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
Fintensity, and this blob 4is signiﬁcantly brighter in 870 μm.
The Astrophysical Journal Letters, 831:L12 (6pp), 2016 November 10 Although circular polarization can in principle be produced by Kataoka et al.
L30 W. R. F. Dent etscattering
al.
of linearly polarized light off non-spherical grains
Hull
et al.
HD 142527 4IM
Lup
CWwellTau(e.g., Tazaki
et al. 2017), StokesV has not beencharacterized for ALMA, so its detection could be spurious
IM Lupand will not be discussed in detail in this paper.
HD 163296
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 of1%the disk, with the southeast morepolarized than the northwest. For 1.3 mm, the absoluteAUdifference between the 200
two
P peaks along the major axis(a)is0.0013 ± 0.0014, which is not statistically signiﬁcant. The
parison between the Cho & Lazarian (2007) model and the SMA 340 GHz observations of HD 163296. The top row shows the prediction forasymmetry
the model for 870 μm is more apparent, although the southeast
HL
TauThe gray scale shows either thepart
HD
DG
Tau
n (left), a simulated observation of the model with the SMA (center), and the 2008 SMA observations(right).
total
offluxthe163296
disk −1does not peak (i.e., it is still increasingat theFigure 1. In
left panel,colorscale
the polarized
in units
ofofmJy
witha log scale,the grayThecontours
show
the continuum emission,
arized flux (center, right), and the blue vectors indicate the percentage
andthedirectionof thepolarized
fluxrepresents
at half-beamintervals. intensity
The centerandedgebottom
thebeamdetectedpolarized
emission).
P peaktowardand thevectorsshow
the polarization
vectors.
thatleftthe tolengthsthethepolarization
are setbe the same.
The alevels
the (a)contours are (3, 10, 30,
he model prediction (center) with the observed SMA data (bottom)
in whiteeach ofthe four
Stokesparameters (I,Q, U,NoteV, fromright).ofContour
levels vectors
northwest
is 0.0077
±to0.0010.
Drawingline offromthe−1
μJyV,beam
StokesI.
Polarization
vectors theare center
plotted ofwhere
the polarized
intensity is larger than
100,−1300,
600, 900,
1500, 1800)×σ
both rows, either multiples of 10% of the peak flux (0.9 Jy beam) in Stokes
I or 1200,
in−1incrementsof 2σ forI(=185
Q, U, and
where) σforis therms
noiseof peak throughnorthwest
the disk,
at the same
1 . The size and orientation of the synthesized beam are indicated
3σPIin=0.128mJy. In
the right panel, the colorscale displays the polarization
fraction
overlaid withpolarization±
vectors.
TheP graythe lower
leftbeamof eachpanel.
radius on
the southeastside,theP=0.0094
0.0010.
risescontours display thecontinuum emission with the same levels of the left panel. The colorscale is only
shown withtheedgesameof the polarization
vectors in the left panel.
to ∼0.0115
at theofthresholdthe disk. Therefore,
the P asymmetryn of this figure is available in the online journal.)
along the major axis for 870 μm is ∼0.002–0.004. These
Hughes et al. 2009
are relatively
moderate, butwith
the fact
that they
