星形成ゼミ#339 1-8 担当：片岡1. Particle acceleration and magnetic field amplification in massive young stellar object jets
Anabella Araudo, Marco Padovani, Alexandre Marcowith
2. A VLA Census of the Galactic H II Region Population
William P. Armentrout, Loren D. Anderson, Trey V. Wenger, Dana S. Balser, Thomas M. Bania ⋆3. ALMA-IRDC: Dense gas mass distribution from cloud to core scales
A. T. Barnes, J. D. Henshaw, F. Fontani, J. E. Pineda, G. Cosentino, J. C. Tan, P. Caselli, I. Jiménez-Serra, C.
Y. Law, A. Avison, F. Bigiel, S. Feng, S. Kong, S. N. Longmore, L. Moser, R. J. Parker, Á. Sánchez-Monge,
K. Wang
4. Fragmentation in the massive G31.41+0.31 protocluster
M. T. Beltrán, V. M. Rivilla, R. Cesaroni, L. T. Maud, D. Galli, L. Moscadelli, A. Lorenzani, A. Ahmadi, H.
Beuther, T. Csengeri, S. Etoka, C. Goddi, P. D. Klaassen, R. Kuiper, M. S. N. Kumar, T. Peters, Á. SánchezMonge, P. Schilke, F. van der Tak, S. Vig, H. Zinnecker
5. Fitting an Analytic Magnetic Field to a Prestellar Core
Gianfranco Bino, Shantanu Basu
6. Self-scattering on large, porous grains in protoplanetary disks with dust settling
Robert Brunngräber, Sebastian Wolf
7. Magnetic fields and Star Formation around HII regions: The S235 complex
R. Devaraj, D. P. Clemens, L. K. Dewangan, A. Luna, T. P. Ray, J. Mackey
8. Water in star-forming regions (WISH): Physics and chemistry from clouds to disks as probed by
Herschel spectroscopy
E. F. van Dishoeck, L. E. Kristensen, the WISH teamSelf-scattering on large, porous grains in protoplanetary disks withdust settling• 原始惑星系円盤の偏光は散乱による事がわかっている。
ダストの空隙率の影響を調べた。
•
Fig. 7. Map of the polarisation degree with superimposed polarisation
vectors for ourreference disk with = 1 (left), 10 1 (centre), and 10 2 fo
• 空隙がある方が偏光度が上がる場(right), and an inclination of 8 = 0° (top) and 45° (bottom). The lengths of an
the polarisation合があることがわかった。
vectors are scaled to the polarisation degree, as indicated poin the top left corner.
0.00

y [au]

200

0.25

0.50

Degree of polarisation [%]
0.75

1.00

1.25

1.50

1.3 %

0
200

=1

= 10 1= 10 2

1.75

2.00

m
de
ag
du
th
20
di
of
fa

th
200
po
to
0
in
2.
(A
200
sp
200
0
200
200
0
200
200
0
200
an
x [au]x [au]x [au]
thFig. 8. Same as Fig. 7, but for a disk with a ten times higher dust mass. no
The lengths of the polarisation vectors are scaled to the polarisation isdegree, as indicated in the top left corner.
be
ar
tim
polarisation degree for inclined disks through the optical depth.
The optical depth and the mass under the assumption of a known

y [au]

Robert Brunngräber, Sebastian Wolf
Observations of protoplanetary disks (PPDs) at sub-mm wavelengths suggestthat polarisation is caused by scattering of thermal radiation. Most of the dust
models that are used to explain these observations have major drawbacks:
They either use much smaller grains than expected from dust evolution
models, or result in lower polarisations than observed. We investigate theeffect of dust grain porosity on the polarisation due to scattering at sub-mm
wavelengths arising from grains of up to mm sizes, as they are expected to bepresent in the midplane of PPDs. Using the effective medium theory, we
calculated the optical properties of porous dust and used them to predict the
behaviour of the scattering polarisation. Subsequently, radiative transfer
simulations for PPDs with porous dust were performed to analyse theadditional effect of the optical depth structure, and thus the effect of multiple
scattering events and inhomogeneous temperature distributions on thepolarisation. We find that porous dust with moderate filling factors of about
10\% increase the polarisation compared to compact grains. For higher grain
porosities, that is, grains with a filling factor of 1\% or lower, the extinction
opacity decreases, as does the optical depth of a disk with constant mass.
Consequently, the unpolarised direct radiation dominates the scattered flux,
and the polarisation drops rapidly. Even though the simulated polarisation ishigher than in the case of compact grains, it is still below the typical observed
values for face-on disks. However, the polarisation can be increased when
crucial model assumptions derived from disk and dust evolution theories, forinstance, dust settling and mm-sized dust grains, are dropped. In the case of
inclined disks, however, our reference model is able to achieve polarisation
degrees of about 1\%, and using higher disk masses, even higher than this.
arXiv:2103.04819 [astro-ph.SR]

Se
in
co
60
th
fa

For the Stokes V images (Figure 2), 3.1 mm appears to be
purely noise. For both 1.3 mm and 870 μm, StokesV has anegative ﬂux blob offset to the top-right of the Stokes I
intensity, and this blob is signiﬁcantly brighter in 870 μm.
Although circular polarization can in principle be produced byscattering of linearly polarized light off non-spherical grains(e.g., Tazaki et al. 2017), StokesV has not been wellcharacterized for ALMA, so its detection could be spurious
and will not be discussed in detail in this paper.
The central polarization vectors for 1.3 mm and 870 μm are
55°. 4 ±0°. 96 and 55°. 3 ±0°. 41, respectively. The major axisof HL Tau’s disk is 138°. 02 ± 0°. 07 (ALMA Partnership et al.
2015), indicating the polarization vectors are ∼7° from being
perpendicular to the major axis. We would expect thepolarization to be perfectly perpendicular to the major axis.
The reason for this small 7° discrepancy is uncertain, but could
indicate asymmetry in the disk.
The right panels of Figure 2 show P for each wavelength.
For 1.3 mm and 870 μm, there is a slight asymmetry for P
along the major axis of the disk, with the southeast morepolarized than the northwest. For 1.3 mm, the absolutedifference between the two P peaks along the major axis is0.0013 ± 0.0014, which is not statistically signiﬁcant. The
asymmetry for 870 μm is more apparent, although the southeast
part of the disk does not peak (i.e., it is still increasing at theedge of the detected polarized emission). The P peak toward
the northwest is 0.0077 ± 0.0010. Drawing a line from thenorthwest peak through the center of the disk, at the sameradius on the southeast side, P=0.0094 ± 0.0010. P rises
to ∼0.0115 at the edge of the disk. Therefore, the P asymmetry
along the major axis for 870 μm is ∼0.002–0.004. These
asymmetries are relatively moderate, but the fact that they
appear both at 1.3 mm and 870 μm suggests they could
be a real feature. If we consider total polarization PT 
Q 2 U 2 V 2 I , the asymmetry is similar; for 1.3 mm, it is0.0012 ± 0.0014, while for 870 μm it is ∼0.001–0.004.
Yang et al. (2017) showed that asymmetry along the major
axis is not expected for scattering-induced polarization in an
axisymmetric disk, so axisymmetry must be broken if thepolarization is really due to scattering even though the total
intensity (Stokes I ) is highly axisymmetric. On the other hand,the polarization along the minor axis is expected to beasymmetric if the dust disk is optically thick and has asigniﬁcant angular width (for a cartoon illustration, see Figure 6
in Yang et al. 2017). HL Tau disk is known to be opticallywhat can we learn from self-scattering?
Albedo = efficiency of scattering

