.30
0.05

logLacc
L
-2.6
-0.5
-1.9
-1.8
-5.1
-4.3
-3.1
-3.2
-2.8
-2.9
-3.5
-3.0
-2.3
-1.2
-3.8
-2.9
-2.8
-1.3
-5.1
-2.3
-2.6
-3.7
-2.0
-3.1
-4.2
-1.7
-4.5
-3.9
-1.9
-2.3
-1.6
-2.3
-1.4
-2.7
-0.3
-3.4

M?
M
0.24
1.63⇤
0.20
0.91
0.42
0.31
0.37
1.24
0.20
0.20
0.19
0.19
0.19
0.31
0.44
0.12
0.20
0.46
0.18
0.82
0.29
0.41
0.61
0.14
0.19
0.60
0.18
0.19
0.27
0.50
0.32
0.20
1.25
0.29
0.81
0.16

Ṁacc
M /yr
3.66·10 10
1.59·10 08
2.03·10 09
8.81·10 106.69·10 13
6.04·10 12
1.48·10 10
3.09·10 113.53·10 102.73·10 105.01·10 11
1.52·10 107.16·10 107.42·10 09
1.08·10 11
2.76·10 103.58·10 106.13·10 09
8.62·10 13
2.65·10 104.56·10 10
1.99·10 117.90·10 10
1.97·10 10
1.02·10 11
1.74·10 09
4.54·10 12
1.45·10 119.77·10 104.75·10 102.93·10 09
1.18·10 09
2.09·10 09
5.17·10 102.44·10 08
8.05·10 11

Acc?
Y
Y
Y
Y
N
N
Y
N
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y
N
Y
Y
N
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y
Y

Mdisk,dust
M
1.69·10 06
1.51·10 04
2.31·10 06
2.74·10 06
9.16·10 07
7.26·10 07...
1.55·10 04
1.27·10 06
1.27·10 06
4.82·10 07
3.88·10 06
4.00·10 06
2.28·10 06
3.25·10 06
6.46·10 06
5.05·10 06
2.69·10 06
1.03·10 06
5.48·10 07
8.39·10 06
5.98·10 07
2.58·10 05
5.01·10 07
1.05·10 06
2.81·10 05
6.78·10 07
1.14·10 06
3.69·10 06
3.65·10 06
5.80·10 07
4.97·10 06
3.15·10 06
7.84·10 07
1.28·10 05
2.76·10 06

Notes. Disk type from Barenfeld et al. (2016); Luhman & Mamajek (2012); Carpenter et al. (2006). Stellar properties obtained using the Bara↵e etal. (2015) evolutionary models, apart from the target 2MASSJ15583692-2257153, for which Siess et al. (2000) models were used since the stellar
mass was higher than the maximum one modeled by Bara↵e et al. (2015). Disk masses are updated from Barenfeld et al. (2016) using the distanceinferred from the Gaia DR2 (Gaia Collaboration et al. 2018) parallaxes. When no uncertainties on the distance is reported, the mean distance tothe targets of 145 pc was adopted. Possible non-accretors are reported with "N" in the Acc? column. The values reported here for the accretion
rate of non-accretors are considered in this work as upper limits.

Fig. 2. HR diagram for objects in Upper Scorpius observed here. The
evolutionary tracks are from Bara↵e et al. (2015), with isochrones for
Fig.
2. HR
diagramforlineobjects
in Upper
Scorpius
observed here. The
1.2, 3, 5,
10, and
30 Myr.
The redis the medianof theL? in di↵erent
T e↵ bins.
evolutionary tracks are from Bara↵e et al. (2015), with isochrones forFig. 1. Disk mass vs stellar mass after correcting the disk masses for theGaia estimated distances and using the stellar masses derived here. Orange
circlesmass
are used
for targets
observedwith X-Shooter,graymasses
symbolsfor the 1.2, 3, 5, 10, and 30 Myr. The red line is the median of the L? in di↵erent
Fig.
1. Diskvs stellar
mass
after correcting
the diskfor estimatedtargets not observed,
for ALMA
detections,
downwardfacingGaia
distancescircles
and using
the stellar
massesderivedhere. Or- T e↵ bins.
triangles for ALMA upper limits. The shaded regions delimit the disk prises various Class III photospheric templates with di↵erent
ange circles are used for targets observed with X-Shooter, gray symbols
mass ranges where the sample is complete with respect to the Baren- spectral types (SpT) from G- to late M-type taken from Manara
forfeldtargets
observed,
ALMA
detections,
downward
et al.not(2016)
one. Thecircles
dashedforgraylines delimit
the bins
used in thefacing
et al. (2013a, 2017b), di↵erent slab models, and extinction valtriangles
for ForALMA
uppernotlimits.
Thewithshadedregionsthedelimit
the disk
prisesvariousClass
III law
photospheric
templateswith di↵erent
discussion.
the objects
observed
X-Shooter,distances,
ues
(A
),
assuming
thereddeningby
Cardellietal.
(1989)
V
mass
rangeswhere
the from
samplecompletewith respect to the Baren- spectral types (SpT) from G- to late M-type taken from Manara
disk,
and stellar
massesthe isliterature
are adopted.