possibility is the difference in the grain size and its effects on asymmetries
power-lawsize distribution
a powerof −3.5 and theappearmaximum
both at 1.3
mm sizeand 870
they
instrumental
leakagethe
calibration.
(3) severalofe able to detect it given that we recover most of
the alignment
efﬁciency;
radiative Furthermore,
torque efﬁciency
grain
150μm
μm,suggests
which isthecould
most efﬁcient grain
be a real feature. If we consider total polarization PT 
MA
340Figures
GHz observations
HD 163296.theThedata
top rowshows
the prediction
for the model
thewithnightswere shared
with
other
SMAsizepolarization
projects,flux.
2 and 3of comparewiththedecreases
decreasing
grain
size
if
the
grain
is
smallersize2 to scatterthe thermal
emission and make it polarized at the2
2
Contours:
Contours:(center), and the 2008 SMA observations (right). The gray scale shows either the total flux
QStokesU
V I , the asymmetry is similar; for 1.3 mm, it isandoursolutionsfortheinstrumental
leakage
betweendelpredictions(describedinSection
4.1
below).
Polarized
Intensity
P
Polarization
thanthewavelength.
Therefore,
if
the
peak
emissionof
StokesI
observed
wavelengthof fraction
λ = P0.87
mm. The density distribution
frac
percentage and direction of polarized flux at half-beam intervals. The center and bottom
0.0012 ± 0.0014, while for 870 μm it is ∼0.001–0.004.
parameters
for theeight
quarter-wave
plates were effectively
ight
figureamountandiscomingfromgrainssmallerthan the wavelengths,
it Yangis etbasedon ashowed
previousof thethecontinuumat the sameottom)panelin eachofofeach
the four
Stokesdisplays
parametersthe(I, Q,U, V, fromleftmainly
to right).
Contour levels
al. (2017)that modelingasymmetry along
major
identical
to thoseby otherobservers,
who successfully
observed
flux
for eachsource,thebeam−1 ) in polarized
Stokes I or inincrementsof 2σ
for Q, whileU, and V,
where σ decreaseis the rms thenoiseof derived
wouldalignment
efﬁciencyandthus the polariza(Mutoet al. 2015). The
densityindistributionof dustaxis
iswavelength
not disk
expectedscattering-induced
polarization
an
Figure 1. 870 µm ALMA maps of the Class
II protoplanetary
IM Lup. The forpeak valuesof the total intensity (Stokes I),
polarized
intensity (P ), and polarization fraction images are 120.74 mJy beam , 1.37 mJy beam , and 0.011, respectively. The rms noise values incatedin the lower
left ofeachonpanel.
Non-detection
TW Hya.
axisymmetricdisk,
sotakenaxisymmetry
must
be broken if thedetect polarization
targets. 2007;
presents
contourmaps
forHD
each163296,
of the individualisGaussiantion fraction(e.g., Cho in&theirLazarian
Lazarian&Σ
the total intensityand polarized intensity
thermal dust grainsemission images
are
= 100
µJy beam toand be= 22 aµJy beam, respectively. distribution in radial andd
The black ellipses in the lower-left corners of all panels represent the ALMA synthesized beam (resolution element), which measures 0. 50 ⇥polarizationdue ofto161 scattering
evenetthough
the total
0. 40 at a position
angle of 76.9 , here
corresponding
to a linear
resolutionisof ⇠really
72 au direction
at a distance±as
