星形成ゼミ
担当：SFN#336 39-45
2021/03/12
総研大 D1 竹村英晃

SFN#336 39-45
39.Atlas of CO-Line Shells and Cavities around Galactic Supernova Remnantswith FUGIN
40.ALMA view of the Galactic super star cluster RCW 38 at 270-AU
resolution
41.SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation
Eﬃciency Across the Galactic Disk

42.The asymmetric inner disk of the Herbig Ae star HD 163296 in the eyesof VLTI/MATISSE: evidence for a vortex?
43.On the pumping of the CS(v=0) masers in W51 e2e44.Constraining the magnetic ﬁeld properties of Bok globule B335 using
SOFIA/HAWC+

45.Search for associations containing young stars (SACY) VIII. -An updated
census of spectroscopic binary systems showing hints of non-universal multiplicity among
these associations-

(Full atlas including all figures is available in this URL: https://nro-fugin.github.io/2020-apjs-CO-Shell-Atlas-SNRFUGIN-IX.pdf)
ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui

An ALMA view of the Galactic super star cluster RCW38 at 270-AU resolution

Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui We report millimeter andsubmillimeter continuum and molecular line observations of the Galactic super star cluster RCW 38, obtained fromthe Atacama Large Millimeter/Submillimeter Array with a minimum angular resolution of 000 .17 ⇥ 000 .15 (' 289 AU ⇥255 AU). The C18 O image reveal many massive condensations embedded within filamentary structures extending
along the northwest-southeast direction in the center of cluster. The condensations have sizes of 0.01-0.02 pc, H2column densities of 1023 -1024 cm 2 , and H2 masses of 10-130 M . In addition, the 233-GHz continuum image reveals
two dense, small millimeter-sources with radii of 460 and 200 AU (Source A and Source B). Source A is embedded
within the most massive C18 O condensation, whereas no counterpart is seen for Source B. The masses of Source A
and Source B are estimated as 13 and 3 M at the optically-thin limit, respectively. The C18 O emission shows avelocity gradient of 2 km s 1 at the central 2000 AU of Source A, which could be interpreted as a Keplerian rotation
with a central mass of a few M or infall motion of gas. Further, the ALMA 12 CO data reveal that Source A andSource B are associated with molecular outflows exhibiting maximum velocities of ⇠30-70 km s 1 . The outflows haveshort dynamical timescales of <1000 yr and high mass outflow rates of ⇠ 10 4 -10 3 M yr 1 . These observationalsignatures suggest an early evolutionary phase of the massive star formation in Source A and Source B.

••

ALMAでsuper star cluster RCW 38を観測（分解能は最高で289 au 255 au）。

SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the• 距離：1.7
Galactic
Disk kpc、~20個のO型星（候補も合わせて）を含む~104個の星が存在、年齢~0.1-0.5 Myr

J. S. Urquhart, C. Figura, J. R. Cross, M. R. A. Wells, T. J. T. Moore, D. J. Eden, S. E. Ragan, A. R. Pettitt,
高密度かつコンパクトな233 GHz 連続波源を2つ検出。
A. Duarte-Cabral, D. Colombo, F. Schuller, T. Csengeri, M. Mattern, H. Beuther, K. M. Menten, F. Wyrowski,
L. D.• Anderson,
J. Barnes, M. T. Beltrán,
S. J. Billington,
L.18Bronfman,
A. Giannetti, J. Kainulainen, J.
Source P.
A・・・半径：460
au、質量：13
Msun、C
O condensationが付随、速度勾配（中心質Kauﬀmann, M. -Y. Lee, S. Leurini, S. -N. X. Medina, F. M. Montenegro-Montes, M. Riener, A. J. Rigby, A.
量が数Msunのケプラー回転もしくはinfall
motion）
Sánchez-Monge,
P. Schilke, E. Schisano, A. Traficante,
M. Wienen By combining two surveys covering a large

fraction of the molecular material in the Galactic disk we investigate thethe spiral arms play in the star formation• Source B・・・半径：200 au、質量：3 Msun、 C18Orolecondensationの付随なし
process. We have matched clumps identified by ATLASGAL with their parental GMCs as identified by SEDIGISM,
12COで~30-70
and usegiant A、Bにおいて、
molecular cloud (GMC)
masses, the bolometricluminosities,
and integrated clump
masses obtained• these
Sourcekm s-1のアウトフロー（dynamicaltimescale<1000yr、
P
in a concurrent paper to estimate P
the
dense
gas
fractions
(DGF
=
M
/M
)
and
the
instantaneous
star
gmc
clump
gmc
-4-10-3 Msun yr-1）の検出。
mass outﬂow
rate ~10
forming eﬃciencies(i.e., SFE
=
L
/M
).
We
find
that
the
molecularmaterial
associated with ATLASGAL
gmc
clump
gmc
clumpsconcentrated
in the spiral arms (⇠60searched for variations in the values of these physical parameters with• is Source
A、Bでの大質量星形成の初期段階
respect to their proximity to the spiral arms, but find no evidence for any enhancement that might be attributable toh"ps://arxiv.org/pdf/1907.07358.pdf

ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui(a) C18O (J=2-1)の積分強度図（12m+ACA、-2 km s-1-14 km s-1）
29' 30''Jy beam–1!km s–1

(a)
1

0

6

12

log (MJy str–1)

(b)

18

2.5

3.0

3.5

Declination (J2000)24(c) H30α（ACA）

5
7J (~1.25 μm)：青、 H (~1.65 μm)：緑、 K (~2.2 μm)：赤

6
8

10
13 12
14

9

11

IRS2

15
16
17
18
2019
21

C18Oall(2.13''×1.74'')

0

18

36

08h 59m 10s

00s

・・・233 GHz 連続波源○・・・C18O condensation

10''

(d)

54

•
Declination (J2000)

・・・O型星候補＋・・・O5.5星（IRS 2）

0.3pc
00s
Jy beam–1!km s–1

(c)

＋

IRS1

0.3pc

-47o 31' 30''08h 59m 10s

29' 30''(b) Spitzer/IRAC 5.8 μm + C18O (J=2-1) 積分強度のコントア(d) VLT JHK + H30α / C18O (J=2-1)の積分強度のコントア

3

30' 30''4.0Bが分布 (a)

30''

0.1pc
30' 30''

フィラメント状の構造に沿って、21個のC18O condensation

•

IRS2

半径0.01-0.02 pc、距離0.05-0.3 pc、H2の柱密度1023-1024 cm-2、質量10-130 Msun

H30αACA(7.22''×5.32'')

A

50''

H30αACA(7.22''×5.32'')

0.3pc-47o 31' 30''08h 59m 10s

00s
Right Ascension (J2000)

08h 59m 05s

C18Oall(2.13''×1.74'')
03s
Right Ascension (J2000)•

IRS 1、IRS 2を囲むようにH30αを観測 (c)•Source AにはC18O condensationが付随し、 Source BにはFig. 1. (a) Intensity distributions of the C18 Oall data integrated over a velocity range of −2–+12 km s−1 . Dashed contours indicate the zero level of the
C18 Oall image. Circles indicate the positions of the C18 O condensations. (b) Contour map of the C18 Oall image in (a) superimposed on the Spitzer/IRAC
5.8 µm image (Wolk, Bourke, and Vigil 2008). Contours start at 1 Jy beam−1 km s−1 with steps of 2 Jy beam−1 km s−1 . (c) H30αACA image, with thedashed contours plotted at the zero level. The velocity integration range is from −20 to 30 km s−1 . In (a), (b) and (c), the large cross indicates the central
O5.5 star IRS 2, whereas small crosses indicate the O star candidates (Winston et al. 2011). (d) A VLT near-infrared J-(blue), H-(green), and Ks -(red) image(Wolk et al. 2006) shown in the region indicated by the white box in (b). Black contours show the C18 Oall distribution plotted in (b), whereas H30αACA isshown in the red contours starting at 10σ (23 Jy beam−1 km s−1 ) with steps of 3σ (6 Jy beam−1 km s−1 ). The two boxes marked by blue lines indicate theregions of Source A and Source B shown in Figures 3(a) and (b), respectively.