The best fit of the Balmer continuum emission
feld et al. (2016) one. The dashed gray lines delimit the bins used in andthe RVet=al.3.1.
(2013a,2017b),
di↵erentslab
andfitextinction valare shown in Figs. C.1-C.6.
The
integrated
fluxmodels,of the best
discussion. For the objects not observed with X-Shooter, the distances,

dard stars at the beginning and at the end of theduring the
Appendix A: Log of★
theobservations,Rを切り分け★と，観測で得られたスペクトルからL★night• 見積もったM
first night, and at the beginning, in the middle, and at the end of
The observations were carried
out
in
two
di↵erent
observingprothe night during the second night. The reduction led to consistent·grams. Eight targets were observed in the Service Mode program results with all standard stars. For the reduction of the data obから，Maccinは，たL
acc
Pr.Id.
097.C-0378 (PI Manara)
the period from July to August tained in the first night, we adopted the standard star observed at2016. Typically, these observations were carried out with image the end of the night. For the data obtained in the first part of thequality in the VIS arm of ⇠1L(see Table
A.1). The standard star second night we used the standard star observed in the middle of
R
·
acc
★
observed at the beginning of the night as part of the standard ca- the night, while we used the one observed at the end of the night
Maccplan∼for 1.25
libration
X-Shooter was used for the で求める。
flux calibration of for the spectra obtained in the second half of the night, starting
the spectra.
from and including 2MASS J16072625-2432079.
GM★
The remaining 28 targets discussed here were observed dur00

質量降着率 vs. Mstar & Mdust ：（下2枚）他の領域との比較
A&A proofs: manuscript no. usco_survey_lang_edited

年齢が異なる領域間で大きな違いなし·
Maccの中央値が必ずしも低下していない·τ ≡ Mdust / Macc も年齢が若いものと同じくらい大きな分散があり，この傾向は古典的質量降着モデルで説明できない（分散がもっと小さくなるはず）

Fig. 3. Mass accretion rate vs stellar mass for the targets in the Upper
Scorpius region (orange points, both in the upper and bottom panels)
and for the targets in the Lupus and Chamaeleon I regions (gray symbols, bottom panel). The downward facing triangles are used for nonaccreting objects, transition disk objects are highlighted with a circle
around their symbols. The cross indicates the typical errors on the measurements.

line emission the results discussed in the following would not
be a↵ected. We nevertheless note that Ṁacc could be underestimated for the objects with the lowest accretion rates, which are
typically below the chromospheric noise.
For some of the targets with the lowest measured accretion
rates, the ratio Lacc /L? falls below the typical values for chromo-

Fig. 4. Mass accretion rate vs disk mass for the targets in the Upper Scorpius region (orange points, both in the upper and bottom panels) and for the targets in the Lupus and Chamaeleon I regions (graysymbols, bottom panel). The dot-dashed lines report di↵erent ratios of
Mdisk / Ṁacc : 0.1 Myr, 1 Myr, and 10 Myr, as labelled. Symbols as inFig. 3.

Figs. C.1-C.6), in line with other estimates to confirm the accretion status of a young stellar object (e.g., Herczeg, & Hillenbrand 2008; de Albuquerque et al. 2020). We note that for
the non-accreting targets in the Upper Scorpius region, the mea˙より若い領域との比較
A&A proofs: manuscript no. usco_survey_lang_edited

Manara et al.: X-Shooter survey of disk accretion in Upper Scorpius

Fig. 6. Comparison between median and percentiles of the mass accretion rates as a function of disk mass for the Lupus and Upper Scorpius
regions (left), and for the Upper Scorpius region and the expectations from viscous models at 1 Myr and 8 Myr (right).

Fig. 5. Mass accretion rates vs disk dust masses for the targets in the Lupus and Chamaeleon I star-forming regions, and in the Upper Scorpius
region. The dashed lines report the 16th and 84th percentiles, and the solid line the median of the distributions.
Table 1. Median values for the Ṁacc - Mdisk relation

Disk mass bin
4.8 · 10 5

1.7 · 10 4

6.47 · 10

4

Lupus

1.7 · 10 4

6.47 · 10 4
1.55 · 102

Chamaeleon I

Upper Sco

Viscous 1 Myr

Viscous 8 Myr

Median

Spr.

Ndata

Median

Spr.

Ndata

Median

Spr.

Ndata

MedianSpr.

Median

Spr.