10 pc (Gaia Collaboration
al. 2016).
It
is
worth
comparing
the
rms
noise
achievedwiththemeters.
Theemissionin
StokesQ
andU(linearazimuthal
Hoang2007).
Top:
Polarization
map
of
IM
Lup.
Contours
are
the
total
intensity
thermaldust
emission,plotted
at
3,
8,
16,
32,
64,
128,
256,
512,
1024
⇥
.
GM Aur and DG Tau (Hughes et al. 2009,
Grayscale is the polarized thermal dust emission,plotted (Stokes
starting at 3 I .) Linesegments are
the polarization orientationof the other hand,
intensity
is highly
axisymmetric.
On thedust vectors
emission,with lengths
proportional Errors
to the polarization
fraction P. Bottom left: Grayscale is the total intensity thermal dust
limiting precisionof thethepolarization
current SMA
polarimeter.
in, as well as in Stokes V (circular polarization), is The azimuthal
direction ofas
shownemission; contours are the polarized intensitythe
P , plotted
at 0.06, 0.1, 0.2, 0.3, 0.4,
0.5, 0.6, 0.7,the0.8, 0.9 ⇥minor
the peak of 1.37
mJy beam. Bottompolarization
along
axis
is expected
to be
⎡
⎤
right:
Grayscaleis the total
intensity thermal dust emission; contours are
the polarization fraction P= P/I, plotted at levels of2013).
2
sigma
detection
on
HL
Tau
⎛
⎞
Figure1.
ALMApolarimetricobservations
at
3.1
mm
(top,
Kataokaetal.
2alignmentofthequarter-wave
platesintroduce
instrumentalithnoise.
As
noted
inSection2,
since
Stokes
V
is
in
the
outer
regions
indicatesthepoloidalmagneticﬁeld
0.005,
0.006,
0.007,
0.008,
0.009,
0.01.
The
coordinates
andthegrayscaleareidentical
in
the
two
bottompanels.
The
ALMA
data
used
to
Figure1.makeSub-mmcontinuumpolarizationin HD163296.
The ettop
show (a) Dentthetotalandintensitypola
R2019,
Kataokaal.the2016,
al.version
2018,
Bacciotti
al.
et
al.
asymmetric
if thedust
diskpanels
is 2018,
opticallythickhas(anda (b)Rthed ) percentage
instrumental leakage calibration. Furthermore,2017),(3)1.3severalof and 870
offigure areHull
availablethein etthe onlineof this publication.
mm (middle),
μm (bottom),et
where
redthisvectors
show
⎢
⎥
⎜
⎟
exp
4angular
4width
40,min(c)marked(4)seed  of
0,min
0,max
⎜
⎟ ⎥(b). The low
conﬁguration,
while
the
radialdirection
on
the
inner
ring
region.
Superimposed
as ellipses
arerepresentations
the
three
rings
in theillustration,
total
intensity
image,“leakage”
betweenStokesparameters,
allowing
some
of
the
the
difference
betweenthemeasured
right
and
left
signiﬁcant(foracartoonFigure
6 2G 21–3 in(Stephens
et
al.
2014)
>3σ
polarization
morphology(i.e.,
vectors
have
not
been
rotated).
Vectorthe nights were shared with other SMA polarization
projects,he⎢⎝In (c), the vectors⎠ are⎣⎦ shown at
Stephens
et
al.
2017,
cf.
Kataoka
et
al.
2017,
Ohashi
et
al.
2018
d
lengthsarelinearly
proportional
to
P.
The
color
scale
shows
the
polarized
flux
intensity
and
(d)
the
polarization
angle,
both
with
the
polarization
vectors
superimposed.
in
Yang
et
al.
2017).
HL
Tau
disk
is
known
to
be
optically
indicatesthetoroidalmagnetic
ﬁeld.
There
has
been
no
flux from Stokes I to bleed into the linear Stokes parameters.
L) circular polarization, the difficulty of calibratingw).