付随していない。(d)ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui
0

Jy beam–1(b)

0.3 0.6 0.9

0

0.3 0.6 0.9

Offset [arcsec]
0

2

1

20

FWHM = 1008AU (0.59'')

-1

-2

1.0
Tb[K]

Jy beam–1

(a)

Tb[K]

29' 30''10

0.5

Offset [arcsec]
0.0

-0.5

-1.0

FWHM = 435AU (0.26'')

20

Tb[K]
0

6

12

0

3

6

9

h m
s
08 59 3.10

9

18

27

s
3.00
Right Ascension (J2000)

s
2.90

Fig. 3. (a, b) Magnified views of the 233-GHz image for Source A and Source B. The white contours in (a) and (b) start at half of the peak intensities, 24.9
and 48.8 mJy beam−1 , with steps of 4 and 8 mJy beam−1 , respectively. The dashed contours are plotted at the zero level. Intensity profiles along the dashed

341GHzACA(4.75''×3.65'')
Declination (J2000)

0mJy beam–1

(d)

18

30''

s
3.70
Right Ascension (J2000)

20

mJy beam–1
0

mJy beam–1
29'23''

27
10

10''

18

0

h m
s
08 59 3.90

9

-1.0

0

00s

Tb[K]

Declination (J2000)

Declination (J2000)-0.5

mJy beam–1

08h 59m 10s(c)Offset [arcsec]
0.0-47o30'46''

Right Ascension (J2000)

29' 30''

0.5

0.3pc

22''

1.0

0.3pc

-47o 31' 30''

1000AU-1

341-GHzACA(4.75''×3.65'')

2000AU

Offset [arcsec]
0233-GHzACA(7.25''×5.29'')44''

233GHz
C18O(0.17''×0.15'') (0.23''×0.21'')1

30' 30''(b) Source B
21''

FWHM = 377AU (0.22'')233GHz
C18O(0.17''×0.15'') (0.23''×0.21'')

2

FWHM = 835AU (0.49'')(a) Source A

40

0

20

0

40

lines across the peaks of the 233-GHz emission are plotted in black lines at the upper- and right-sides of the panels in units of Tb , with the best-fit Gaussian

B

233GHz(0.17'×0.15'')

curves overlaid in dashed red lines.

(a) Source A 233 GHz（12m、拡大図）

30' 30''(b) Source B 233 GHz（12m、拡大図）
50''*両ソースともにシングルピーク

A
0.1pc

0.3pc

-47o 31' 30''08h 59m 10s

00s

08h 59m 05s

Right Ascension (J2000)

03s
Right Ascension (J2000)

大質量星形成の初期条件を満たす
Table 5. Physical properties of Source A and Source B

Fig. 2. (a) A 233-GHzACA image in the same region in Figures 1(a)–(c), with contours starting at 5σ (0.19 Jy beam−1 ) with steps of 10σ (0.37 Jy beam−1 ).

Source(b) A 345-GHzACA image, with contours starting at 5σ (0.07 Jy beam−1 ) with steps of 10σ (0.15 Jy beam−1 ). Contours with dashed lines are plotted at the

rdust
(pc)

the same levels in (b).

(a) 233 GHz（ACA）(a)

peak NH2
23

(×10

b

−2

cm

MH2
)

233-GHz (0!! .17 × 0!! .15)b

rdust apeak NH2 b
24

cm

MH2 b

(M# )

(AU)

0.02

10.7 – 26.8

32(3.9) – 81(9.9)

460

65(3.8)

B

0.03

4.9 – 12.2

16(5.1) – 41(13.1)

200

62(1.2)

25

(×10

−2

A

zero level. (c) A 233-GHz image. (d) An enlarged image of the 233-GHz data toward the box in (c). Contours show the 341-GHzACA distribution, plotted at(b) 341 GHz（ACA）

341-GHzACA (4!! .75 × 3!! .65)
a

)

(M# )

rdust was measured at the half maximum. (b) The NH2 and MH2 were estimated at the optically-thin limit.

The Td of 20 and 40 K were assumed in the NH2 and MH2 estimated from the 341-GHzACA image to constrain(c) 233 GHz（12m）

MH2 ranges of the sources, whereas in the 233-GHz image, 100 K (see the text) was used as Td for Source A and(d) 233 GHz（12m）拡大図

Source B, respectively (Figure 3). The statistical uncertainty of the derived MH2 is shown in parentheses. Note that24

systematic errors arising from uncertainties of the distance measurement and the assumption of gas-to-dust
ratio/opacity give factor of ∼2 uncertainty further.

ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui
Source A

40(a) C18O (J=2-1) moment-0

500AU

0

20

500AU

40

0

20

+233 GHzコントア（12m）• 連続波源Source AにはC18O

44.5''

Declination (J2000)

Declination (J2000)

44.5''

20

100condensation が付随-60
40-47o30'45.5''-47o30'45.5''

C18O

233GHz(0.17''×0.15'') (0.23''×0.21'')233GHz(0.17''×0.15'') (0.23''×0.21'')h ms
08 59 3.80

h ms
08 59 3.80s
3.70
Right Ascension (J2000)

km s–1–10

1

244.5''

2

3

4

が、直接の証拠はない

08h 59m 3.80s

3.70s
Right Ascension (J2000)
mJy beam–1
0

35

40

20

10

-60

-40

-20
0
Vlsr [km s–1]

Fig. 5. Spectra of the 12 COACA and C18 OACA data near the peaks of the 233-GHz map for Source A and Source B. The intensities of the 12 COACA are
doubled.

(c)
70

12CO ではC18Oよりも広がった構造（30-70 km s-1）が存在

3.70s
Right Ascension (J2000)

42''(c) PV図（233 GHzピーク）
mJy beam–1

0

35

(a)

(b)

mJy beam–1→アウトフロー？（詳しくは次のスライド）
0

10

20

70

3000AU
Declination (J2000)

Vlsr [km/s]

Vlsr [km/s]

2012CO 2-1 x2

5

5

40

• Infall motionとも解釈できる

C18O(0.23''×0.21'')

C18O(0.23''×0.21'')200• 速度勾配は回転のように見える-47o 30' 45.5''-47o 30' 45.5''(c)1• ＋は233 GHzのピーク

Declination (J2000)

44.5''

08h 59m 3.80s

0

500AU

500AU
Declination (J2000)

–1

–1

4

-20
0
Vlsr [km s–1]

C18O 3-2s
3.70
Right Ascension (J2000)

km s (J=2-1) momen-1(b) C18O

(b)
3

-40

(b) Source B

30
I [Jy/beam]

C18O

(b)

12CO 2-1 x2

C18O 3-2

40

I [Jy/beam]

(a)

(a)

(a) Source A

30

mJy beam–1 km s–1

mJy beam–1 km s–1

46''(–22, –10)(+17, +40)00-47o 30' 50''08h 59m 4.0s

resolution

resolution
-5

-5
1000

0
–1000
Distance [AU]

1000
–2000

0
–1000
Distance [AU]km s–1

12CO
233GHz(0.50''×0.49'') (0.17''×0.15'')

10

3.6s
Right Ascension (J2000)

3.2s

20

08h 59m 4.0s

30
3.6s
Right Ascension (J2000)

3.2s

–2000

Fig. 6. (a) 12 CO intensity distributions of the blueshifted and redshifted outflow lobes shown in blue and red contours, respectively, with the 233-GHz back-O data(colorand blackdistribution
contours) andtheC18233-GHzcontour)for Source
A. Thevelocity
integration
range for Source A. The velocity integration range
Fig. 4.
(a) Intensityof theO data map(color(whiteand blackcontours)
and the233-GHzmap(white contour)

ground image. The velocity ranges of the lobes are shown in parentheses in units of km s−1 . The contours stars at 3 σ with steps of 1.5 σ, where the 1 σ

−1
−1
−1
−1
start
at 3σ
mJy
beam
km s−1) with stepsof 3σ
1σ (135(45 mJykm),) whereas
theofwhite
contours
start
is −2
to (135
+4 km
s−1
. The
black
contours
start atmJybeambeam−1
kmss−1
with steps
1σ (45mJy beam
km s−1 ), whereas the white contours startvalue of the blueshifted and redshifted maps of Source A are 30 and 65 mJy beam−1 km s−1 , respectively. The triangle depicts the approximate position of−1−1
5σ (2.25
beam). (b)
Firstof themJy
C18beamO data(color)superimposed
withofthe
map (color)of the superimposed with the contour map of theat 3σmJy
(1.35
mJy−1
beam
) withmoment
steps ofmap
5σ (2.25). (b) First
moment map
thecontour
C18 O data

the vibrational transition of the H2 emission detected by VLT (DeRose et al. 2009). (b) Color map of |vmom1 − vsys | for the lobes, with the 12 CO contour

8

peak
the 233-GHz
Thethetwopeakarrowsindicate
the233-GHz
directionsimage.
of the The
blueshifted
and indicate
redshifted
C position
O data ofin (a).
The crossimage.
indicates
positionof thetwo arrows
theoutflow
directions of the blueshifted and redshifted outflow

maps plotted in (a), where vmom1 is the first moment generated from the 12 CO data for the velocity ranges shown in panel (a), whereas the systemic velocity

citylobes
diagramof theC18 O 6.
data.
C18 O data weresmoothed
beOa data.
velocityof were1.5 kms−1 . Thetoverticalshown
in Figure(c) thePosition-velocity
diagramof theto
C18
theresolution
C18 O datasmoothed
be a velocity resolution of 1.5 km s−1 . The verticalof the source vsys was measured from the C18 OACA data as 2 km s−1 (Figure 5). The cross indicates the 233-GHz peak for Source A.