9.77

0.55

5/0/1

10.26

1.22

14/0/11

10.30

1.87

15/4/0

10.01

1.11

11.04

1.56

9.94

1.58

15/0/0

9.48

2.56

20/4/49.10

1.12

14/0/0

9.53

0.76

10.50

1.18

9.17

1.13

32/4/08.39

1.93

36/1/0

8.93

1.81

6/1/0

8.91

1.35

9.65

1.08

Notes. Mdisk = 100 * Mdisk,dust in M . The table reports the values of log Ṁacc , reported in M /yr, for the median, and for the spread of the distribution,
defined as the di↵erence between the 16th and 84th percentile of the distribution in a given bin of Mdisk . The latter is equivalent to a 2 spread. The
Ndata columns report the total number of tagets included in the bin / the number of non accretors in the bin / the number of undetected disks in thebin.

region of Chamaeleon I, and larger than the viscous evolution
model evolved at 8 Myr. This is particularly true when we compare the spread obtained fitting the model at 8 Myr using thelinmix tool, = 0.4 dex, with the data in the Upper Scorpius
region, that have a spread with this method of = 1.3 dex (seeSect. 4). We thus observe that the models able to reproduce theOn the other hand, external photoevaporation would predict that the disks have low mass, while still low values of
Mdisk / Ṁacc ⇠0.1 Myr (see Fig. 4, and Rosotti et al. 2017; Sellek
et al. 2020), more in line with what is observed here. In this context, it is worth mentioning that the environment of Upper Scorpius is di↵erent than the one of Chamaeleon I and Lupus, having

observed. Further studies on variability in older star-forming regions is needed. The mass-budget issue given the observed Ṁacc
and Mdisk is therefore important. Indeed, under the assumptionthat Ṁacc is constant with time, these high values of Ṁacc would
imply that over the lifetime of the disk 10 8 M /yr ·107 yr ⇠ 0.1
M of disk gas mass is accreted from the disk onto the star.
Assuming a gas-to-dust ratio of 100, this means that a total of
Mdisk,dust ⇠ 10 3 M was accreted. This value is in line with themost massive disks observed in the Lupus and Chamaeleon I
disks, which could indeed be the progenitors of the survived
disks observed here. Such high mass would probably imply thatthese disks were gravitationally unstable at the beginning of theirlives (e.g., Kratter & Lodato 2016). One possibility could bethat accretion becomes active at later ages, as predicted by somemodels of MHD disk winds driven accretion (e.g., Armitage etal. 2013, for the case of a constant di↵erential magnetic flux).

making the disks in Upper Scorpius as massive as the younger
disks in Chamaeleon I and Lupus, it is unclear whether a better
match with the observed age-independent spread could be obtained with such models. Mulders et al. (2017) already showedthat a simple scatter in the values of the gas-to-dust ratio alone
cannot reproduce the observed scatter in the Ṁacc -Mdisk relation
with no need for other sources of scatter, such as accretion variability.

Lupus(age ~1Myr)の結果をよく再現する円盤分布を8Myr進化させてもUScoの結果を再現しない
5.4. Does the mass accretion rate decrease with time?

The detection of strong accretors at later ages, and the connected
hint of a lack of a general decrease of accretion rates with time
when the targeted stars are still hosting a disk, has already been
observed in di↵erent older star-forming regions: the nearby loose
associations TWA (Venuti et al. 2019) and ⌘-Cha (Rugel et al.
2018), the more distant -Velorum cluster (Frasca et al. 2015),
Orion OB1b and Orion OB1a associations (Ingleby et al. 2014),
and even the very massive regions like NGC3603 or 30 Doradus(De Marchi et al. 2017). Individual targets have also been foundto be still accreting at age>20 Myr (e.g., Mamajek et al. 2002;
Zuckerman et al. 2014; Murphy et al. 2018; Lee et al. 2020).
While this appears to be in contrast with evidence of a decreaseof Ṁacc with individual ages of young stars (e.g., Hartmann et al.
1998; Sicilia-Aguilar et al. 2010; Antoniucci et al. 2014; Hartmann et al. 2016), it should be noted that Da Rio et al. (2014)
showed that correlated uncertainties on the determination of stellar parameters from the HR diagram can lead to spurious corre˙

古典的粘性降着だけでない要因？

(Photoevaporation, dust 進化，MHD wind)5.3. The need to account for dust evolution

It is worth mentioning again that the assumption Mdisk =100·Mdisk,dust even after ⇠5-10 Myr of disk evolution is possibly
incorrect. As shown by global models of dust and gas evolution(e.g., Birnstiel et al. 2010; Rosotti et al. 2019), dust radial drift
is in general more efficient than gas accretion, implying that the
dust-to-gas ratio is a decreasing function of time. Depending on
the disk parameters (such as the efficiency of grain growth andthe disk size), there can be an initial period of time, lasting ⇠12 Myr, in which the assumption Mdisk ⇠ 100·Mdisk,dust is almost