Millimeter polarization

ALMA polarization (Class II disks)- before ALMA -upper limit ~ 1 %

1

1

I

1

P

1

00

00

I

P
frac

1

frac

and our solutions for the instrumental leakage between Stokes

what can we learn from self-scattering?
Assumption: spherical dust grains

Albedo = efficiency of scattering

a=λ/2πgrain size

Polarization at single scattering

a=λ/2πScattering-induced polarization is detectable only when a=λ/2πgrain size

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

0.19

0.33

0.1

2.1

4.2

6.3

0.1

0.3

0.5

200

276.46
0.11 0.13

0.37

0.61

0.1

1.3

2.5

3.7

0.1

0.3

0.5

275.9 406.5 598.8 882.2 17

0.12 0.08

0.28

0.48

0.1

1.1

2.1

3.1

0.1

0.3

0.5 190.6

T (K)
18
19

20

3.62
3.28

1.3
4.7

20.0
19.8

703

1318 1718

19276.48

18

19

20

17

18

19

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

4460.03300.0375

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

343.5

618.9 1115.4

16

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
A. T. Barnes
et al.

Number of cores

20

20

17.5
15.0

MedianquiescentMedianstar forming

All sources
Quiescent

.
ub

ds

This work - QuiescentunroVirial parameter [ vir ]Non-thermal velocity disp. [ NT ; km s 1 ]15
smallest cores, could be a result of several factors. Firstly, there
are
Sanchez-Monge 12.5
et al. (2013)
several systematics15to consider. L81 defined the size as the is maxStar
Kauffmann et al. (2013)
10.0
This work - Quiescentimum linear extent of the structure (i.e. larger than A eff used10 here).
forming
This work - Star7.5
forming
Secondly, the L81 10
mass estimates are determined from fundamen100
5.0
tally different observations than used in this work. In this work,
we5
5
analysis the dust continuum emission to estimate the masses, yet
2.5
L81 used molecular0line (e.g. CO) excitation arguments and 0simple
0.0
0
1
2
3
4
5
0.00
0.05
0.10
0.15
0.0
0.5
1.0
1.5
2.0
2.5
assumptions about the geometry (see e.g. Lada & Dame 2020 for
Reff [pc]log(M ) [MCM
] 95 - s
Sonic Mach number Ms [km s 1 ]further discussion of how different mass tracers can cause scatter inmasHigh 0.21
R
the mass-size relation). Thirdly, the ALMA continuum observations
0.72-1
cs~0.25kms
3. Distributionof the properties
determinedwhichacross may
the core sample (section 3.3). Shown in panels from left to right is the effective
radius ('eff ), thepresented withinFigure
this work
suffer fromspatial filtering,
Thermal sound speed
8(logarithm)
mass
("
),
and
sonic
Mach
number(M=f/f)
calculated
acrossthecore
sample(section
3.3).
Shown
as
black,
red and blue stepped
3
s
.
NT
T
0
affect this comparison to sizes and masses determined primarily
R
profiles are histograms for the whole sample, the star-forming cores and quiescent (non-starcores, respectively. The red
and ablue
ss vertical dashed lines.43 8forming)
1
1
53.
from single-dishshow
observations.
0Lthe median values for the star-forming cores and quiescent cores.
R wm
4
6
Lo
With the above systematics in mind, it is then worth considering
0.
0
1
2
3
95
M
that there may be an alternative interpretation of this result. On theMass [log(M )]C
1
2of eachwe usethe integrated flux 10
We10 finddust median temperaturesof )dust = 17.8+1.4
K.0
largest scale (>1 pc),Towecalculate
see thatthe
themass
cloudsare core,in broadagreement
10 1
0.910
(( a ; see(seeTablelower
A1) following,
These dust temperatures foreach corewith the L81 relation
panel of Figure 6). The increase
Radius[pc] are then used to calculate+32.5
a median mass of " = 18.4 15.0 M across the core sample
2 ( 'at smaller A
in " relative to L81 3seeneff (<0.1 pc) could then be
a gd
4 We assumed aSanchez-Mongeet al.
(2013)= et al. (2013),(1)(" b = 4.6+16.8
M).
typically⇠
10
percent in
Sanchez-Monge
explained bythe" fragmented
structureoftheclouds,
and,
on
the
3.7
^ a ⌫ a ()dust )Kauffmann et al. (2013)
1the absolute flux scale of the ALMA observations, and, following
Kauffmannetal.
(2013)
10
smallest scales,where
the densityof the coresthemselves
This work - Quiescent
3 is theprofiles
source distance(see Table1), 'gd (see= 141 is the