0.1
n size [cm]

-0.2
1

-0.2
0.001

-0.2

what can we learn from self-scattering?
0.01
0.1
Maximum grain size [cm]1

Anisotropies ofcattering angle of 90 and the albedo ! = sca /(abs + sca ) as a function of maximum grain size.
radiation
fi
eld
at
the
Polarization
Grain
size
=
3.5
ve power law of n(a) / a. The wavelengths are assumed to be 1.3 mm for the left panel andindicatethemaximum
grain
size
which
has
the
most
efficientpolarizationby
90
scattering.
observed wavelengths.
1.2

for observed

1grain size amax

r each wavelengthpolarization due to 90
gnificant only when theve a large albedo (§2.2)

0.8

Pω

mm
0 µm
0 µm
0 µm

Prediction of polarization fraction
3.1 mm0.87 mm (Band 3)(Band 7)7 mm
(Band 1)0.34 mm(Band 10)

0.6
0.4
0.2
0
0.001

0.01
0.1
Maximum grain size [cm]1
Kataoka et al., 2015

(grain size) ~ λ/2π

R. Brunngräber and S. Wolf: Self-scattering on porous grains
100

smax = 100 µm
smax = 1 mm
smax = 10 mm

Parameterp Ŝ11 ( = 90°)

10 1

Table 1. Disk and stellar parameter values of our reference case.

10 2

10 3

10 4 4
10

10 3

10 2

Filling factor

=1

10 1

P

100

Fig. 2. Product of albedo l, single-scattering polarisation ?, and probability density function (ˆ11 for the scattering angle \ = 90°. All quantities
are calculated for a grain size distribution with @ = 3.5, Bmin = 5 nm
and three different maximum grain sizes Bmax = 100 µm (blue), 1 mm(red), and 10 mm (green) at a wavelength of 850 µm for astrosil; see
Sect. 2.3 for the corresponding discussion.

with the vertically integrated surface density

Variable

ValuesInner radius'in [au]
0.1
Outer radius'out [au]300• アルベドとpとS11をかけたもの
Reference radius'0 [au]100
がだいたいの偏光の効率を表して
Truncation radius'trunc [au]100いる。
Reference scale height ⌘ref [au]10
• 1mmとか 1cmダストでは、充Density profileW
1.1
Flaring parameterV1.1
填率10%くらいだと偏光度が高い
Dust mass"dust [M ] 10 4
Settling parameter
Usettle10 2
Inclination
8 [°]0Stellar luminosity!¢ [L ]
0.9
Stellar temperature)¢ [K]
4050
Silicate: 62.5 %Dust composition
Graphite? : 25 %(mass fraction)
Graphite k : 12.5 %Silicate:
3800
rbulk
Grain bulk density
Graphite: 2200
3[kg m ]
Mixture: ⇠ 2986
Minimum grain size
Bmin [nm]5
Maximum grain size
Bmax [mm]1
Size distribution@
3.5

⇠sca + ⇠abs

A&A proofs manuscript no. selfscattering_porous_protoplanetary_disks

10 2

0.2

10 3

0.8

P

p=

0.4

=1
Filling factor

=1

0.6

1.0

10 1

S12/S11

P

0.8

10 1

Filling factor

100

1.0

0.6

Albedo

100

10 2

0.4
10 3

0.2

0.0

100

P

16

=1

8

Filling factor

100

101

Maximum grain size smax [mm]

102

0.0

100

32

10 1

10 1

10 1
10 1

4

10 2

2

10 2
10 2

Ŝ11( = 90°)

Maximum grain size smax [mm]

102

P

101

=1

100

Filling factor

10 1

10 4

2
1
ext [cm g ]10 4

The albedo and extinction opacity, which determines thdepth, are shown in the upper right and lower left
Fig. 1, respectively. For large maximum grain sizes, the
左上：１回散乱の偏光度。1だとfirst increases •withdecreasing filling factor and decreas
after reaching a偏光が100%。この数字が高いほ
maximum value. Whereas the grain
proportional to thefilling factor, the extinction cross-sectio
ど偏光が受かりやすいa more complex behaviour. At first, it decreases much slo• 右上：アルベド。散乱の効率。こdecreasing filling
factor than the grain mass. However, ffactors . 10 2の数字が高いほど散乱しやすい。
, the cross-section is roughly proportionand thus decreases
faster than the grainmass. Hence, theopacity。そ• 左下：extinction
is maximised at the transition of these two regions. In
もそもの天体の明るさ。
the opacity is constant for a constant product Bmax · , w• 右下：90度散乱の効率。この数
been shown previously
by Kataoka et al. (2014) and Taz(2019), for example.
字が低いと前方散乱→あんまりう
Scattering angle dependency: If a scattering even
れしくないthe actual intensity of the radiation scattered in the dof the observer is also a function of the scattering ang
single scattering, this dependency is given by the firs
matrix element (11 (\). The probability density functionØ c (11 ( \)for a scattering angle of \ = 90° is show0

(11 sin \ d\lower right panel of Fig. 1. For large maximum grain si
10 3
10 3
probability decreases for increasing grain size because lar
1tend to scatter predominantly in the forward direction
4100.5\ ⇡ 0°. With increasing grain size, the intensity of the r
4
4
10
10
1
0
1
2
10
10
10
10
10 1
100
101
102
that is scattered at 90° decreases. As a consequence, the fl
Maximum grain size smax [mm]Maximum grain size smax [mm]of scattered to direct radiation and hence the degree of pol
Fig. 1. Degree of polarisation for single scattering ? = ( /( for a scattering angle of \ = 90° (upper left), albedo l (upper right), extinction decreases as well.
12

11

10 3

opacity (lower left), and probability density function (ˆ11 for the scattering angle \ = 90° (lower right). All quantities are calculated for a grain size
distribution with @ = 3.5, Bmin = 5 nm and different upper grain size limits Bmax at a wavelength of 850 µm for astrosil.