18

33-GHzSourcethe
A. direction of the 233-GHz peak of Source A.
whitepeakline ofindicates

27

I[

10

ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
0

-60

-40

-20
0
Vlsr [km s–1]

20

40

Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui

Fig. 5. Spectra of the 12 COACA and C18 OACA data near the peaks of the 233-GHz map for Source A and Source B. The intensities of the 12 COACA are

Source A（C18O condensationが付随）

doubled.

𝑣!"!# − 𝑣$%$

12CO moment-0 + 233 GHz
42''

(a)

10

• Source A：10-4 Msun yr-1、Source B：10-3 Msun yr-1(b)

mJy beam–10• Mass outﬂow rateが大きい• 大質量星形成のモデルから示唆される値とconsistent

20

Declination (J2000)

3000AU• Dynamical timescaleが短い（<1000 yr）(–22, –10)46''

• 大質量星形成の初期段階？

(+17, +40)

• Accretion rateはoutﬂow rateの数倍で(Beuther et al. 2002;-47o 30' 50''km s–1

12CO

233GHz(0.50''×0.49'') (0.17''×0.15'')08h 59m 4.0s

10

3.6s
Right Ascension (J2000)

20

Maud et al. 2015) 、 accretion rateを 10-4 - 10-3 Msun yr-1

30

08h 59m 4.0s

3.2s

3.6s
Right Ascension (J2000)

3.2s

Fig. 6. (a) 12 CO intensity distributions of the blueshifted and redshifted outflow lobes shown in blue and red contours, respectively, with the 233-GHz background image. The velocity ranges of the lobes are shown in parentheses in units of km s−1 . The contours stars at 3 σ with steps of 1.5 σ, where the 1 σvalue of the blueshifted and redshifted maps of Source A are 30 and 65 mJy beam−1 km s−1 , respectively. The triangle depicts the approximate position of
the vibrational transition of the H2 emission detected by VLT (DeRose et al. 2009). (b) Color map of |vmom1 − vsys | for the lobes, with the12とすると、降着するガスの温度は100-500 K（Hosokawa andOmukai 2009のモデルより）• しかし、近赤外線や中間赤外での顕著な増光は見られない

CO contour

maps plotted in (a), where vmom1 is the first moment generated from the 12 CO data for the velocity ranges shown in panel (a), whereas the systemic velocityof the source vsys was measured from the C18 OACA data as 2 km s−1 (Figure 5). The cross indicates the 233-GHz peak for Source A.

• ダストによる減光？
• episodic accretion scenario？

27

Source B

12CO moment-0 + 233 GHz
18''

(a)

• より詳細にA、Bについて調べるためにはさらに高空間分解能観測𝑣!"!# − 𝑣$%$をする必要がある。

(b)• 低周波数観測（Band3など、dust opacityの影響が小さく

Declination (J2000)

12CO
233GHz(0.50''×0.49'') (0.17''×0.15'')なる）→質量の推定(+17, +50)
22''

• 2-3倍程度分解能の良い観測→質量の推定3000AU08h 59m 3.5s

010

• 他の分子輝線→ケプラー回転の証拠が得られれば、原始星km s–1mJy beam–1(–68, –16)
-47o 30' 26''

20
2.5s

Right Ascension (J2000)

3000AU1040

7008h 59m 3.6s
3.2s
Right Ascension (J2000)

Fig. 7. The same plots shown in Figure 6 are presented here, but for Source B. In (a), the 1 σ levels of the 12 CO maps are 90 and 100 mJy beam−1 km s−1
for the blueshifted and redshifted lobes, respectively. In (b), only the blueshifted lobe is plotted.

の質量を求めることができる。＆Source BのCO
depletionの可能性についても調べることができる。

ALMA view of the Galactic super star cluster RCW 38 at 270-AU resolution
Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui
C18O (J=2-1) moment-0
29' 30''Jy beam–1!km s–1

(a)
1

0

6

12

C18O condensationについて5.8 μmlog (MJy str–1)

(b)

18

2.5

3.0

3.5

•

フィラメント状の構造（幅~0.1-0.2 pc）に沿ってC18O

4.0

condensationが見つかったが、Source Aに付随するもの以外

Declination (J2000)

2

は、連続波源に付随していない。

3
4

•

5
7

30' 30''

6
8

10
13 12
14

9

量も大きい（> 10 Msun）。

11

IRS2

15
16
17
18
2019
21

C18Oall(2.13''×1.74'')

Declination (J2000)

29' 30''

H30α
0

18

36

08h 59m 10s

このような構造は近傍の小質量星形成領域では見られない(e.g., Onishi et al. 2002; Tachihara et al. 2002)。

0.3pc
00s
Jy beam–1!km s–1

(c)

•
IRS1

0.3pc

-47o 31' 30''08h 59m 10s

CondensationのH2柱密度は大きく（>1023 cm-2）、質VLT JHK00s•

ただし、質量やサイズは銀河面のIRDCで見られるコアの

10''

(d)

54

ものとconsistent (e.g., Ohashi et al. 2016)•B

いずれ、（中）大質量星へ進化しうる。

フィラメントの形成メカニズムについて

30''

0.1pc
30' 30''• O型星がフィードバックにより、シェル状の構造を周囲に形成し、IRS2星形成を誘発するモデルが存在する(Elmegreen and Lada
H30αACA(7.22''×5.32'')-47o 31' 30''

A

50''

H30αACA(7.22''×5.32'')

0.3pc

08h 59m 10s

00s
Right Ascension (J2000)

08h 59m 05s

C18Oall(2.13''×1.74'')
03s
Right Ascension (J2000)

1977)。
• この天体に関しては、O型星とフィラメントの距離が近すぎる。

Fig. 1. (a) Intensity distributions of the C18 Oall data integrated over a velocity range of −2–+12 km s−1 . Dashed contours indicate the zero level of the
C18 Oall image. Circles indicate the positions of the C18 O condensations. (b) Contour map of the C18 Oall image in (a) superimposed on the Spitzer/IRAC
5.8 µm image (Wolk, Bourke, and Vigil 2008). Contours start at 1 Jy beam−1 km s−1 with steps of 2 Jy beam−1 km s−1 . (c) H30αACA image, with thedashed contours plotted at the zero level. The velocity integration range is from −20 to 30 km s−1 . In (a), (b) and (c), the large cross indicates the central
O5.5 star IRS 2, whereas small crosses indicate the O star candidates (Winston et al. 2011). (d) A VLT near-infrared J-(blue), H-(green), and Ks -(red) image(Wolk et al. 2006) shown in the region indicated by the white box in (b). Black contours show the C18 Oall distribution plotted in (b), whereas H30αACA isshown in the red contours starting at 10σ (23 Jy beam−1 km s−1 ) with steps of 3σ (6 Jy beam−1 km s−1 ). The two boxes marked by blue lines indicate theregions of Source A and Source B shown in Figures 3(a) and (b), respectively.