Sanhueza et al. (2017), we assume an uncertaintyof ⇠ 30 per centThis work - Quiescent
This
work
Star
forminge.g.0 Kauffmannetal.
2010a,b;Ballesteros-Paredes
et
al.
2012
for
total- (gasplus dust)-to-(refractory-component-)dust-mass ratio (asThis
work
Star
forming
dust
opacity.
These
uncertainties
in
the
dust
opacity,
dust
emission10
discussion). Thesuming
L81 arelationto accountfor muchtypical fails
interstellar
compositionof H, of
He,thisand metals;
fluxes, and the distance propagate to give an uncertainty of ⇠ 50 per
3complex structure
we 2011),
now know
existswithinmolecular
cloudsDraine⌫ a ()dust) is thePlanckfunction fora dust tempercent in masses. The histogram distribution of the size and mass of
vir=2)dust , at2007
a representative
frequencyof significant
a = 90.664 GHz, and(see e.g. McKeeature,& Ostriker
for references
showing
theαstar-formingand quiescent core samples is shownUnboundin Figure 3.
V21^ a = relations),^0 (a/a0 ) ⇡0.175
cming our, whenassuming a(seeBound
0 = 230 GHz,
It is worth noting that the dust temperatures measured
here arescatter in mass-size
andseen
observations2
1
00
V = not
1.75all(Battersby
et al. 2011)
and
^find
cm g (Ossenkopf
0 = 0.899- thenaverages sampled by the beam of the PPMAP maps (⇠12 ; shown
5 isFigure 2).
It
toosurprising
that
we
this
deviation,
9
M
0
C
s
10
mas & Henning 1994 result for an MRN size distribution with thin ice
as a black circle in the rightmost panel of Figure 2), and the beam ofhgiparticularlyas
here
investigate
scalesbelow
thatprobed
by
L81.
H
5
6
3
0.21
mantles after 10 yr of coagulation at a density of 10 cm ).
the ALMA observations used in this work is significantly smaller
.72R
0
-1
Also over-plottedon
Figure6arethehigh-mass
star
formac~0.25kms
To obtain an estimate of the sdust temperature towards each(⇠ 3 00 ; see Table 2). Therefore, smaller-scale temperature variations
Thermalsound speedtion relations
proposedby Krumholz& McKee(2008,
henceforth
8ofthecores,
we
make
use
of
spectralenergydistribution (SED)
3.
due to e.g. cold cores or embedded stellar objects could cause us
0R
3 Kauffmann
KM08)1.4and
Pillai (2010,
henceforth
KP10).
defits to the&far-infrared
dust
continuum
observedWhen
with Herschel(Hissto over- or under-estimate the temperature, respectively (e.g. see
1
a
53
.
8
0
L
m
R
Molinari
et al.
2016).
We use the mean line-of-sight results
termining their Gal;
relation,
KP10
reducedRagan et al. 2011; Dirienzo et al. 2015; Sokolov et al. 2017).
64 Lowthe Ossenkopf & Henning.
0
fromused
the for
PPMAP
project
shown
in1 Figure
2 (Marsh
et al. 2016,0
3 of 1.5(1994) dust opacity
the massby2a factor
To quantify the extent that the cores containing embedded
95determinationM
Mass[log(M)]2017).
To
check thethe PPMAP
columndensityand tem- 10 1protostars may have underestimated temperatures in the PPMAP
Cfidelity ofDunham
(see
et
al.
2010a).
Followingetal.
(2011),
we
10 1 Kauffmann
1
10 1
100
101
102
10 2 perature maps, where10possible,
100corresponding
we compare to themaps, we follow a procedure outlined in Peretto et al. (2020) for
determine the relation to account for the dust opacity used throughMass[M angular]maps from Nguyen
RadiusLuong[pc] et al. (2011), Lim et al. (2016), Zhang
determining
thetemperatureathigherresolution (⇠5 00 ). For
out this work (seesection
3.2).
The
KP10
scale
relationishence
et al. (2017),
and Soam et al. (2019). These authors independently
Figure 7.
Dynamical
state ofthe identified
cores.
panel
shows
this
estimate, theseauthors
require that
the Thecoresupper
have anidentified
Sanchez-Mongeet al.the(2013)
given as " = produced
580 ' 1.33
,
where
the
"
and
'
are
mass
and
column density and temperature maps following a more
the non-thermalvelocity
dispersion
function
effective
70 m point
source(the 70 asma flux
givenofinthe
Table
A1). radius,
The 70andm
Kauffmannetal.
(2013)
radiusin units ofconventional
M and parsec.
It is worth
just briefly
here
far-infrared
SED fitting
routine,noting
whilst accountingfor101
the lowerflux
panel
the virial toparameter
as a functionof the mass.
Shown
is shows
then converted
a bolometricflux following
the relationg
ckba