= 850 µm
Scattering only

Smaller smax

= 850 µm
Scattering only

= 850 µm
Scattering only

= 850 µm
Scattering only

60

50

50

40

40

30

30

20

20

10

10

0
101

0
101

100

100

10 1

10 2

10 2

10 3

10 3

0.5

0.5

0.4

0.4

0.3

0.3

0.2

0.2

0.1

0.1

0.0

50

100

r [au]

150

200

50

100

r [au]

150

200

50

100

r [au]

150

200

50

100

r [au]

150

0.4

0.6

0.8

1.0

1.2

200

0.8 %

0
200

= 10 1=1= 10 2

200
0
200
200

0

x [au]

200

200

0

x [au]

200and disk evolution models, however, nor do these models predict

2000x [au]

200

Fig. 7. Map of the polarisation degree with superimposed polarisation
vectors for our reference disk with = 1 (left), 10 1 (centre), and 10 2(right), and an inclination of 8 = 0° (top) and 45° (bottom). The lengths ofthe polarisation vectors are scaled to the polarisation degree, as indicatedin the top left corner.

0.0
200

Fig. 3. Radial profiles of different quantities for disks with different filling factors (colours). The reference disk model (solid lines) is shown in acomparison with different disk models (dashed lines): A disk with a dust mass "dust = 10 5 M lower by a factor ten (first column), a disk with adust mass "dust = 10 3 M higher by a factor ten (second column), a disk with a maximum grain size Bmax = 100 µm smaller by a factor ten (thirdcolumn), and a disk with a maximum grain size Bmax = 10 mm larger by a factor ten (fourth column). Top: Degree of polarisation for scattering only,
i.e. without the unpolarised direct re-emission. Middle: Flux ratio of scattered to direct re-emission. Bottom: Net degree of polarisation of the total
flux, i.e. scattered and direct radiation, including vertical error bars.

to a reduced overall degree of polarisation (lower left panel of

0.2

Larger smax

y [au]

60

Larger Mdust

0.0

10 1

Degree of polarisation [%]Scattered / direct flux

Degree of polarisation [%]Smaller Mdust= 10 2y [au]= 10 1=1Degree of polarisation [%]• 偏光度は高くても0.5%程度• スペクトル指数は合わないALMA-IRDC: Dense gas mass distribution from cloud to core scales
A. T. Barnes, J. D. Henshaw, F. Fontani, J. E. Pineda, G. Cosentino, J. C. Tan, P.
Caselli, I. Jiménez-Serra, C. Y. Law, A. Avison, F. Bigiel, S. Feng, S. Kong, S. N.
Longmore, L. Moser, R. J. Parker, Á. Sánchez-Monge, K. Wang
Infrared dark clouds (IRDCs) are potential hosts of the elusive early phases of highmass star formation (HMSF). Here we conduct an in-depth analysis of thefragmentation properties of a sample of 10 IRDCs, which have been highlighted assome of the best candidates to study HMSF within the Milky Way. To do so, we haveobtained a set of large mosaics covering these IRDCs with ALMA at band 3 (or 3mm).
These observations have a high angular resolution (~3arcsec or ~0.05pc), and highcontinuum and spectral line sensitivity (~0.15mJy/beam and ~0.2K per 0.1km/s
channel at the N2H+(1-0) transition). From the dust continuum emission, we identify
96 cores ranging from low- to high-mass (M = 3.4 to 50.9Msun) that are
gravitationally bound (alpha_vir = 0.3 to 1.3) and which would require magnetic fieldstrengths of B = 0.3 to 1.0mG to be in virial equilibrium. We combine these resultswith a homogenised catalogue of literature cores to recover the hierarchicalstructure within these clouds over four orders of magnitude in spatial scale (0.01pc
to 10pc). Using supplementary observations at an even higher angular resolution, we
find that the smallest fragments (<0.02pc) within this hierarchy do not currently
have the mass and/or the density required to form high-mass stars. Nonetheless,the new ALMA observations presented in this paper have facilitated the identificationof 19 (6 quiescent and 13 star-forming) cores that retain >16Msun without further
fragmentation. These high-mass cores contain trans-sonic non-thermal motions,
are kinematically sub-virial, and require moderate magnetic field strengths forsupport against collapse. The identification of these potential sites of high-mass starformation represents a key step in allowing us to test the predictions from highmass star and cluster formation theories.
arXiv:2103.09122 [astro-ph.GA]• 10個の赤外線暗黒星雲IRDCについて、ALMA 3 mm波で観測した。

• 空間分解能 3 arcsec = 0.05
pc
• ダスト連続波からは、96個のコアを見つけた。質量は3.450.9Msun• alpha̲vir = 0.3-1.3 (重力的に束縛されている)• ビリアル平衡に達するためには、
B=0.3-1.0 mGくらい必要• 0.02 pc以下のフラグメントは、その質量が大質量星を作れないほ
ど小さかった。
• だが、それでも19のコアが（16Msun以上）同定された