ガス・ダスト比のばらつき？

The Age-Dependence of Mid-Infrared Emission Around Young Star Clusters#36

Zesen Lin1,2,3 , Daniela Calzetti3 , Xu Kong1,2 , A. Adamo4 , M. Cignoni5 et al.
Key Laboratory for Research in Galaxies and Cosmology, Department of Astronomy, University of Science andTechnology of China, Hefei 230026, China; 2 School of Astronomy and Space Sciences, University of Science andTechnology of China, Hefei, 230026, China; 3 Department of Astronomy, University of Massachusetts, Amherst, MA
01003, USA; 4 Department of Astronomy, Oskar Klein Centre, Stockholm University, AlbaNova University Centre,
SE-106 91 Stockholm, Sweden; 5 Department of Physics, University of Pisa, Largo Pontecorvo, 3, I-56127 Pisa, Italy
1E-mail contact: zesenlin at mail.ustc.edu.cn系外銀河の星形成領域における8μmダスト(PAH)放射の年齢依存性を議論
Lin et al.
NGC4395

NGC4449

1

NGC1313

1

NGC3738

1

1

Accepted by ApJ

1Figure 1. GALEX two-color (far-UV and near-UV) images for our five galaxies. The name of each galaxy is listed at theupper-left corner of each panel. The white bar indicates the 10 scale for each panel.
Table 1. Properties of Selected Galaxies

PAH Emission as a Function of Age
F657N

NGC F275W
7793 NGC 4395

Properties
Morph.

F555W
SAd
F814W
3.44

a

Dist. (Mpc)b

12 + log(O/H)

c

NGC 4449

NGC
1313
IRAC4

9

SAm

IBm

SBd

Im

4.30

4.31

4.39

4.90

8.50

8.19

8.26

8.21

8.10

6.0 ⇥ 108

1.1 ⇥ 1092.6 ⇥ 1092.4 ⇥ 108

378

108

284

674

164

N (IRAC4 clumps)f

48

17

25

47

9

NGCN4395-S(final) 622

39

17

16

23

2

N (YSC)eg

8

Lin et al.
NGC7793
NGC4395
NGC4449
NGC1313
NGC3738

NGC 3738

3.2 ⇥ 109

M⇤ (M )d

http://arxiv.org/pdf/2005.02446

a Morphological type as listed in the NASA/IPAC Extragalactic Database.

F657N
F275W
IRAC4
b Redshift-independentor flow-corrected redshift-dependent distance in Mpc, taken from Calzetti et al. (2015).

10

2

101

100

F555W
c Central oxygen abundances of the galaxies derived
from the empirical calibrations (Pilyugin & Thuan 2005; Pilyugin et al.
F814W

10 1

e Number of identified YSCs; see Section 2.2.
f Number of IRAC4 sources identified as YSC counterparts; see Section 3.2.
g Number of IRAC4 sources in the final sample; see Section 4.

10

NGC7793
NGC4395
NGC4449
NGC1313
NGC3738

2

101

100

agemax /agemin > 3

2012; Pilyugin et al. 2013) or the direct method (Pilyugin et al. 2010, 2012). The values for NGC 7793, NGC 4395, and NGC
1313 are taken from Pilyugin et al. (2014), while the ones for NGC 4449 and NGC 3738 are taken from Pilyugin et al. (2015).
d Stellar masses in M , taken from Calzetti et al. (2015).

24
L ,8/M [L /M ]NGC7793

L ,8/M [L /M ]4

Using the star cluster catalogs from the Hubble Space Telescope program Legacy ExtraGalactic UV survey (LEGUS)
and 8 µm images from the IRAC camera on the Spitzer Space Telescope for 5 galaxies within 5 Mpc, we investigate
how the 8 µm dust luminosity correlates with the stellar age on the 30–50 pc scale of star forming regions. We
construct a sample of 97 regions centered at local peaks of 8 µm emission, each containing one or more young starcluster candidates from the LEGUS catalogs. We find a tight anti-correlation with a Pearson correlation coefficientof r = −0.84 ± 0.05 between the mass-normalized dust-only 8 µm luminosity and the age of stellar clusters younger
than 1 Gyr; the 8 µm luminosity decreases with increasing age of the stellar population. Simple assumptions on a
combination of stellar and dust emission models reproduce the observed trend. We also explore how the scatter of theobserved trend depends on assumptions of stellar metallicity, PAH abundance, fraction of stellar light absorbed bydust, and instantaneous versus continuous star formation models. We find that variations in stellar metallicity havelittle effect on the scatter, while PAH abundance and the fraction of dust-absorbed light bracket the full range of thedata. We also find that the trend is better explained by continuous star formation, rather than instantaneous burst
models. We ascribe this result to the presence of multiple star clusters with different ages in many of the regions.
Upper limits of the dust-only 8 µm emission as a function of age are provided.

fU,sel /fU,ap < 10%

10 1

L⌫,8 / 8 < 3

106

107

108

109

ageM [yr]

106

107

108

ageM [yr]109NGC 1313-E 640

Figure 4. ⌫L⌫,8 /M⇤ as a function of ageM⇤ for all 146 identified IRAC4 counterparts (left) and our final sample (right). The