Non-thermal velocity disp. [ NT ; km s 1 ]Number of cores

14

20.0

Number of cores8• 10個の赤外線暗黒星雲IRDCについて、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

•

cm
-3
10 8

nH

2 =10

3

102

cm
-3

101

10 1

106
105
104
103

cm
-3

10 4

10 6

cm
-3

100

Mass [M ]

103

cm
-3

104

cm
-3

cm
-3

10 4

10 6

ALMA-IRDC: dense gas mass distribution

L18
This work

ALMA
ALMA

Not I/J
1.3
LarsonAll (1981)31
2.9

Sanchez-Monge et al. (2013)
Peretto et al. (2013)010
4.2.1 Mass-size relation

Kauffmann et al. (2013)
Chen et al. (2020)& Tan(2013)
Figure 6 presents the mass as a function ofKainulainen
the effective
radiusfor13
10 1

the core sample2 identified within
this work
compared 1to sources 2
1
0
10
10
10
10
10
10 2
10 1
100
taken from the literature. Shown as black-outlined red points are
Radius[pc]
TableHomogenised
3. Observational
properties [pc]of the literature sample (section 4.1).
themasses
(")
and
as
red
pointsarethebackground-subtracted
Radiuscatalogue
Homogenised catalogue
TabulatedRathborne
is the typereference, the telescope used to studyFigure
themasses6.(") determined
3.2 sizes(see Table
A2).
Shown
as
A bcomparison
of inthesection
mass andfor thecores
determined
et ofal. structure,(2006)
Rathborne et al. (2006)
106
core samples,the clouds
covered, and the wavelength and angular resolution
Henshaw
et al. (2016)
Henshaw et al. (2016)
the frommasseseffective
fromrelations
the
inthe
thiscoloured
work, to points
samplesare
taken
the andliterature,and radii
mass-sizeof the observations
(mean
The
Henshaw et al.
(2017)of major and minor axis of beam size). andHenshaw et al. (2017)
thresholds formassivecore
starcatalogue
formation.(section
Shown4.1).
as black-outlined
homogenised
literature
For a repre- red
5
Liu reference
et al. (2018)are Rathborne et al. (2006, R06), Henshaw et al.
10abbreviated
Liu et al. (2018)
pointsare thecores studied
this plot
workon("Figure), whichare connectedto theirsentative
comparison,
weinalso
7 thecore samples
This
work
(2016a, H16), Henshaw et al. (2017, H17) and Liu et al. (2018b, L18).
This work
correspondinglower(" b ; see
Tablefrom Kauffmannet background-subtractedal. (2013, K13), Perettomasses
et al. (2013,
P13),
andA2).
Highlighted
as coloured
points SM13).
are the mass and sizes of the cores and/or
Sánchez-Mongeet al. (2013,
104
Structure
Ref.
TelescopeSample_\ clumpsFigure
determinedby Liuetsmalleral. (2018b),
Henshaw
et al. (2016a),
Henshaw
6
shows
that
cores
are
typically
lessmassive,00
mm
etwithinal. (2017),
and Rathborne
et al.scatter(2006),inwhich
an orderof magnitude
masshavefor crossover
any givenwith
ef- the103
sampleexamined
thisdiagonal
work. The
properties
thesearecores
have been
fectiveradius.
Showninas
dottedgrey oflines
constant
Clump
R06
IRAM-30m
All
1.3
11 IRDCusingdensities.
the same assumptions
outlinedin thiswork (section
mean number
We see thatthe corefragments
have4.1).
Core
H16
IRAM-PdBI
H
3
3.75recalculated
High-mass
5 cmdirectly
3.
grey open
arehavethe properties
star mass limit
H17
ALMA
H(6)
1
1 (lower=H2 ⇠panel)
106 7Shown
cm 3 ,aswhereas
the points
clumps=H2 ⇠ 104taken
102 Core-frag.
Larsonthat(1981),
Kauffmann
et al.are(2013),
Peretto etyetal. (2013),
SánchezCore-frag.
L18
ALMA
Not I/J
1.3
1 fromShowing
the smallest
cores
less massive,
significantly
etthanal. (2013),
andChen
et al. (2020).
OverlaidThe
as diagonalsolid blackCore
This work
ALMA
All
3
2.9 Monge
densertheir
more