• H30αのリング状の構造のsemi-major / minor axisの長さはそれぞれ~0.3 pc / 0.2 pc。
• モデル計算では、圧縮層の密度は1、2桁大きくなる (e.g.,
Hosokawa and Inutsuka 2006)。
• Source AとBはO型星が周囲を電離する前に形成されたと考Galactic Disk

J. S. Urquhart, C. Figura, J. R. Cross, M. R. A. Wells, T. J. T. Moore, D. J. Eden, S. E. Ragan, A. R. Pettitt,

SEDIGISM-ATLASGAL:
DenseF.Gas
Fractionand Star
the F.
Galactic
Disk
A. Duarte-Cabral, D. Colombo,
Schuller,
T. Csengeri,
M. Formation
Mattern, H. Eﬃciency
Beuther, K.Across
M. Menten,
Wyrowski,L. D. Anderson, P. J. Barnes, M. T. Beltrán, S.
J. Urquhart
Billington,
J. S.
et L.
al. Bronfman, A. Giannetti, J. Kainulainen, J.
Kauﬀmann, M. -Y. Lee, S. Leurini, S. -N. X. Medina, F. M. Montenegro-Montes, M. Riener, A. J. Rigby, A.
Sánchez-Monge, P. Schilke, E. Schisano, A. Traficante, M. Wienen By combining two surveys covering a large•

fraction of the molecular material in the Galactic disk we investigate the role the spiral arms play in the star formation
process. We have matched clumps identified by ATLASGAL with their parental GMCs as identified by SEDIGISM,
and use these giant molecular cloud (GMC) masses, the bolometric luminosities,
and integrated clump masses obtainedP
in a concurrent paper to estimate P
the dense gas fractions (DGFgmc =Mclump /Mgmc ) and the instantaneous star
forming eﬃciencies (i.e., SFEgmc =Lclump /Mgmc ). We find that the molecular material associated with ATLASGAL
clumps is concentrated in the spiral arms (⇠60searched for variations in the values of these physical parameters withrespect to their proximity to the spiral arms, but find no evidence for any enhancement that might be attributable tothe spiral arms. The combined results from a number of similar studies based on diﬀerent surveys indicate that, whilespiral-arm location plays a role in cloud formation and HI to H2 conversion, the subsequent star formation processes
appear to depend more on local environment eﬀects. This leads us to conclude that the enhanced star formation
activity seen towards the spiral arms is the result of source crowding rather than the consequence of a any physicalprocess.

ATLASGALとSEDIGISMのカタログを使って、銀河の渦状腕が星形成に与える影響を調べた。
•

The APEX Telescope Large Area Survey of the Galaxy (ATLASGAL)：大質量クランプの870μmを用いた銀河面サーベイ（サイズは~ 0.5 pc、質量は~ 500 Msun）•

Structure, Excitation and Dynamics of the
Inner Galactic Interstellar Medium (SEDIGISM )：14
13CO、C18OのJ=2-1輝線、ショックや高密度ガスのトレーサーを用いたサーベイ (APEX)•

Clump-渦状腕の間の距離とclumpの物理量の関係を調べたが、渦状腕に近づいたことによる物理量の変化は見られなかった。

•（他の研究結果と組み合わせると）渦状腕はHIからH2への変換には寄与する（進める）が、その後の星形成活動には、個別の環境の影響が大きいと考えられる。

•

渦状腕で星形成が誘発されているように見えるのは、単に混み合っているからだと思われる。
h"p://arxiv.org/abs/2012.01464v1

SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the Galactic Disk
J. S. Urquhart et al.

ATLASGALのクランプとSEDIGISMの輝線データのマッチング（速度の決定）•

全部で5754個のATLASGALクランプ（870 μm）•

各クランプの870 μmのピークを中心に30”の範囲で13COとC18O (2-1)のintensityの合計を求める。
•

606個のATLASGALクランプでは「ベースラインが正しく決められていない」and/or「> 30 km/sを超える速度幅を
DGF and SFE Across the Galactic Disk
持つ（すべて銀河中心）」という理由でサンプルから外す。（残り5148クランプ）

5

Table 1. Fitted Gaussian parameters to CO spectra extracted towards AT-

clumps.
•LASGAL輝線データから3σ以上のローカルピークを抜き出
CSC Name

Transition

Tmb

vlsr

1

FWHM
1

(K)
(km s )(km s )す（ガウシアンでフィット）。
CO 10.75 ± 0.11 8.51 ± 0.03 4.48 ± 0.10

AGAL300.504 00.176

13

13 CO

4.64 ± 0.12

27.26 ± 0.07 3.84 ± 0.19

5.32 ± 0.16

36.86 ± 0.03 1.64 ± 0.07

7.28 ± 0.13

37.73 ± 0.03 2.75 ± 0.08

C O
1.64 ± 0.12 8.51 ± 0.20 4.02 ± 0.58
1. シングルピーク（~
25%）：クランプと輝AGAL300.748+00.097
CO 20.74 ± 0.14 36.98 ± 0.01 2.32
± 0.02
18

13

C18 O
13 CO
C18 O
13 CO
C18 O

26.18 ± 0.10 39.62 ± 0.02 5.92 ± 0.06
線データをマッチング
8.47 ± 0.11
39.61 ± 0.04 4.67 ± 0.14

AGAL301.136 00.226
AGAL301.279 00.224

2.78 ± 0.17
37.54 ± 0.05 1.56 ± 0.12
2. 複数のピーク＆一番強いピークが他に比べてNotes: Only a small portion of the data is provided here, the full table will
be available in electronic form at the CDS.

2倍以上強い：もっとも強いピークをクラン
プとマッチング（C18Oが検出されたらそれも参考に）

3. 複数のピーク＆一番強いピークとそれ以外の差があまりない：目で確認•4998個のクランプについて速度を決定（サンプルの97%）•
Figure 3. Examples of the 13 CO (black line) and C18 O (grey) spectra extracted towards three dense clumps identified from the ATLASGAL survey;
the former is offset by 2 K from zero intensity. The results of the automatic
Gaussian fitting are overlaid in red and blue and the velocity assigned to theclump is indicated by the green vertical dash-dotted line.

Figure 4. Example of how integrated emission maps can be used to determine the most likely velocity component. The two maps above correspondto the two velocity components seen towards AGAL342.901 00.081 (seelower panel of Fig. 3); the upper panel shows the integrated emission for the⇠ 86 km s 1 component while the lower panel shows the integrated emission at ⇠ 3.5 km s 1 . The contours show the distribution of the ATLASGAL 870 µm emission. The emission at ⇠3.5 km s 1 is strong, compact
and correlated with the position of the ATLASGAL source and, therefore,
is considered to be the most likely component. The SEDIGISM beam size
is shown in the lower left corner of each map.

1108個のクランプについて速度を初めて決定した。

•

Figure 5. Histogram of the differences between previously assigned velocities (Urquhart et al. 2018) and the velocities obtained from analysis of the
SEDIGISM data discussed in this paper. The red curve shows the Gaussian
fit to the distribution, which has a FWHM of 1.1 km s 1 . The bin size is
0.2 km s 1 .

4998個のクランプが1709個のGMCに分類された。

In other cases where the integrated intensities towards a par-

DGF andEﬃciency
SFE AcrossAcrossthe Galactic
7 DiskSEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation
the Disk
Galactic
J. S. Urquhart et al.
背景：12COの積分強度

Figure 6. Galactic longitude-velocity distribution of all ATLASGAL sources located between 300 < ` < 60 for which we have been able to assign a velocity.
The greyscale image shows the distribution of molecular gas as traced by the integrated 12 CO (1-0) emission for comparison (Dame et al. 2001). The red circles
mark the positions of clumps where the velocity has been drawn from Urquhart et al. (2018), while the yellow circles mark the positions of clumps where avelocity has been determined from the work presented in this paper. The location of the spiral arms are shown as curved solid lines (colours are as given in
Fig. 1). The positions of the four main spiral and local arms have been taken from the model of Taylor & Cordes (1993) and updated by Cordes (2004).