Figure 5. F657N (H↵; left), optical RGB (middle), and IRAC4 (right) images of two outliers in the right panel of Figure 4:
cyan circles, magenta upward-pointing triangles, orange squares, green diamonds, and gray pentagons indicate IRAC4 sources
NGC 4395-S
622 (top)andclass=1,
NGC 1313-E(bottom).
NGC 4395-S11and
NGCMost1313-E
indicate
the pointings,
while public
the following
clustercandidates
with
2, or 640
3) in
our sample,
Mpc.
of the
LEGUS
galaxies with
YSC
included in our final sample from NGC 7793, NGC 4395, NGC 4449, NGC 1313, and NGC 3738, respectively. In the left panel,
numbers are
sourceof thecentralareclustersin the (Press
correspondingcatalogs
Adamo et(2017)
catalogs.
TheLVLcolorsof thebysourceIDs
indicating
that
the IDs
best-fitmodelsacceptable
areal.also
includedin thesampledesign.
the pink hexagons, blue stars, and yellow downward-pointing triangles are sources excluded by three additional criteria; see theareal.
the2007).
same asThethoseresultingof the corresponding
in theandright panelTheof Figure
The RGB
images areto
composedof theofF275W
etsample has points
a median
IRAC 4.
imagesare resampled
a pixel scale
0.00 75,(blue), F555W (green),and F814W +1.89(red) images. The small
red circles and large cyan circles are the same ast-3099(el-344.2g-344.23r-344.542.0.9(.903llo)2519(WBa(o)-040.9(eo(c)40.9(eob3)-0.9l3r-344.5textv)50.9(atin)-1(g-344.2on3)-0.9lc)40.99(of(Adamome)-742s5(47)-1(re)z0.9(me)-842s50)-11)-1
0.9left(me)-342s5(l,F)-0.8(ig)-0.9the.9724distribution clearly.
+0.49
200 Figure ou33-0.900
for moredetails.
The uncertainties
in 01
age
are
not plottedin(re)!
thepanel01.184G0.9
in orderto show

68% scatter of 0.31 0.28 and 1.20 0.87 for Q and red ,respectively. See Adamo et al. (2017) for more details.
3. IMAGE PROCESSING AND PHOTOMETRY

and the FWHMs of their PSFs are ⇠ 1. 6 and ⇠ 1. 9 for
IRAC1 and IRAC4, respectively (Dale et al. 2009).
The IRAC1 images are convolved to match the PSF
of IRAC4 images using the Photutils package (BradleyM⇤

for the aperture. To decrease the mismatch between
the age of the clusters and the age of the population

The exclusion of the two outliers gives a slightly stronger
correlation with r = 0.84 ± 0.05. Thus, older popula-https://arxiv.org/pdf/1912.10012#38

Magnetic fields in the early stages of massive star formation as revealed by ALMA
Junhao Liu1,2,3 , Qizhou Zhang3 , Keping Qiu1,2 , Hauyu Baobab Liu4 , Thushara Pillai5,6 et al.
School of Astronomy and Space Science, Nanjing University, 163 Xianlin Avenue, Nanjing, Jiangsu 210023, People’s
Republic of China; 2 Key Laboratory of Modern Astronomy and Astrophysics (Nanjing University), Ministry of
Education, Nanjing, Jiangsu 210023, People’s Republic of China; 3 Center for Astrophysics | Harvard & Smithsonian,
60 Garden Street, Cambridge, MA 02138, USA; 4 Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box
23-141, Taipei 10617, Taiwan; 5 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany;
6
Institute for Astrophysical Research, 725 Commonwealth Ave, Boston University Boston, MA 02215, USA
1E-mail contact: kpqiu at nju.edu.cn

The Magnetic Field in G28.34

7

Figure 2. Dust continuum and magnetic orientation maps. The Stokes I of the ALMA 1.3 mm continuum is shown in grayscales and in contours. The ALMA 1.3 mm continuum of three clumps are also shown in contours. The contour levels are (±3,6, 10, 20, 30, 40, 50, 70, 90, 110, 150, 180, 210, 250, 290, 340, 390, 450) ⇥ I , where I is the rms noise of the Stokes I maps
(see Section 2). Line segments represent the orientation of the plane-of-sky magnetic field and have arbitrary length. Red andcyan line segments correspond to data with P I/ P I > 2 and P I/ P I > 3, respectively. Blue dashed regions indicate the areasof MM1-Core1 and MM4-Core4 in which the polarization data are used for the angular dispersion function analysis. The sizeof the synthesized beam are indicated in the lower left corner of each panel.