massivecounterparts.
largerclumps(KP10) (KM08)
1
Larson (1981)
10
lines
arehowever,
the mass-radiusrelations
takenfrom densities
Larson (1981,
L81),
and
could,
also
achievelargervolume
at
the
same
-2
1.33
1gcm
Sanchez-Monge et al. (2013)(L81) 580R
high-mass
starofformationthresholds
takencore
fromfragments,
Krumholz &McKee(2008,
size
scalesthesmaller
cores
and
in
the
likelyPeretto et al. (2013)
460R1.9
KM08)
andthey
Kauffmann& Pillai(2010, KP10).
Thecould
KP10imply
relation
4.2.1 Mass-size relationcase thatare centrally
concentrated.
This
thathasthebeen
Kauffmann et al. (2013)100
scaledby avolume
factor of
1.5 to matchthedustclump
opacityusedthroughout
thispaper.
Chen et al. (2020)
deriveddensityofthelargerissmaller
becausemost
Figure 6 presents the mass as a function of the effective radiusThefor horizontal dotted black line shows the mass threshold of high-mass stars
Kainulainen & Tan (2013)of their volume is at lower densities.
thecoresampleidentified
withinthisworkcomparedtosources1(>8 M ). The dashed diagonal grey lines show constant number densities
10
2
1
0
1
2
2
1
0
Overlaid on Figure 6 are several commonly adopted thresholds
10 from the 10
10 as black-outlined
10
10
10
10
taken
literature. Shown
red10points (asarelabelled).
for massive-star formation, which we can compare to the observed
themasses
(")
and
as
red
pointsarethebackground-subtracted
Radius[pc]
Radius[pc]
Homogenised catalogue
cores. The horizontal dashed black line, and dark shaded region isbmasses(")
determinedinsection
3.2
(see
Table
A2).
Shown
as
Rathborne et al. (2006)
Figure 6. A comparison of the mass and sizes for the cores determined
the limit above which a star can be considered high-mass (⇠ 8M ).
Henshaw et al. (2016)
thecolouredpointsarethemassesandeffective
radii
from
the
in this work, to samples taken from the literature, and mass-size relations
Note that this mass limit does not account for a star formation
Henshaw et al. (2017)
and thresholds
for massive
star formation.
Shown as(section
black-outlined
red a reprehomogenised
literature
core catalogue4.1). For
MNRAS
000,and,
1–25(2020) represents a lower limit for the required
efficiency,
therefore,
Liu et al. (2018)
points aresentative
the cores studied
in this work(" ),plot
which
areFigureconnected
to coretheir samples
comparison,
we
also
on
7
the
This work
mass within a structure. Ultimately, any core with a mass lower thancorresponding
background-subtracted
masses(" b ; etseeal.Table
A2).P13), andfrom lower
Kauffmann
et al. (2013, K13),
Peretto(2013,
this threshold can not in its current state form a high-mass star. Also
Highlighted
as coloured points
are(2013,
the mass
and sizes of the cores and/or
Sánchez-Mongeet al.
SM13).
overlaid on Figure 6 as several black diagonal lines are often quoted
clumps determined by Liu et al. (2018b), Henshaw et al. (2016a), Henshaw
Figure 6 shows that smaller cores are typically less massive,
et al. (2017), and Rathborne et al. (2006), which have crossover with the
mass-size relations.
withinanorderofmagnitude
scatter
in
mass
for
any
given
efIRDC sample examined in this work. The properties of these cores have been
The first of these mass-size relations is from Larson (1981,
fectiveradius.
Shown
asdiagonal
dotted
grey
lines
are
constant
recalculated using the same assumptions outlined in this work (section 4.1).
henceforth L81), and is given as " = 460' 1.9 where the normal10 1