Table 2. Assigned velocities and derived physical parameters for the ATLASGAL
clumps. The clumps’ star formation efficiencies are taken directly from（上）：ATLASGALクランプの分布（範囲内のすべて、速度を決Urquhart et al. (2018).
CSC Name
AGAL341.126
AGAL340.392
AGAL340.311
AGAL341.131
AGAL340.349
AGAL339.176
AGAL339.403
AGAL340.304

Transition
00.34700.43100.43600.42100.43400.39100.41400.376

13 CO
13 CO
13 CO
13 CO
13 CO
13 CO
13 CO
13 CO

めていないものも含む、8948個）

13 CO(T ) C18 O(T ) FWHM
vlsr
SFEclump
Rgc
GMC Name
13 CO
mb
mb
1
1
(km s )(K)
(K)(km
) (Lbol / M clump ) (kpc)• s 黄色：新たに速度を決めたクランプ
41.4
14.6
8.0
4.32
8.68
5.07 SDG344.929+0.3022
45.9
4.4
1.7
5.02
0.94
4.94 SDG343.133 0.4493
48.1
5.5
3.1
4.31
0.52
4.85 SDG343.133 0.4493（左）：渦状腕の分布、色は上の図と対応
35.0
3.6
1.4
4.31
0.56
5.39 SDG344.257 0.3774
47.9
3.9
1.5
6.31
0.48
4.86 SDG343.133 0.4493•3.11薄いグレー：ATLASGALの観測領域
37.3
11.4
5.4
13.84
5.45 SDG341.016 0.1252
39.1
4.5
1.9
3.37
1.19
5.34 SDG341.016 0.1252• 濃いグレー：SEDIGISMの観測領域
51.6
7.2
4.1
4.17
1.00
4.71 SDG343.133 0.4493

8

SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the Galactic Disk
J. S. Urquhart et al.

J. S. Urquhart et al.

SFEgmc = Â Lclump /Mgmc ). These two parameters, as defined, areindependent of heliocentric distance. The distances of clumps andthe GMCs havebeen
determined using different rotation curves（上）：ℓ >
10°のクランプ（赤）とGMC（青）と、もっとも近い渦状腕との速度の差
and velocities so distances of matched objects can be different(±0.5 kpc). We can eliminate this by first dividing the catalogueの絶対値の累積分布。
masses and luminosities by the catalogue distance squared to obtain cloud and clump properties per kpc2 before taking the ratios of• 速度差が10 km/sのところで傾きが急になっている。
clump and cloud values(i.e. all normalised
to 1 kpc).
The clump masses and cloud masses have been calculated us• クランプ、GMCはともに渦状腕に沿って分布しているが、クランプの方がその傾向がing different dust opacity values. The ATLASGAL survey uses anopacity of 1.85 cm2 g 1 at n0 = 350 GHz (Schuller et al. 2009 andreferences therein) while the SEDIGISM survey uses a CO-H2 con強い→高密度クランプを持つGMCの方が渦状腕との相関が良い。
version (CO X factor ⇠ 1.08±0.19⇥1021 cm 2 (K km s 1 ) 1 ; see
Paper I for details) that is derived from the dust opacity used by• 渦状腕の位置はpulsar
dispersion measurementから求めており、分子ガスと相関が
the Hi-GAL column density maps (Elia et al. 2013; Schisano et al.
212020; k0 = 0.1 cm g at n0 = 1200 GHz; Hildebrand 1983). Fol良いのは自明ではない。
lowing the reference of Elia et al. (2013) the opacity is assumed toscale as:

• 大質量生形成のトレーサーの検出の有無でクランプを分類して比較しても、速度分布
✓

n

◆b

kn = k0に優位な差は見られなかった。
n0

Figure 7. The upper panel shows the cumulative distribution of the absolute
velocity separations of all ATLASGAL clumps (red curve) and SEDIGISM
GMCs (blue curve) from their nearest spiral-arm locus in ` v space. The
lower panel shows the frequency distribution of velocity differences of all
GMCs (grey) and those associated with dense gas (red).

clouds through the spiral arms (⇠7-10 km s 1 ; Burton 1971; Stark& Brand 1989; Reid et al. 2009).
4.2 Dense gas fraction
In this and the following subsection we will restrict our analysis to
clouds associated with dense gas that are located between ` = 300(1)

If we use the values used by Elia et al. (2013) stated above
and set n = 350 GHz and set b = 1.75 we obtain a value ofk350 µm = 8.5 ⇥ 10 3 cm2 g 1 for the dust opacity at the ATLAS（下）：GMC（グレー）と高密度ガスが付随するGMC（赤）と渦状腕の速度差のヒス
GAL frequency. The difference in the dust results in a discrepancy
of a factor of 1.61 in the masses and to compensate for this we have
トグラム
multiplied the ATLASGAL clump masses by this factor. We alsoneed to apply this factor to the 5s column densities threshold for
the ATLASGAL survey ⇠ 7.5 ⇥ 1021 cm 2 at 20 K (Schuller et al.
• クランプの方が相関が良いこと以外は、大きな違いはない。
2009), which is now ⇠ 1.5 ⇥ 1022 cm 2 . The column density sensitivity of SEDIGISM is ⇠ 0.95 ⇥ 1021 cm 2 , assuming a 5s noise• GMCの分布をガスシアンでフィットすると1σ~
9 km/s。
of 3.5 K, channel width of 0.25 km s 1 and the mean CO X factor.
The material traced by ATLASGAL has column densities approxi円運動（Brand&Blitz
km/s）渦状腕間の分子雲の移動に伴う流mately •15 times
larger than the GMCs identified by1993、12.8
SEDIGISM.
In Fig. 8 we show the distribution of the DGF with respect to1971;& arms.
Brand 1989; Reid et al. 2009、~7-10 km/s）
the velocityれ（Burton
offset between the clouds
and theirStark
nearest spiral
We note that there is a large scatter in the values of the DGF (dex= 0.4) and の値と似ている。
this results in approximately 20 per cent of the clouds
having a DGF value above 1. Values of the DGF above unity are
somewhat unexpected and might suggest the presence of an observational bias such as line-of-sight projection effects that might
result in the dust continuum emission being systematically pushed
to higher values. However, the spatial filtering of the large-scale
dust emission that occurs as part of the data reduction effectively
removes any contribution from large-scale diffuse emission and theoverall number of dense clumps means that the chance of multiple
clumps lying along the same line-of-sight is negligible. An alternative explanation for these high values for the DGF is the large
uncertainties in the parameters that go into the calculation of theSEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the Galactic Disk
DGF and SFE Across the Galactic Disk

J. S. Urquhart et al.

9of a spiral arm (red histogram). It is clear from this plot that there isno significant difference between these two populations (a KS-test※銀河中心（渦状腕の位置が正確には決まっていない）をサンプルから除くために、
gives a p-value of 0.06, which is a bit less than 2-sigma). This isconsistent with the results reported by Dib et al. (2012), Eden et al.

300°
ℓ<
350°
𝑏 ≈ 0.5°）にあるサンプルも除く。
(2012, <2013);
thesewill beに注目し、観測領域の端（
discussed in detail in Sect. 5.
The mean value for the DGF is 0.35, which is considered
to be an upper limit. This is in excellent agreement with themean value determined from the ratio of ATLASGAL masses to
CHIMPS clouds
by Rigby（DGF、∑
et al. (2019). It is𝑀
useful,at thisDense
gasreported
fraction/𝑀*!& ）&'(!)
point, to remember that only 11-17 per cent of GMCs identified in
SEDIGISM (Paper III) are associated with an ATLASGAL clump.
（上）DGFのヒストグラム（赤は渦状腕との速度差が10km/s以下のもの）
We can obtain an estimate for the global DGF by considering themass of clumps and GMCs located between 300 and 350 in• total20%のGMCて、DGFが1を超えている。
longitude, yielding a value of 0.16. This is also likely to be an upper
limit as the extraction algorithms used to identify coherent struc• data 視線方向の射影などの観測バイアスの影響？
tures in the
cubes are only able to allocate approximately twothirds of the 13 CO emission to clouds (e.g. Rathborne et al. 2009;• globallarge-scaleの放射はデータリダクションの段階で除かれており、クBarnes et al. 2016). The
DGF is therefore significantly lower
than the upper limit derived here. It is also interesting to note thatthe overall shape of the distribution
is log-normal distribution (the
ランプの総数（が領域の大きさに比べて少ないこと）は、視線方向
blue curve overplotted on Fig. 8 shows the result of a log-normal
fit to the full sample: a KS-test
gives a p-value of 0.59 indicatingに複数の構造が重なる確率を無視できることを意味している。
there is no significant difference between distribution of the data
and a log-normal distribution). This could indicate that variations•
質量推定の不定性？
in the DGF from cloud to cloud are the result of a collection of essentially random processes (Eden et al. 2015), the effects of which• in thisCO-X
factorやwouldダストの
opacity の影響が大きいと考えられる
are multiplicative andcase the extreme
sources
not beabnormal. The statistical properties of the DGF are given in Table 4.
（数倍の精度しかない）。
The lower panel of Fig.
8 shows the DGF distribution as afunction of offset from the nearest spiral arm: this plot reveals there• is no “全サンプルの”DGFの平均値は0.35でenhancement of dense gas in clouds with respect to their prox- Rigby et al. (2019) の結果とconsistentimity to the spiral arms. A linear least-squares fit to the data returns
13
a slope that is very close to zero ( 1.8 ± 0.97 ⇥ 10 3 ), confirm-•

Figure 8. Distribution of the GMC dense-gas fraction (DGF) determinedby summing the masses of all embedded clumps and dividing by the GMC
mass. The upper panel shows the distribution of the whole GMC sample(grey histogram) while the red histogram shows the distribution of clouds
located within 10 km s 1 of a spiral arm. The lower panel shows the distribution of DGF of each cloud as a function of velocity offset from theirnearest spiral arms. The dashed blue line shows the results of a linear leastsquares fit to the log10 [DGF] and the velocity offset; this has a slope of( 1.8 ± 0.97) ⇥ 10 3 , which is essentially flat. The bin size used in the upper panel is 0.25 dex. The blue curve is the result of a log-normal fit to the
distribution of the whole GMC sample.