Dec (J2000)
12.70
12.68

G034.43+00.24

G035.39-00.33

Galactic Longitude

Cloud F

Cloud H

G034.77-00.55

Galactic Longitude

35.5
Galactic Longitude

Cloud G

0.09 0.05

T (K)
18
19

20

3.62
3.28

1.3
4.7

20.0
19.8

0.33

0.1

2.1

4.2

6.3

0.1

0.3

0.5

200

375

703

1318 17

0.28

0.48

0.1

1.1

2.1

3.1

0.1

0.3

0.5 190.6

343.5

618.9 1115.4

18

19

276.48276.46

Dec (J2000)
04.06
04.04

34.5

280.72

34.0

3.6
3.3

13.5
12.3

20
0

280.70

Cloud D

Galactic Longitude

04.02

00.8

4460.03300.00.19

1pc

Table 1. Table of the global properties across the cloud sample. Shown in columns are the cloud names from Butler & Tan (2012), IDs from Rathborne
et al. (2006), Cloudthe systemic velocity (Eradiideterminedfrom nearand mid-M
infraredextinction
IDsys ), cloud kinematic
vsys distances⇡cl(⇡cl ),'effective"cl('eff,cl ),fmasses
Uvir,cl)dust,clfNT,cl=ceff,cl
cl
NT,cl
13
1
1
1
mapping ("cl ), velocity“G(longitude)(latitude)”dispersions from CO (1-0)km
emission(fpc
parametersM(Uvir,cl ),kmhavecl ), virial pcss all been- taken fromK Kainulainen
km s & Tan -(2013, Table# 1).
Also given is the mean Herschel derived dust temperature measured over the ALMA coverage ()dust,cl ), the non-thermal velocity dispersion and sonic MachCloud
A MNT,cl ;G018.82-00.28
59-69of cores
4800
10.4
18500.0continuum
2.04 observations
1.4
18.3 3.1).2.0
7.9
8number (fNT,clandsection 3.3), and the number
identified
in the ALMA
(section
Cloud B
G019.27+00.07
22-32
2400
2.71
2200.0
1.6
2.2
17.8
1.6
6.3
10
Cloud C
G028.37+00.07
73-83
5000
15.4
53200.0
3.72
2.4
17.7
3.7
14.8
16
Cloud
ID
vsys
⇡cl
'eff,cl"cl
fcl
Uvir,cl)dust,clfNT,cl
MNT,cl=c
Cloud D
G028.53-00.25
81-91
5700
16.9
74300.0
1.85
0.5
16.9
1.8
7.5
16
1
1
1
“G(longitude)(latitude)”
km
s
pc
pc
M
km
s
K
km
s
#
Cloud E
G028.67+00.13
75-85
5100
11.5
28700.0
4.32
1.1
19.1
4.3
16.5
4
Cloud
G034.43+00.24
52-62 4800
3700.0 10.43.5 18500.0
4460.0 2.043.62 1.4 1.3 18.320.0 2.0 3.6
Cloud
A F
G018.82-00.28
59-69
7.913.5 8 20
Cloud
G034.77-00.55
35-45 2400
2900 2.713.06 2200.0
3300.0 1.63.28 2.2 4.7 17.819.8 1.6 3.3
Cloud
B G
G019.27+00.07
22-32
6.312.3 10 0
Cloud
G035.39-00.33
38-48 5000
2900 15.49.69 53200.0
16700.0 3.722.03 2.4 0.7 17.719.5 3.7 2.0
Cloud
C H
G028.37+00.07
73-83
14.87.6 16 10
Cloud
G038.95-00.47
38-48 5700
2700 16.93.73 74300.0
2700.0 1.851.65 0.5 1.2 16.918.0 1.8 1.6
Cloud
D I
G028.53-00.25
81-91
7.5 6.4 16 9
Cloud
G053.11+00.05
17-27 5100
1800 11.5
0.755 28700.0
200.0 4.320.96 1.1 1.5 19.119.0 4.3 0.9
Cloud
E J
G028.67+00.13
75-85
16.53.5 4 3
3.5
3.06

276.56

Cloud C

Table1. 1.
of the image
globalofproperties
acrossthewherecloudseveral
sample.
Showndark
in columns
are the cloud& Tan(2012),
from Rathborne
FigureA Tablethree colour
the Galactic
plane
infrared
clouds (IRDCs)
can benames
seen asfrom
dark Butler
extinction
features.
In IDsthis image,red is
(2006),
themsystemic(Esys), cloudfrom
kinematic
distances(⇡cl ), survey
effective
radii et('al.
), masses
determinednearandinfrared extinctioneff,cl
8 et
m,al.
green
is 5.8and bluevelocity
is 4.5 memission
the Spitzer
GLIMPSE(Carey2009).
Labelled
are threefrom
IRDCs
that
aremidinvestigatedwithin13("cl ),F,velocity
dispersions
from COG034.77-00.55,(1-0) emission (f(Uvir,cl), have
all been
Kainulainen& Tan (2013, Table 1).
cl ), virial parametersthismapping
work: Clouds
G and H(or G034.43+00.24,
G035.39-00.33).
The panels
show
zoom-insof takenthese from
IRDCs
for more detail.
Also given is the mean Herschel derived dust temperature measured over the ALMA coverage ()dust,cl ), the non-thermal velocity dispersion and sonic Mach
number (fNT,cl and MNT,cl ; section 3.3), and the number of cores identified in the ALMA continuum observations (section 3.1).

3700.02900

NH2 (1022 cm 2 )118.2 179.3 271.9 412.4

0.12 0.08

16

17

18

19

Dec (J2000)
04.00
03.98

35.0
Galactic Longitude

52-62
35-45

(g cm 2 )
0.1
0.3
0.5

1pc

276.50Galactic Latitude

Galactic Latitude

00.8

276.58G034.77-00.55

G034.43+00.24
G034.77-00.55

WN2 H+ (1 0) (K kms 1 )0.1 1.2 2.3
3.5

1pc

Cloud G

G034.43+00.24

G035.39-00.33

35.5

F3mm (mJy beam 1 )0.11 0.02 0.15 0.28

5

Cloud B

Figure 1. A three colour image of the Galactic plane where several infrared dark clouds (IRDCs) can be seen as dark extinction features. In this image, red is
8 m, green is 5.8 m and blue is35.5
4.5 m emission from the Spitzer GLIMPSE
survey (Carey et al. 2009). 34.5
Labelled are three IRDCs that are investigatedwithin35.0
34.0
this work: Clouds F, G and H (or G034.43+00.24, G034.77-00.55, G035.39-00.33).
The panels show zoom-ins of these IRDCs
for moredetail.
Galactic
Longitude
Galactic Longitude