10 8

Mass [M ]s [M ](L81)
460R1.9

100(KP10) (KM08)
580R1.33 1gcm-2

Core-frag.
1
Core
10

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

Fig. 2. Overlay of the ALMA 3.5 mm continuum emission (white contours) on the 1.4 mm continuum emission (colors) toward
cores NE and Main in G31. The white contours are 3, 6, 10, 15, 20, 30, and 40 times 1 , which is 0.16 mJy beam 1 . The 1.4 mm
and 3.5 mm synthesized beams are shown in the lower left-hand and right-hand corner, respectively. The names of the sources inwhich the emission of the Main core has been resolved (see Sect. 3.1 and Table 2) are indicated. Dashed black lines indicate thedirection of the E–W and N–S outflows mapped in SiO by B18.

3. Results and Analysis
3.1. Continuum emission of the Main and NE cores
Figure 1 (left panel) shows the dust continuum emission of
the G31 HMC at 1.3 mm obtained by Beltrán et al. (2019)
in their dust polarized emission observations carried out
with ALMA at an angular resolution of ⇠000.24, which haveclearly resolved the emission of G31 into two cores, Main
and NE. In this figure, the Main core shows a monolithic
and homogeneous appearance. The right panel shows the
new high angular resolution (⇠000.1) observations at 1.4 mm
overlaid on the ⇠000.24 observations. In the higher angularresolution map, the emission from both cores is clearly detected and the emission of the Main core starts to show
a less homogeneous morphology revealing the presence ofclumpiness. In addition, the observations have also revealed
an elongation to the north of the Main core that was alreadyconfusion between the larger cores observed with an angular resolution of ⇠000.24 and the cores resolved with higher
angular resolution, hereafter we will use the term “source”
for the latter. Sources A and B, which are located at thecenter of the core and are separated by ⇠000.20 (⇠750 au),coincide with the two free-free continuum emission sources
first detected at 1.3 cm and 7 mm by C10. The morphologiesof these two sources are quite similar to each other. Source
C lies ⇠000.30 (⇠1100 au) east of source B, while source D
is located ⇠000.45 (⇠1700 au) south of source A. This latter source is located close to the intersection of two of the
SiO outflows, the N–S and the E–W, detected by B18 (seeFig. 2). The fact that the source is located along the direction of the E–W SiO outflow and close to its center ofsymmetry suggests that source D is likely its driving source.
The NE core and the northern elongation are marginallyvisible at 3.5 mm.

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
0.08
0.08
polarization maps obtained from the model also provides a good fit
0.06
0.06
to the observed polarimetry. Our model fits the striking feature of
0.04
0.04
significant curvature of magnetic field lines in the outer part of FeSt
0.02
0.02
1-457. Combined with independent column density estimates, we00infer that the core of size Rgas has a mildly supercritical mass-to-0.02
-0.02
-0.04
-0.04
flux ratio and may have formed through dynamical motions starting
-0.06
-0.06
from a significantly larger radius R. A breakdown of flux-freezing
-0.08
-0.08
through neutral-ion slip (ambipolar diffusion) could be responsible-0.06 -0.04 -0.02
00.02 0.04 0.06
-0.06 -0.04 -0.02
00.02 0.04 0.06
for effecting such a transition from a large-scale magnetic fieldstructure to a more compact gas structure.
Figure 2. Model magnetic field fits with a line spacing not indicative of the field strength. Top: Model with R = 0.1 pc. Bottom: Model with R = 0.2 pc. Eachmodel is shown in the region where there are observed polarization segments, overlaid with polarization segments on the left, and without polarization segments
arXiv:2103.03324 [astro-ph.SR]on the right.
0.08

0.08

0.06

0.06

0.04

0.04

0.02

0.02

0

0

-0.02

-0.02

-0.04

-0.04

-0.06

-0.06

-0.08

-0.08

-0.06 -0.04 -0.02

0

0.02

0.04

0.06

-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。