We present 1.3 mm ALMA dust polarization observations at a resolution of ∼0.02 pc of three massive molecular clumps,
MM1, MM4, and MM9, in the infrared dark cloud G28.34+0.06. With the sensitive and high-resolution continuum
data, MM1 is resolved into a cluster of condensations. The magnetic field structure in each clump is revealed by thepolarized emission. We found a trend of decreasing polarized emission fraction with increasing Stokes I intensities
in MM1 and MM4. Using the angular dispersion function method (a modified Davis-Chandrasekhar-Fermi method),the plane-of-sky magnetic field strength in two massive dense cores, MM1-Core1 and MM4-Core4, are estimated tobe ∼1.6 mG and ∼0.32 mG, respectively. The virial parameters in MM1-Core1 and MM4-Core4 are calculated tobe ∼0.76 and ∼0.37, respectively, suggesting that massive star formation does not start in equilibrium. Using thepolarization-intensity gradient-local gravity method, we found that the local gravity is closely aligned with intensity
gradient in the three clumps, and the magnetic field tends to be aligned with the local gravity in MM1 and MM4
except for regions near the emission peak, which suggests that the gravity plays a dominant role in regulating the gas
collapse. Half of the outflows in MM4 and MM9 are found to be aligned within 10◦ of the condensation-scale (<0.05pc) magnetic field, indicating that the magnetic field could play an important role from condensation to disk scale in
the early stage of massive star formation. We found that the fragmentation in MM1-Core1 cannot be solely explained
25by thermal Jeans fragmentation or turbulent Jeans fragmentation.
Accepted by ApJ
http://arxiv.org/pdf/2005.01705

VLT/SPHERE survey for exoplanets around young, early-type stars including systems
with multi-belt architectures
大質量星形成領域G28.34+0.06のM. Lombart1,2 , G. Chauvin1,3 , P. Rojo4 , E. Lagadec5 P. Delorme3 et al.

ALMAダスト偏光観測による磁場構造議論

Unidad Mixta Internacional Franco-Chilena de Astronomía, CNRS/INSU UMI 3386 and Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile; 2 Univ Lyon, Univ Lyon1, Ens de Lyon, CNRS, Centre
de Recherche Astrophysique de Lyon UMR5574, F-69230, Saint-Genis,-Laval, France; 3 Univ. Grenoble Alpes, CNRS,
IPAG, F-38000 Grenoble, France; 4 Departamento de Astronomía, Universidad de Chile, Camino El Observatorio
1

5

#39

pc) magnetic field, indicating that the magnetic field could play an important role from condensation to disk scale in
the early stage of massive star formation. We found that the fragmentation in MM1-Core1 cannot be solely explainedby thermal Jeans fragmentation or turbulent Jeans fragmentation.
Accepted by ApJ
http://arxiv.org/pdf/2005.01705

VLT/SPHERE survey for exoplanets around young, early-type stars including systems
with multi-belt architectures
M. Lombart1,2 , G. Chauvin1,3 , P. Rojo4 , E. Lagadec5 P. Delorme3 et al.
Unidad Mixta Internacional Franco-Chilena de Astronomía, CNRS/INSU UMI 3386 and Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile; 2 Univ Lyon, Univ Lyon1, Ens de Lyon, CNRS, Centre
de Recherche Astrophysique de Lyon UMR5574, F-69230, Saint-Genis,-Laval, France; 3 Univ. Grenoble Alpes, CNRS,
IPAG, F-38000 Grenoble, France; 4 Departamento de Astronomía, Universidad de Chile, Camino El Observatorio
1515, Las Condes, Santiago, Chile; 5 Université Cote d’Azur, OCA, CNRS, Lagrange, France
1

E-mail contact: maxime.lombart at ens-lyon.fr

VLT/SHPEREによる
若いearly type starsの惑星・デブリ円盤のイメージングサーベイ

mass accretion rates, while external photoevaporation might explain the low disk masses and high accretion rates. A
partial possible solution to the discrepancy with the viscous models is that the gas-to-dust ratio of the disks at ∼5-10
Myr is significantly different and higher than the canonical 100, as suggested by some dust and gas disk evolution
models. The results shown here require the presence of several inter-playing processes, such as detailed dust evolution,
external photoevaporation and possibly MHD winds, to explain the secular evolution of protoplanetary disks.

#41

Accepted by Astronomy & Astrophysics
https://arxiv.org/pdf/2004.14232

Maria Jose Maureira1 , Jaime E. Pineda1 , Dominique, M. Segura-Cox1 , Paola Caselli1 , Leonardo Testi2 ,Giuseppe Lodato3 , Laurent Loinard4,5 and Antonio Hernandez-Gomez4,6,7
Max-Planck-Institut für extraterrestrische Physik (MPE), , Giessenbachstr. 1, D-85741 Garching, Germany; 2European Southern Observatory, Karl-Schwarzschild-Strasse 2, 85748 Garching, Germany; 3 Dipartimento di Fisica,
Universita degli Studi di Milano, Via Celoria 16, I-20133 Milano, Italy; 4 Instituto de Radioastronomia y Astrofisica,
Universidad Nacional Autónoma de México Morelia, 58089, Mexico; 5 Instituto de Astronomia, Universidad Nacional
Autonoma de Mexico,
Apartado Postal 70-264, Ciudad de Mexico 04510, Mexico; 6 Max-Planck-Institute für Radioastronomie Auf dem
Hügel 69, 53121 Bonn, Germany; 7 IRAP, Université de Toulouse CNRS, UPS, CNES, Toulouse, France
1