サンプルを限定したときは、0.16（分子雲の同定では COのemissionの2/3しing there is no significant correlation between these two parameters(see dashedblue line on the lower panel of Fig. 8).
か分子雲に割り当てられていないので、上限値になる）
The lack of any enhancement of the DGF with proximity to
the spiral arms is also consistent with the results of Csengeri et al.
（下）もっとも近い渦状腕との速度のオフセットとDGF(2016a), who calculated the fraction of emission from dense gas
by comparing the emission detected in the ATLASGAL survey• to the渦状腕に近づいても（速度が近くても）、DGFは変わらない。
total dust emission detected by the Planck space mission(⇠5 per cent). They also refer to this as a dense gas fraction, but• over
最小二乗法でフィットすると、傾きは0に近くなり、両者に優位な相関関
it is averaged
the whole line of sight and includes a contribution from diffuse material not associated with the star formation
process. This therefore
represents a lower-limit, and so is similar to
係はない。
the global DGF referred to above.
4.3 Star formation efficiency

Log[SFE GMC] 936
Log[SFE CSC] 2829

0.27
0.02
0.75
0.27
0.01
0.80
VLSR < 10 km s 1
0.47
0.02
0.47
0.32
0.03
0.79
0.30
0.02
0.82

0.21
0.27

2.58
2.09

1.94
2.72

2.11
1.30
SEDIGISM-ATLASGAL:0.48
Dense
Gas Fraction and Star Formation Eﬃciency Across the Galactic Disk
0.26
2.58
1.94

Log[DGF]
513
Log[SFE GMC] 513
Log[SFE CSC ] 1685

10

0.30

2.09

2.72

J. S. Urquhart et al.

GMC

Table 4. Summary of physical properties of the whole population of clumps
and the four evolutionary subsamples identified. In Col. (2) we give thenumber of clouds in each subsample, in Cols. (3-5) we give the mean values, the error in the mean and the standard deviation, in Cols. (6-8) we givethe median and minimum and maximum values of the samples.
Parameter#

Log[DGF]
936
Log[SFE GMC] 936
Log[SFE CSC] 2829
Log[DGF]
513
Log[SFE GMC] 513
Log[SFE CSC ] 1685

x̄

ps

(N)

s

All
0.45
0.01
0.45
0.27
0.02
0.75
0.27
0.01
0.80
VLSR < 10 km s 1
0.47
0.02
0.47
0.32
0.03
0.79
0.30
0.02
0.82

xmed

xmin

xmax

0.45
0.21
0.27

2.11
2.58
2.09

1.30
1.94
2.72

0.48
0.26
0.30

2.11
2.58
2.09

1.30
1.94
2.72

J. S. Urquhart et al.

ATLASGAL clump

Star formation eﬃciency（SFE*!& =∑ 𝐿&'(!)/𝑀*!& ~SFE）
IMFは普遍的と仮定（上） SFEgmcのヒストグラム（赤は渦状Figure 10. As Figure 8 but for the ATLASGAL clump star formation efficiency. The dashed blue line shows the results of a linear least-square
fit to the Log10 [SFEclump ] and the velocity offset; this has a slope of(9.1 ± 1.5) ⇥ 10 4 , which is essentially flat. The bin size used in the upper panel is 0.25 dex.

腕との速度差が10 km/s以下のもの）（下）もっとも近い渦状腕との速度のオフセットとSFEgmc

Figure 9. As Figure 8 but for the star formation efficiency. The dashed blueline shows the results of a linear least-squares fit to the log10 [SFE] andthe velocity offset; this has a slope of 6.9 ⇥ 10 5 ± 1.7 ⇥ 10 3 , which isessentially flat. The bin size used in the upper panel is 0.25 dex.

variations in the IMF, stochastic or intrinsic, are affecting our own•
observations remains unclear, and is presently difficult to fold in
our interpretation of the (L/M) ratio. Therefore, we make the assumption that the IMF is universal, as a working hypothesis.
The upper panel of Fig. 9 shows the SFE distribution of all
clouds (grey) and the subsample located within 10 km s 1 of a spiral arm. The distributions of these two populations are very similar
to each other, and a KS-test confirms that they are not significantly
different (p-value of 0.79). The lower panel of Fig. 9 shows the SFE
as a function of velocity offset of the clouds from their nearest spiral arm. This plot shows no evidence for any significant change in
Figure 10. As Figure 8 but for the ATLASGAL clump star formation efthe SFE
as theblue
velocitydifference
increases
betweenthe GMC
gmcdashed
ficiency.
Theline shows
the resultsof a linear
least-squareandthespiralarms.
A
linearleast-squares
fit
to
the
velocity
offfit to the Log10 [SFEclump ] and the velocity offset; this has a slope of
3
4 , whichhas
set± andlog
zero gradient(6.9bin⇥ size10 5used
± 1.7
⇥ 10
(9.1
1.5) ⇥
1010 [SFE]is essentiallyflat. Thein theup- ;bluedashedperseepanel
is 0.25dex. line on the lower panel of Fig. 9), indicating that
there is no correlation between these two parameters. From this we
can conclude that there is no dependence of the star formation effivariations
inathe
IMF, proximity
stochastic tooraintrinsic,
are affecting our own
ciency on
cloud’sspiral arm.
observations
remains unclear,
andprovide
is presently
difficult
toformation
fold inFor completeness,
we alsoplots ofthe star

SFEと渦状腕との速度差に相関はない。

DGF Star
and SFE
Across the Galactic
Disk Across
11
SEDIGISM-ATLASGAL: Dense Gas Fraction andFormation
Eﬃciency
the Galactic Disk
J. S. Urquhart et al.
GMCの銀河中心からの距離のヒストグラム（全サンプル、渦状腕との速度差が10 km/s以下、高密度ガスを含む、渦状腕との速度差が10 km/s以下で高密度ガスを含む）• サンプル全体では特徴的な分布はない• 渦状腕との速度差が10 km/s以下のGMCでは、~ 4 kpcと~5.75 pcにピークがある。
• ~ 4 kpc：Norma armの（太陽から）近い側• ~ 5.75 pc：Scutum armの近い側とNorma armの遠い側• 渦状腕との速度差が10 km/s以下で高密度ガスを含むGMCについても、
Figure 11. Distribution of GMCs as a function of Galactocentric distance.
The distribution of the whole sample is shown in grey, those located within
10 km s 1 of a spiral arm, are shown in blue. The clouds associated with
dense gas are shown in green and the population of these clouds that arewithin the 10 km s 1 of a spiral arm are shown in red. The bin size used is
0.5 kpc.