Cloud F
Cloud G

ALMA-IRDC: dense gas mass distribution

Dec (J2000)
12.08

Galactic Latitude

Galactic Latitude
Galactic Lati

Cloud F

Cloud H
35.5
Galactic Longitude

3

Cloud A

00.4

Galactic Latitude
Galactic+00.0 Latitude+00.4
Galactic Latitude
00.4
+00.0

ALMA-IRDC: dense gas mass distribution

1pc

281.08281.06
RA (J2000)

Figure 2. (first column) Three colour images from the Spitzer GLIMPSE survey (Churchwell et al. 2009). (second) The ALMA 3 mm dust continuum (without
the primary beam correction), and (third) N2 H+ (J,F1 ,F = 1,0,1 ! 0,1,2) integrated intensity maps. (fourth) Combined near- and mid-infrared extinction derived
mass surface density maps (Kainulainen & Tan 2013). (fifth) Far-infrared Herschel derived column density, and (sixth) dust temperature maps (Marsh et al.
2016, 2017). The yellow circles overlaid on each panel show the positions and sizes of 70 m emission point sources (Molinari et al. 2016; Marton et al. 2017).
The coloured contours overlaid on each panel show boundaries of all the “leaves” (or cores) identified using the dendrogram analysis on the dust continuum
maps (see section 3.1). The red and blue contours indicate the cores that have been classified as star-forming and quiescent, respectively (section 3.1). The cyan
contours indicate the cores that have been identified as being quiescent, high-mass and without further fragmentation (see section 4.3.2). Shown in the lower
left of the ALMA observation panels is the beam size, and in the lower left of the three colour image (first panel) is a scale bar adjusted for the distance of eachcloud (see Table 1).

within the dendrogram boundary) of the 96 identified cores can befound in Table A1. We find that around 5 to 15 cores are present
within each cloud, with the exception of Cloud G where no cores
have been identified. This was likely because the mapped regionof Cloud G is, by design, focused on the eastern shocked region
explored by Cosentino et al. (2018, 2019), and not the main dust extinction/continuum feature(s) previously identified within this cloud(e.g. Rathborne et al. 2006; Butler & Tan 2012; Kainulainen & Tan
2013).
Whereas the dendrogram structure itself has been calculatedusing the non-primary beam corrected continuum maps, all fluxes
quoted in this work have been corrected for the primary beam response. Moreover, when determining the total primary beam corrected flux for each core, we consider that these are not isolated
structures, but rather that they are sitting within the complex three-A. T. Barnes et al.
25

20

Number of cores

Number of coresMedianquiescent

20.0

20
15
10

15

10

5

5

Medianstar forming

17.5

Number of cores

8

All sources
Quiescent

15.0
12.5

Star
forming

10.0
7.5
5.0

• 10個の赤外線暗黒星雲IRDCにつ

2.5
0
0.00

0.05

0.10

Reff [pc]

0.15

0

0.0

0.5

1.0

1.5

log(M ) [M ]

2.0

2.5

0.0

0

1

2

3

4

5

Sonic Mach number Ms [km s 1 ]Figure 3. Distribution of the properties determined across the core sample (section 3.3). Shown in panels from left to right is the effective radius ('eff ), the(logarithm) mass (" ), and sonic Mach number (Ms = fNT /fT ) calculated across the core sample (section 3.3). Shown as black, red and blue stepped
profiles are histograms for the whole sample, the star-forming cores and quiescent (non-star forming) cores, respectively. The red and blue vertical dashed lines
show the median values for the star-forming cores and quiescent cores.

To calculate the mass of each core, we use the integrated flux(( a ; see Table A1) following,"=

3 2 ( a 'gd^ a ⌫ a ()dust ),(1)

where 3 is the source distance (see Table 1), 'gd = 141 is thetotal (gas plus dust)-to-(refractory-component-)dust-mass ratio (assuming a typical interstellar composition of H, He, and metals;
Draine 2011),3 ⌫ a ()dust ) is the Planck function for a dust temperature, )dust , at a representative frequency of a = 90.664 GHz, and^ a = ^0 (a/a0 ) V ⇡ 0.175 cm2 g 1 , when assuming a0 = 230 GHz,V = 1.75 (Battersby et al. 2011) and ^ 0 = 0.899 cm2 g 1 (Ossenkopf& Henning 1994 result for an MRN size distribution with thin ice
mantles after 105 yr of coagulation at a density of 106 cm 3 ).
To obtain an estimate of the dust temperature towards eachof the cores, we make use of spectral energy distribution (SED)
fits to the far-infrared dust continuum observed with Herschel (HiGal; Molinari et al. 2016). We use the mean line-of-sight resultsfrom the PPMAP project shown in Figure 2 (Marsh et al. 2016,2017). To check the fidelity of the PPMAP column density and temperature maps, where possible, we compare to the corresponding
maps from Nguyen Luong et al. (2011), Lim et al. (2016), Zhang
et al. (2017), and Soam et al. (2019). These authors independently
produced column density and temperature maps following a moreconventional far-infrared SED fitting routine, whilst accounting for

We find dust median temperatures of )dust = 17.8+1.4
K.
0.9
These dust temperatures for each core are then used to calculate
a median mass of " = 18.4+32.5
M across the core sample
15.0
4 We assumed a typically ⇠ 10 per cent in(" b = 4.6+16.8
M).
3.7
the absolute flux scale of the ALMA observations, and, following
Sanhueza et al. (2017), we assume an uncertainty of ⇠ 30 per centdust opacity. These uncertainties in the dust opacity, dust emission
fluxes, and the distance propagate to give an uncertainty of ⇠ 50 per
cent in masses. The histogram distribution of the size and mass ofthe star-forming and quiescent core samples is shown in Figure 3.
It is worth noting that the dust temperatures measured here are
averages sampled by the beam of the PPMAP maps (⇠12 00 ; shown
as a black circle in the rightmost panel of Figure 2), and the beam of
the ALMA observations used in this work is significantly smaller(⇠ 3 00 ; see Table 2). Therefore, smaller-scale temperature variations
due to e.g. cold cores or embedded stellar objects could cause us
to over- or under-estimate the temperature, respectively (e.g. see
Ragan et al. 2011; Dirienzo et al. 2015; Sokolov et al. 2017).
To quantify the extent that the cores containing embedded
protostars may have underestimated temperatures in the PPMAP
maps, we follow a procedure outlined in Peretto et al. (2020) for
determining the temperature at higher angular resolution (⇠5 00 ). For
this estimate, these authors require that the cores have an identified
70 m point source (the 70 m flux given in Table A1). The 70 m
flux is then converted to a bolometric flux following the relation