E-mail contact: maureira at mpe.mpg.de
We present 3 mm ALMA continuum and line observations at resolutions of 6.5 au and 13 au respectively, toward theClass 0 system IRAS 16293-2422 A. The continuum observations reveal two compact sources towards IRAS 16293-2422
A, coinciding with compact ionized gas emission previously observed at radio wavelengths (A1 and A2), confirming thelong-known radio sources as protostellar. The emission towards A2 is resolved and traces a dust disk with a FWHM
size of ∼12 au, while the emission towards A1 sets a limit to the FWHM size of the dust disk of ∼4 au. We also detectspatially resolved molecular kinematic tracers near the protostellar disks. Several lines of the J = 5 − 4 rotationaltransition of HNCO, NH2 CHO and t-HCOOH are detected, with which we derived individual line-of-sight velocities.
Using these together with the CS (J = 2 − 1), we fit Keplerian profiles towards the individual compact sources andderive masses of the central protostars. The kinematic analysis indicates that A1 and A2 are a bound binary system.
Using this new context for the previous 30 years of VLA observations, we fit orbital parameters to the relative motion
between A1 and A2 and find the combined protostellar mass derived from the orbit is consistent with the masses
derived from the gas kinematics. Both estimations indicate masses consistently higher (0.5 < M1 < M2 < 2 M" )than previous estimations using lower resolution observations of the gas kinematics. The ALMA high-resolution data
provides a unique insight into the gas kinematics and masses of a young deeply embedded bound binary system.
Accepted by ApJ
27

What did the seahorse swallow? APEX 170 GHz observations of the chemical conditionsin the Seahorse infrared dark cloud
Oskari Miettinen11

近接連星系をなしているとみられる162932422 Aの多波長観測。固有運動から中心星質量

Orbital and mass constraints of the young binary system IRAS 16293-2422 A

https://arxiv.org/pdf/2005.11954

IRAS16293-2422 A-B原始星システムのうち，

Academy of Finland, Hakaniemenranta 6, P.O. Box 131, FI-00531 Helsinki, FinlandE-mail contact: oskari.miettinen at aka.fi
Context. Infrared dark clouds (IRDCs) are useful target sources for the studies of molecular cloud substructure
evolution and early stages of star formation. Determining the chemical composition of IRDCs helps to constrain theを議論。分子輝線によるガス運動も議論。
4

Maureira et al.

Sanhueza et al. 2012; Vasyunina et al. 2014; Gerner et al. 2014;
Zeng et al. 2017; Colzi et al. 2018; Taniguchi et al. 2019).
Observations of IRDCs have shown that they often appear to be filamentary structures, which could be the results of converging atomic gas flows or collisions between precursor interstellar clouds that led to the formation of the IRDC filament (e.g. Jiménez-Serra et al. 2010;cf. Heitsch et al. 2006). Moreover, the IRDC filaments appear
to be fragmented along their longest axis, potentially via theso-called sausage-type fluid instability in cylindrical geometry, and the substructures thus formed, known as clumps, arefound to be hierarchically fragmented down to smaller units,
the dense cores (e.g. Jackson et al. 2010; Wang et al. 2011;
Beuther et al. 2015; Mattern et al. 2018; Svoboda et al. 2019).
Some of the dense substructures of IRDCs are found tobe associated with early stages of high-mass star formation (e.g. Rathborne et al. 2006; Beuther & Steinacker 2007;
Chambers et al. 2009; Battersby et al. 2010). This makes IRDCsvery useful targets for the star formation research because theprocess(es) of high-mass star formation still remains inconclusive (e.g. Motte et al. 2018 for a review; Padoan et al. 2020).
In this paper, we present a study of the chemical properties ofthe filamentary IRDC G304.74+01.32. This IRDC has been thetarget of submillimetre and millimetre dust continuum studies(Beltrán et al. 2006; Miettinen & Harju 2010; Miettinen 2018),
and molecular spectral line observations (Miettinen 2012). Most
recently, Miettinen (2018) found that 36% of the clumps
along G304.74+01.32 are fragmented into smaller cores, andthat 65% of those cores are associated with IR sources detected with the Wide-field Infrared Survey Explorer (WISE;
Wright et al. 2010). Those results strongly point to a hierarchicalfragmentation of the filament. Miettinen (2018) also found that
G304.74+01.32 has a seahorse-like morphology in the Herschel1(Pilbratt et al. 2010) far-IR and submillimetre images (Fig. 1;
cf. Fig. 2 in Miettinen 2018) and, therefore, proposed the nickname Seahorse Nebula for the cloud. Moreover, there appear
to be elongated, dusty striations that are roughly perpendicularly projected with respect to the main Seahorse filament. The
striations are also seen in absorption in the WISE mid-IR images (Fig. 2; cf. Fig. 3 in Miettinen 2018). Such features could
be an indication that the filament is still accreting gas from its
surrounding medium, but this hypothesis needs to be tested via
spectral line imaging observations.
Only one of the clumps in the Seahorse IRDC appears tobe associated with high-mass star formation, namely the one associated with the IR source IRAS 13039-6108. Sánchez-Mongeet al. (2013) found that IRAS 13039-6108 shows radio continuum emission characteristic of an optically thin H ii region. The
IRAS 13039-6108 region is also surrounded by extended anddiffuse IR emission, most notably seen in the WISE 12 µm imd
i#42