~ 4 kpcと~5.75 pcにピークがある。
• 各距離ビンにおける、ほとんどの高密度ガスを含むGMCが、渦状腕との速度差が10 km/s以下で高密度ガスを含むGMCにも分類されている。

efficiency within individual ATLASGAL clumps (as calculated in
Urquhart et al. 2018) in Fig. 10. The distribution of the clump SFE
is well modeled by a log-normal distribution (the log-normal fit to
the data is shown by the blue curve overplotted on Fig. 10; a KS-test
gives a p-value of 0.97 indicating there is no significant difference
between the them) and its mean value is approximately an orderof magnitude larger than the GMC SFE but the fraction associatedwith the spiral arms is approximately the same and has a similar
distribution with respect to the spiral arms.
The statistical properties of the SFE for both the clumps and
GMCs are given in Table 4.
4.4 DGF and SFE as a function of Galactocentric distance
We have used the radial velocities obtained from the CO analysis in• 高密度ガスを含むGMCは渦状腕に起因する？

Figure 12. Distribution of the GMC dense gas fraction and star formation
efficiency as a function of distance from the Galactic centre. The verticalshaded areas on both plots correspond to regions in Galactocentric radius
where the cloud population is dominated by the near sides of the Norma and
Scutum arms (4 kpc and 5.75 kpc, respectively). The green circles indicate

SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the Galactic Disk
DGF and SFE Across the Galactic Disk

11

J. S. Urquhart et al.
（上）DGFと（下）SFEの銀河中心からの距離についてのプロット• 影がついているのは、Norma arm と Scutum arm の位置。
• 緑の点は、0.5 kpcごとの平均値• 渦状腕がDGFやSFEを大きくする場合、ピークが見られるはず。
（上）DGF• 全体的にフラットで、渦状腕がDGFを大きくする証拠は得られなかった。
（下）SFE• 正の相関が見られるが、中心から離れるにつれて明るいGMCにバイアスがかかっている可能性がある。
• 渦状腕の位置でピークは見られず、渦状腕がSFEを大きくする証拠は得ら
れなかった

まとめ
• 渦状腕がDGFやSFEを大きくする証拠は得られず、GMC間のDGFやSFEのバリエーションも小さい。
• 個々のGMCでの局所的な要因の方が星形成に与える影響が大きいと考えられる。
Figure 12. Distribution of the GMC dense gas fraction and star formation
efficiency as a function of distance from the Galactic centre. The verticalshaded areas on both plots correspond to regions in Galactocentric radius
where the cloud population is dominated by the near sides of the Norma and
Scutum arms (4 kpc and 5.75 kpc, respectively). The green circles indicate
the average value calculated for each 0.5 kpc.

• 渦状腕と分子ガスの分布には強い相関があることが分かったが、星形成を誘発している証拠は得られなかった。
• ただし、物質を集める役割はある。

of the CS masers in high mass star forming regions might be the result of a required high CS abundance as well as
due Constraining
to attenuation ofthe
the magnetic
maser emission
insideas well as outside
the hot core.
ﬁeldproperties
of Bokofglobule
B335 using SOFIA/HAWC+
Niko Zielinski, Sebastian Wolf, Robert Brunngräber

Constraining the magnetic field properties of Bok globule B335 using SOFIA/HAWC+
Niko Zielinski, Sebastian Wolf, Robert Brunngräber Thanks to their well-defined shape and mostly isolated loca-

tions, Bok globules are suitable objects for studying the physics of low-mass star formation. To study the magnetic
field of the prototypical Bok globule B335, we obtained a spatially resolved polarization map with SOFIA/HAWC+ at
a wavelength of 214 µm. For the first time, these observations reveal that polarization holes in Bok globules, that is,the decrease in polarization degree towards their dense centers, also occur in the far-infrared wavelength regime. The
observed polarization pattern is uniform with a mean polarization angle of 48 ±26 and a magnetic field strength of⇠ 142 µG. Moreover, we use complementary polarimetic data for B335 obtained at near-infrared to millimeter wavelengths to analyze and constrain the magnetic field across diﬀerent scales. By applying the 3D Monte-Carlo radiativetransfer code POLARIS (Reissl et al. 2016), we developed a model for the density and magnetic field structure aswell as for the dust properties of this globule. We conclude that the column density towards the center of B335 is too
low to cause the observed polarization hole in B335 via dichroic absorption (Brauer et al. 2016). Furthermore, we
conclude that the eﬀect of self-scattering has no significant impact on the observed polarization. Adopting dust-grainalignment via the radiative torque mechanism, a combination of the interstellar radiation field and the central star asradiation sources is consistent with the decrease in polarization degree at the outer regions of B335 (⇡ 104 au from thecore). However, the model fails to explain the low polarization degree within the inner 5000 au.

•

Bok globule B335の偏光観測をSOFIA/HAWC+ 214μmで行った。
•

Bok globule：HII領域内の、星を形成しうる可視で暗い領域（<8 104 au~0.4 pc）。

•

遠赤外線の波長域で、偏光度が中心に向かって減少するpolarization holeを、初めて観測した。

•

Self-scatteringは観測された偏光にさほど寄与しておらず、dichroic absorptionを起こすには中心部の密度が低すぎる。

•

15

観測された偏光はダストの輻射トルクによる整列、星間輻射場と中心星の輻射の組合せで再現できる。
•

ただし、中心から~104スケールについて。<5000 auのスケールの低い偏光度はこのモデルでは説明できない。
h"p://arxiv.org/abs/2012.05889v1

Constraining the magnetic ﬁeld properties of Bok globule B335 using SOFIA/HAWC+
Niko Zielinski, Sebastian Wolf, Robert Brunngräber