いて、ALMA 3 mm波で観測した。
• ダスト連続波からは、96個のコアを見つけた。質量は3.450.9Msun• 速度はN2H+を使って求めた• Ms = 0.9 ­ 2.7, and virial
parameters of 𝛼vir = 0.3 ­1.3•

ALMA-IRDC: dense gas mass distribution

13

A. T. Barnes et al.
R06 clumps
L18 cores
B21 fragments
L18 FOV
B21 FOV

12.08

R06 clumps
L18 cores

Hierarchical association

Clump

Dec (J2000)
12.08

Dec (J2000)

12.06

Cloud BCores

B21 fragments

103

Clump
Cores

102
Fragmentation
links

Core-fragments

104

101

Mass [M ]16100Core-fragments
276.50276.48RA (J2000)

276.46

RA (J2000)

10 2

10 1

100

Radius [pc]Figure 8. A schematic diagram showing the fragmentation linking procedure, which we use to determine a hierarchy of structures similar to a dendrogram
across multiple datasets covering our cloud sample (section 4.3). The left panel shows the near- and mid-infrared extinction derived mass surface density map
for Cloud B in greyscale, overlaid with circles representing the position and sizes of the cores contained within the homogenised catalogue (section 4.1). The
red circles show the Rathborne et al. (2006, R06) clumps, the solid magenta circles show the cores identified in this work (B21), and the solid green circles
show the core-fragments identified by Liu et al. (2018b, L18). The dashed magenta and green lines show the field of view (FOV) of the ALMA observations
presented in this work and in L18, respectively. The centre panel shows a zoom-in of the MM2 R06 core within Cloud B, which is highlighted on the left panel
by a solid black box. Here the clump, core and core-fragment structures have been labelled. The right panel shows the mass as a function of radius for theexample MM2 region. Here we show the fragmentation links between the host clump, and the cores and core fragments contained within its boundary.

105
Cloud A

Hierarchical structureeach core. Cores with =cf = 0 are included, as these represent theexample MM2 region. Here we show the fragmentation links between the host clump, and the cores and core fragments contained within its boundary.

105

104

=1
03

nH

2

Error
High-mass limit

100Core-frag. Cores

10 1

(K

0)
1
P

cm 3

(L
81
)

10 7

101

cm 3

cm 3

102
10 8

Mass [M ]103

Hierarchical structure
(K
M
08
)

Cloud A
Cloud B
Cloud C
Cloud D
Cloud E
Cloud F
Cloud H
Cloud I
Cloud J

10 2

Clouds(IRDCs)Clumps

10 1

100

101

Radius [pc]Figure 9. The clump-to-core fragment masses within each cloud as a function of the size-scale. The circles, crosses and plus sign markers represent
the clumps from R06, cores from this work and the clumps from L18. The
straight lines connect the symbols for each core/core-fragment to the larger,
host clump/core and then cloud (see Figure A1 in the appendix for all clouds
plotted separately). The lines and symbols have been coloured by cloud as
indicated in the legend located in the upper left. The dotted diagonal greylines show constant number densities (as labelled), and the black lines show
the high-mass star formation thresholds shown in Figure 6. In the lower right,
we show a representative uncertainty range ⇠ 15 and ⇠ 50 per cent on theradius and mass, respectively.

each core. Cores with =cf = 0 are included, as these represent thelowest level of the hierarchy;
i.e. those that are high-mass, yet do
過去文献から同じエリアのデータを引っ張ってきて合わせて
not appear to fragment. Cores with =cf = 1 are also included, as the
再解析した。その結果、質量-半径のプロットをしてみると、
core-fragment corresponds
to a central emission peak, and hence
these also represent cores that do not fragment. Cores with =cf > 1
the cloud (∼10 pc) to clump (∼0.5 pc), and to the core (<0.1 pc)
are included only if at least one of the core-fragments has a massscalesのスケールであまりシンプルなパワーローには乗ってof "cf > 16 M . We stressonce morethat the cores could have acomplex three-dimensional, density and temperature structures that
いなかった。
could influence this analysis.
For example, the emission peak(s)could correspond to a central density peak, or, in the case of astar-forming sources, an increase in temperature (e.g. those with a
70 m point source; see section 3.2). A more in-depth study of thekinematic and chemical composition of these cores will be used toaddress this in the future, and here we make the simple, generalassumption that the identified
hierarchical structure is dueto the pc (~2000 au) 以下では大質量星
小さいコアに着目。<
0.02
physical fragmentation. We separate this high-mass core sample bythose with and without signsof ongoing star formation (i.e. have noを作れない。
associated infrared point source emission; see section 3.2).
ここからhierarchicalを使って、質量は星を作る
Out of the sample of
96 cores used for this analysis, westructurefindthat 19 cores satisfy the above
requirements as potential candidates
ほど大きいが、これ以上fragmentしない天体１９個を同定しfor high-mass star formation. Out of these, we find that 6 show nosigns of active star formation,
た。whilst 13 have signs of active starformation.9 Figure 10 shows the mass and radius of the high-masscores with the connectionそのうち１３個は星形成の兆候あり。残り６個はいい候補か。
to their host clumps and contained corefragments. We plot separately
the cores that have been classified
マッハ数の中間値2.4くらいのtranssonicas quiescent and star-forming in the upper and lower panels, respectively. For reference,kinematically
we also show the cores subvirial
identified in this(medianα =0.5)
analysis that do not meet the background-subtracted mass threshold
なので、支えるには磁場が必要（median 𝐵 of 930 𝜇 G）
of ("cb > 16 M ).
We first focus on the 6 quiescent high-mass cores. We see that

•
•

•
•
•
•

9

The quiescent high-mass cores are: A3c3, C2c1, C1c1, D8c1, D6c5, D6c4.
The high-mass star-forming cores are: A1c1/2, B2c10, C2c2, C2c3/5, C2c6,
D5c7, C6c1/2, D9c1/2, F4c5, F4c8, F4c10, H3c3, I1c1.
MNRAS 000, 1–25 (2020)