IRDC G304.74+01.32の

ガス化学に関する包括的な仕事。APEXによる。

分子間のアバンダンスの相関解析から，[C2H]の低下などが良い進化の指標になりそう。

Accepted by ApJ
https://arxiv.org/pdf/2005.11954

What did the seahorse swallow? APEX 170 GHz observations of the chemical conditionsin the Seahorse infrared dark cloud
Oskari Miettinen1
1

Academy of Finland, Hakaniemenranta 6, P.O. Box 131, FI-00531 Helsinki, FinlandE-mail contact: oskari.miettinen at aka.fi
Context. Infrared dark clouds (IRDCs) are useful target sources for the studies of molecular cloud substructure
evolution and early stages of star formation. Determining the chemical composition of IRDCs helps to constrain theinitial conditions and timescales (via chemical clocks) of star formation in these often filamentary, dense interstellar
clouds.
Aims. We aim to determine the fractional abundances of multiple different molecular species in the filamentary IRDC
G304.74+01.32, nicknamed the Seahorse IRDC, and to search for relationships between the abundances and potential
evolutionary trends.
Methods. We used the Atacama Pathfinder EXperiment (APEX) telescope to observe spectral lines occurring at about
170 GHz frequency towards 14 positions along the full extent of the Seahorse filament. The sample is composed of fiveclumps that appear dark in the mid-IR, eight clumps that are associated with mid-IR sources, and one clump thatis already hosting an HII region and is, hence, likely to be in the most advanced stage of evolution of all the target
sources. We also employed our previous 870 µm dust continuum imaging data of the Seahorse.
Results. Six spectral line transitions were detected (≥ 3σ) altogether, namely, SO(NJ = 44 − 33 ), H13 CN(J = 2 − 1),
H13 CO+ (J = 2 − 1), SiO(J = 4 − 3), HN13 C(J = 2 − 1), and C2 H(N = 2 − 1). While SO, H13 CO+ , and HN13 C
were detected in every source, the detection rates for C2 H and H13 CN were 92.9% and 85.7%, respectively. Only one
source (SMM 3) showed detectable SiO emission (7.1% detection rate). Three clumps (SMM 5, 6, and 7) showed theSO, H13 CN, H13 CO+ , HN13 C, and C2 H lines in absorption. Of the detected species, C2 H was found to be the mostabundant one with respect to H2 (a few times 10−9 on average), while HN13 C was found to be the least abundantspecies (a few times 10−11 ). We found three positive correlations among the derived molecular abundances, of whichthose between C2 H and HN13 C and HN13 C and H13 CO+ are the most significant (correlation coefficient r # 0.9).
The statistically most significant evolutionary trends we uncovered are the drops in the C2 H abundance and in the[HN13 C]/[H13 CN] ratio as the clump evolves from an IR dark stage to an IR bright stage and then to an HII region.
Conclusions. The absorption lines detected towards SMM 6 and SMM 7 could arise from continuum radiation froman embedded young stellar object and an extragalactic object seen along the line of sight. However, the cause ofabsorption lines in the IR dark clump SMM 5 remains unclear. The correlations we found between the different
molecular abundances can be understood as arising from the gas-phase electron (ionisation degree) and atomic carbon
abundances. With the exception of H13 CN and H13 CO+ , the fractional abundances of the detected molecules in the
Seahorse are relatively low compared to those in other IRDC sources. The [C2 H] evolutionary indicator we foundis in agreement with previous studies, and can be explained by the conversion of C2 H to other species (e.g. CO)
when the clump temperature rises, especially after the ignition of a hot molecular core in the clump. The decrease of[HN13 C]/[H13 CN] as the clump evolves is also likely to reflect the increase in the clump temperature, which leads toan enhanced formation of HCN and its 13 C isotopologue. Both single-dish and high-resolution interferometric imaging
of molecular line emission (or absorption) of the Seahorse filament are required to understand the large-scale spatial
distribution of the gas and to search for possible hot, high-mass star-forming cores in the cloud.
Accepted by Astronomy & Astrophysics
https://arxiv.org/pdf/2005.08250.pdf

The big sibling of AU Mic: a cold dust-rich debris disk around CP-72 2713 in the β Pic