Particle acceleration and magnetic field amplification in massive
young stellar object jets
Anabella Araudo, Marco Padovani, Alexandre Marcowith
Synchrotron radio emission from non-relativistic jets powered bymassive protostars has been reported, indicating the presence ofrelativistic electrons and magnetic fields of strength ~0.3-5 mG. We
study diffusive shock acceleration and magnetic field amplification in
protostellar jets with speeds between 300 and 1500 km/s. We showthat the magnetic field in the synchrotron emitter can be amplified by
the non-resonant hybrid (Bell) instability excited by the cosmic-ray
streaming. By combining the synchrotron data with basic theory of Bellinstability we estimate the magnetic field in the synchrotron emitterand the maximum energy of protons. Protons can achieve maximum
energies in the range 0.04-0.65 TeV and emit gamma rays in theirinteraction with matter fields. We predict detectable levels of gamma
rays in IRAS 16547-5247 and IRAS 16848-4603. The gamma ray flux
can be significantly enhanced by the gas mixing due to RayleighTaylor instability. The detection of this radiation by the Fermi satellite
in the GeV domain and the forthcoming Cherenkov Telescope Array athigher energies may open a new window to study the formation ofmassive stars, as well as diffusive acceleration and magnetic field
amplification in shocks with velocities of about 1000 km/s.
Comments:
17 pages, 8 figures. Accepted for publication in MNRAS
arXiv:2102.11583 [astro-ph.HE]•• 大質量星からの非相対論的ジェットから、シンク
ロトロン電波放射が報告されている。
• diﬀusive shock accelerationと原子星ジェットにおける磁場による増幅でこれが再現されるか検討してみた• non-resonant hybrid (Bell) instabilityによって磁場が増幅される。
• シンクロトロン放射体のプロトンは0.040.65TeVに達する。
• IRAS 16547-5247 と IRAS 16848-4603からはガンマ線が検出可能であることを予言する

A VLA Census of the Galactic H II Region Population
William P. Armentrout, Loren D. Anderson, Trey V. Wenger, Dana S.
Balser, Thomas M. Bania
The Milky Way contains thousands of H II region candidates identified by theircharacteristic mid-infrared morphology, but lacking detections of ionized gas
tracers such as radio continuum or radio recombination line emission. Thesetargets thus remain unconfirmed as H II regions. With only ∼2500 confirmed H
II regions in the Milky Way, Galactic surveys are deficient by several thousand
nebulae when compared to external galaxies with similar star formation rates.
Using sensitive 9 GHz radio continuum observations with the Karl G. Jansky
Very Large Array (VLA), we explore a sample of H II region candidates in order
to set observational limits on the actual total population of Galactic H II
regions. We target all infrared-identified "radio quiet" sources from the WISE
Catalog of Galactic H II regions between 245∘≥ℓ≥90∘ with infrared diameters
less than 80′’. We detect radio continuum emission from 50% of the targeted
H II region candidates, providing strong evidence that most of the radio quiet
candidates are bona fide HII regions. We measure the peak and integrated
radio flux densities and compare the inferred Lyman continuum fluxes using
models of OB-stars. We conclude that stars of approximately spectral type B2and earlier are able to create H II regions with similar infrared and radio
continuum morphologies as the more luminous H II regions created by Ostars. From our 50% detection rate of "radio quiet" sources, we set a lower
limit of ∼7000 for the H II region population of the Galaxy. Thus the vast
majority of the Milky Way's H II regions remain to be discovered.
Comments:
11 pages of text, 146 infrared/radio images of faint HII region candidates(or arXiv:2103.02049v1 [astro-ph.GA] for this version)• 銀河系においてHII領域の候補
は、 中間赤外で発見されている。
だが、電波源や再結合線が見えずにHII領域と特定されていないものが多い。

• VLAを使って、9GHzの電波連続
波のHII領域候補を調べた。

• WISE(赤外)で同定されたが、電波で見えていない天体を全部見た。
• 結果、50%の天体から電波を受けた。

• 星のモデルと比較した結果、B2
かそれより早い星ならHII領域を作れることがわかった。
• 50％の検出率から、少なくとも銀河系には7000のHII領域があると見積もった。これは、これまで確認された2500に対して、まだまだ見つかっていない領域がたくさんあることを示している。

Fragmentation in the massive G31.41+0.31 protocluster
M. T. Beltrán, V. M. Rivilla, R. Cesaroni, L. T. Maud, D. Galli, L. Moscadelli, A.
Lorenzani, A. Ahmadi, H. Beuther, T. Csengeri, S. Etoka, C. Goddi, P. D. Klaassen, R.
Kuiper, M. S. N. Kumar, T. Peters, Á. Sánchez-Monge, P. Schilke, F. van der Tak, S.
Vig, H. Zinnecker
Context. ALMA observations at 1.4 mm and 0.2'' (750au) angular resolution of the Maincore in the high-mass star forming region G31.41+0.31 have revealed a puzzling
scenario: on the one hand, the continuum emission looks very homogeneous and thecore appears to undergo solid-body rotation, suggesting a monolithic core stabilized by
the magnetic field; on the other hand, rotation and infall speed up toward the corecenter, where two massive embedded free-free continuum sources have been detected,
pointing to an unstable core having undergone fragmentation. Aims. To establish
whether the Main core is indeed monolithic or its homogeneous appearance is due to a
combination of large dust opacity and low angular resolution, we carried out millimeter
observations at higher angular resolution and different wavelengths. Methods. We
carried out ALMA observations at 1.4 mm and 3.5 mm that achieved angular resolutionsof 0.1''(375 au) and 0.075'' (280 au), respectively. VLA observations at 7 mm and 1.3 cm
at even higher angular resolution, 0.05'' (190 au) and 0.07'' (260 au), respectively, werealso carried out to better study the nature of the free-free continuum sources detectedin the core. Results. The millimeter continuum emission of the Main core has beenclearly resolved into at least four sources, A, B, C, and D, within 1'', indicating that thecore is not monolithic. The deconvolved radii of the dust emission of the sources,
estimated at 3.5 mm, are 400-500au, their masses range from 15 to 26 Msun, and their
number densities are several 1E9 cm-3. Sources A and B, located closer to the center ofthe core and separated by 750 au, are clearly associated with two free-free continuum
sources, likely thermal radio jets, and are the brightest in the core. The spectral energy
distribution of these two sources and their masses and sizes are similar and suggest acommon origin.
arXiv:2103.04953 [astro-ph.GA]• 大質量原始星G31.41+0.31という天体を高分解能ALMA観測したら、よく分解して見れた→単体ではない。

Beltrán et al.: Fragmentation in G31.41+0.31

Fitting an Analytic Magnetic Field to a Prestellar Core

• 原始星コアの偏光データをフィッティングするための公式を作ってみた
Gianfranco Bino, Shantanu Basu
• この公式は密度プロファイルによらない。
We deploy and demonstrate the capabilities of the magnetic fieldmodel developed by Ewertowski & Basu (2013) by fitting observed• 輻射輸送を使って確認した。
polarimetry data of the prestellar core FeSt 1-457. The analytic
FeSt 1-457という天体にも応用した•
hourglass magnetic field function derived directly from Maxwell'sFigure 1. Polarimetry data for the core FeSt 1–457 (Kandori et al. 2017). Each dark line corresponds to polarization in the near-infrared due to dichroic extinction
light from a backgroundstar. The direction of each polarization segment is taken to be parallel to the magnetic field. The lengths represent the degree of• theこのコアはサブクリティカル
equations yields a central-to-surface magnetic field strength ratio inofpolarization,withlongest line corresponding to about 5% polarization.
the equatorial plane, as well as magnetic field directions with relative
magnitudes throughout the core. This fit emerges from a comparisonof a single plane of the model with the polarization map that resultsfrom the integrated properties of the magnetic field and dust
throughout the core. Importantly, our fit is independent of any
assumed density profile of the core. We check the robustness of the
fit by using the POLARIS code to create synthetic polarization mapsthat result from the integrated scattering and emission properties ofthe dust grains and their radiative transfer, employing anobservationally-motivated density profile. We find that the synthetic
polarization maps obtained from the model also provides a good fit
to the observed polarimetry. Our model fits the striking feature ofsignificant curvature of magnetic field lines in the outer part of FeSt
1-457. Combined with independent column density estimates, we
infer that the core of size Rgas has a mildly supercritical mass-toflux ratio and may have formed through dynamical motions startingfrom a significantly larger radius R. A breakdown of flux-freezing
through neutral-ion slip (ambipolar diffusion) could be responsible
for effecting such a transition from a large-scale magnetic fieldstructure to a more compact gas structure.
arXiv:2103.03324 [astro-ph.SR]
0.08
0.06
0.04
0.02
0

-0.02
-0.04
-0.06
-0.08

-0.06 -0.04 -0.02

0

0.02

0.04

0.06

Magnetic fields and Star Formation around HII regions: The S235
complexR. Devaraj, D. P. Clemens, L. K. Dewangan, A. Luna, T. P. Ray, J.
MackeyMagnetic fields are ubiquitous and essential in star formation. In
particular, their role in regulating formation of stars across diverse
environments like HII regions needs to be well understood. In thisstudy, we present magnetic field properties towards the S235 complex
using near-infrared (NIR) H-band polarimetric observations, obtainedwith the Mimir and POLICAN instruments. We selected 375 backgroundstars in the field through combination of Gaia distances andextinctions from NIR colors. The plane-of-sky (POS) magnetic fieldorientations inferred from starlight polarization angles reveal a curved
morphology tracing the spherical shell of the HII region. The largescale magnetic field traced by Planck is parallel to the Galactic plane.
We identified 11 dense clumps using 1.1mm dust emission, with
masses between 33−525M⊙. The clump averaged POS magnetic fieldstrengths were estimated to be between 36−121μG, with a meanof ∼65μG. The mass-to-flux ratios for the clumps are found to be
sub-critical with turbulent Alfvén Mach numbers less than 1,
indicating a strongly magnetized region. The clumps show scaling ofmagnetic field strength vs density with a power-law index of
0.52±0.07, similar to ambipolar diﬀusion models. Our results indicate the S235

• S235のNIR偏光結果• 375の背景星を利用• HII領域を沿うような磁場構造を発見 field in the S235 complex
Magnetic

complex is a region where stellar feedback triggers new stars and the magnetic fields
regulate the rate of new star formation.
arXiv:2103.02956 [astro-ph.GA]

Figure 8. Magnetic field morphology in the S235 complex. The background image is the same Spitzer IRAC
composite map shown in Figure 1. The NIR background starlight polarizations are indicated by red-yellow pse
The lengths of the vectors indicate the polarization percentages and a reference vector length is shown at the bott

Water in star-forming regions (WISH): Physics and chemistry from clouds todisks as probed by Herschel spectroscopy
E.F. van Dishoeck (Leiden), L.E. Kristensen (Copenhagen), the WISH team (50 coauthors)(abridged) Data and results from the WISH key program are summarized, designed toprovide a legacy data set to address its physics and chemistry. WISH targeted ~80
sources along the two axes of luminosity and evolutionary stage: from low- to highmass protostars and from pre-stellar cores to protoplanetary disks. Lines of H2O,
HDO, OH, CO and [O I] were observed with the HIFI and PACS instruments,
complemented by molecules that probe UV, X-ray or grain chemistry. Most of the
far-infrared water emission from protostars is found to be compact, originating fromwarm outflowing and shocked gas at high density and temperature in at least two
physical components. This gas is not probed by low-J CO lines, only by J>14. Water isa significant, but not dominant, coolant. Its abundance is universally low, of order
H2O/H2=2E-6, pointing to shock and outflow cavity models that include UV radiation
at 100-1000 times the ISRF. In cold quiescent pre-stellar cores and envelopes, thewater abundance structure is accurately probed through velocity-resolved line
profiles, confirming basic chemistry networks. The gaseous HDO/H2O ratio of 0.025,
much higher than that of bulk ice, is representative of the outer photodesorbed ice
layers and cold chemistry. Water abundances in the inner hot cores are high, but withvariations from 5E-6 to 2E-4. Combined analyses of water gas and ice show that up
to 50% of the oxygen budget may be missing, with possible explanations discussed.
Water vapor emission from disks is weak, indicating that water ice is locked up inlarger pebbles early on and that these pebbles have settled and drifted inward by theClass II stage. Quantitatively, many oceans of water ice are available. Extragalactic
low-J H2O emission is mostly compact and collisionally excited. Prospects for future
mid- to far-infrared missions are given.
arXiv:2102.02225 [astro-ph.GA]• WISHのサマリー論文。
• ハーシェルを使って、80天体を観測• 殆どの場合、遠赤外の水の輝線
は、コンパクトで、温かいショッ
クを経験したガスから。
• J>14のCO輝線でのみ検出される。
• 水は冷却輝線だが、支配的では似。