Bok globuleのpolarization hole（ミリ波・サブミリ波で観測例あり）•

仮説1：密度や温度が高くなるにつれてダストが整列しなくなる•

仮説2：中心部の複雑な構造がみれるほどの分解能に達していない•

仮説3：中心部にいくほど、ダストが細長くなっていない•

仮説4：密度高い領域に整列していないグラファイト粒子が集まっている•

仮説5：光学的厚み、ダストサイズ、化学組成の組合せ

Bok globule B335•

距離：~100 pc•

サイズ：~0.1 pc•

East-west方向にoutﬂow•

dynamical timescale~3 104 yr•

中心星の質量~0.4 Msun•

Emvelope mass~5 Msun

観測結果•

中心部ほどemissionは強くなる
が、偏光度は減少。

•

中心部の北東領域の偏光度のupper limitは0.25 %。

•

偏光度は非一様（平均は48 ）。

Minimum grain radius
amin
5 nm
sizes amax 2 [0.25, 50] µm and globule masses Mgas 2 [1.0, 5.0] M .
Maximumgrainamaxanalysis(l) 250
50 µm
To provide
a basisfor radiusa profoundof nm
the -polarization
maps, we first derive constraints on the structure, mass, and dust
b
the
ﬁeld
propertiesof theet Bok
globuleB335.
For this purpose,
we useproperties of Bok globule B335 using SOFIA/HAWC+
Notes. aConstraining
Kandorial. 2020;
Evans
et al.magnetic2015
a model based on a Bonnor-Ebert density distribution, which is
Zielinski,
Robert Brunngräber
motivated by the rather simple circular shapeNikoof theA&A
object.
For Sebastian
proofs:manuscript Wolf,
no. B335_SOFIA_HAWC_plus
the same reason, it is often used in modeling studies of Bok glob(T Tauri
witheta al.
luminosityof L? et= 1.5
L andBrauer
a radiusofules(e.g.,star)
Harvey2001; Kandori
al. 2005;
et al.
B335のモデル（MonteCarlo
3D輻射輸送コードPOLARIS）
R
=
1.5
R
.
These
valueswere
introduced
by
Evans
et
al.
(2015)?
Table
1.
Overviewoftheparametersconsideredin
2016; Kandori et al. 2020). The radially dependent density struc-the Bok globuletocalculatetheinfallparameters
rate of B335.
観測：SOFIA＋JCMT
model
for 密度分布
B335.
are marked with the symbol "l".
tureis •givenby massFree
compact, homogeneous, spherical dust grains consistgles We adopt
8
with ing of silicate and graphite
Olofsson2011) with a>⇢0 ·(Olofsson
R0 2 , if &r 
R0
>
>
1
>
an- mass abundance
<
ratio
of
62.5%
silicate
and
37.5%
graphite. We
2
Parameter⇢(r) = >(4) Value⇢0 · r , if RSymbol
0 < r  Rout>
The apply the canonical>gas-to-dust
ratio
of
100:1
and
assumethat>:0,
if r > R .
re- the grain sizes are distributed
Fig. 5. Radial profile (I vs. r) for a= 2 µm and Mgas = 4.6 M in
Globule structuresimilarly to theoutpower-law distri- comparison to the observed profile atmax= 214 µm.
bution•of theダストサイズ分布（ISMと似た分布、形状は球）ISM:
The Outerdensity radius
structure is characterized R
byoutthea truncation ra-13 100 au
dius R0 , Truncation
defining
an innerregiondensity.
(l)dn(a)/ radius
a 3.5
da , with aconstant<Ra <amax , The trun(5) 1 000 au
in cation radius is a free parameter and wemin
obtain0 its value
below by computed based the temperature maps. The procedure is as folGas
mass
Mgas
(l) this region, 1-5
MFor a given maximum dust grain size, the globule masslows:
5 % fitting
radialbrightness
profile ofofB335.
Outsidewhere•thedn(a)represents
the number
grainswith a radiusinthe
ダストの組成：62.5%
silicate
and
37.5%
graphite（質量）isadjusted
to match the observed peak emission. Therefore, insity thedensity
towards
the outerradiusRoutthe. Following
range[a,Central
a +decreasesda] (Mathis
et al. 1977).
While
we setminimum
star
this
case,
the
globule mass is closely connected to the maximum
um- Kandorietal.
(2020),
we
use
a
value
of
R
=
13
100
au.
For
dust grain
size
to amin = 5 nm, as often usedoutforb the description• Luminosity
ダストサイズの下限：ISMでよく使われる値
dust
grain
size.
This yields di↵erent combinations of maximumL
1.5
L
ion thecentralheatingsourceweapplyapre-mainsequence
starof the dust phase of the ISM, it is questionable?whether the same
b
dust
grain
sizes
and globule masses. Constrained by the range
R? (amaxpage
argument
applies
for the maximum grain
size
= 250
nm 1.5 R• Radius
ダストサイズの上限：ISMでは250
Article number,
3nm、of 11 of possible maximum grain radii amax 2 [250 nm, 50 µm], thefor the ISM; see e.g., Mathis et al. 1977). On the one hand,
Fig.mass
4. Overview
observed fluxes and the corresponding fluxes resultcorresponding globule
range is Mof
gas 2 [1, 5] M .
Dustproperties
B335の0.87
mmのALMA観測より数十μm
Bok globules
represent
an early
stage of star formation withing from
POLARISsimulationsforcombinations of maximum grainTofurther
constraintheparameters
amax and Mgas
, wedensitiesMinimumthat – in contrast
the conditionsin protoplanetary
grain toradiusamin
5 nm the numericallycomparederivedfluxeswithobservational
sizes
a
2
[0.25,
50]µm
andglobulemasses Mgas 2 [1.0, 5.0] M .
max
auスケール）disks – Maximum
are （~100
not favorable
forradius
grain growth.
On(l)the other
hand,
grain
a
250
nm
50
µm
results
at
a
wide
range
of
wavelengths
obtainedwithsingle-dishmax
Yen et al. (2020) proposed the presence of grains with sizes of awith a resolution similar to that of SOFIA/HAWC+.
A&A proofs: manuscript
no. B335_SOFIA_HAWC_plus• of ↕はフリーパラメーター
few tensµm in B335 based on 0.87 mm ALMA
observations, telescopes
In
particular,
B335 was observed with the 15-m James Clerkabbuton a size
scale etof al.
⇠ 100
au, thatis, onsignificantly
Notes.
Kandori
2020;
Evans
et scales
al. 2015
Table 1. Overview of the parameters considered in the Bok globule
Maxwell
Telescopeat 450 µm, 600 µm, 800 µm, 850 µm,
smaller thanmodel
representive
for our
SOFIAareobservation.
However,for B335. Freeparameters
marked with thesymbol "l".1100 µm, and 1300 µm (Chandler et al. 1990; Launhardt et al.
given this uncertainty, we consider maximum dust grain sizes
for our model inParameter
the range of 250 nm - 50
µm. An overviewValueof 2010). The result is given in Fig. 4 and Table B.1. The model
ベストフィット
Symbolthe considered parameters and chosen values for our model is with a maximum dust grain size of amax = 2 µm and a cor(T Tauri
star)with astructure
luminosity of L? = 1.5 L andrespondinga radius massof of Mgas = 4.6 M best fits the observations
given
in Table
1. Globule
2
a
(
=
0.35,
see
Outer radius
Rout
100 au et al. (2015) Table B.1 and Fig. 5 for the corresponding radial
R? = 1.5 R . These
values were introduced
by13Evans
Truncation
radius theRof(l)000 au brightness profile). The resulting best-fit model has an opticalIn the following,
we constrain
mass
B335 and the 1max0B335.
to calculatethemassinfall
rate
of
Gas
mass
M
(l)
1-5
M depth of only ⌧214µm ⇡ 0.1 towards the center.
gas brightness profile
imum grain size by fitting the observed radialWe
adopt
compact,
homogeneous,
spherical
dust
at = 214 µm. For
this purpose,
we perform radiative transfergrains consistCentralstar
An withalternative
ing of silicate
andpublicly
graphite(Olofsson& Olofsson
2011)
a and commonly used approach to derive thesimulations
usingLuminosity
theavailable
Monte
L? b Carlo 3D radia1.5 L mass of 1interstellar clouds relies on the assumptions of optically2btivetransfer
code
POLARIS(Reissletal.
2016).
Wefind
thatmass abundance
ratio of 62.5% silicate
and 37.5%
RadiusR?
1.5 Rgraphite . We
thin dust emission and constant temperature. Here, one calcuaapply
truncationof R0 =gas-to-dust
1000 au provides
the ofbest100:1
agreement
Fig.that
4. Overview of observed fluxes and the corresponding fluxes resultthe radiuscanonical
ratio
and assumehydrogengas mass MH basedon the measured
fluxgrain
Dustproperties
between the shape of the derived radial brightness profile and the lates theing from POLARIS
combinations
Fig. 5.simulations
Radialforprofile(I vs.ofr)maximum
for amax= 2 µm and Mgas = 4.6 M in
Minimumgrainradius
a
5
nm
the
grain
sizes
are
distributed
similarly
tothepower-lawdistridensity
byminsizes amax 2 [0.25, 50] µm and globule masses Mgas 2 [1.0, 5.0] M .
observed one. First, the temperature profile of B335 is calculatedcomparison
totheobserved
profileat=
214 µm.
Maximumgrainradius
a
(l)
250
nm
50
µm
max
bution of the on
ISM:
self-consistently
the basis of the embedded
radiation source
2S⌫ D
and the dust distribution.
Subsequently, the re-emission maps are MH = Mdust X =このとき中心部で、𝜏
X,(6)+#, -! ~0.1
Notes. a Kandori et al. 2020; b Evans et al. 2015d (⌫) B⌫ (⌫, T d )3.5
1dn(a) / ada ,amin < a < amax ,(5)

in a hydrogen mass of ⇡ 3.5 M . Considering helium andheavy elements, the total gas mass amounts to Mgas ⇡ 1.36
MH ⇡ 4.7 M , which is very close to the value derived without
Constraining the magnetic
ﬁeld properties
oftheBok
the additional simplifications.
We conclude thatderivedglobule B335 using SOFIA/HAWC+
mass based on the 214 µm SOFIA/HAWC+ observation is
NikoZielinski,
Sebastian
Wolf,
Brunngräbercomparable
to values
publishedin earlier studies(⇡ 3.5 Robert- 5M ;
Wolf et al. 2003; Kandori et al. 2020).

Self-scatteringが効くか→大きくは寄与しない

properties of the Bok globule B335

磁場

A&A proofs: manuscript no. B335_SOFI

Fig. 7). W
モデル（self-scattering）
3.4. Origin of the polarization at 214 µm: Polarization due to1000 auof 214 µm
with oblat
Fig. 6. Synthetic polarization map of B335 at a wavelengthn,self-scattering in the context of Bok globules?
an oblatederived under the assumption of self-scattering. Total intensity is shownesphere. Th(re-emission and scattered light) with overlaid polarization vectors int
Until a few years ago, it was generally accepted that the red. Polarization due to self-scattering causes a centrosymmetric poVoshchinne(sub)millimeter polarization observed in Bok globules was larization pattern for our model. The resulting degrees of polarization
See Fig. 8e
観測結果
caused by the emission orabsorption of magnetically aligned for our self-scattering model are very low (see reference vector in thetion map bsnonspherical dust grains. However, as spatially resolved po- upper right corner of the polarization map). The isocontour lines markAs the).
larization maps of selected circumstellar disks (e.g., HL Tau, 20%, 40%, 60%, and 80% of the maximum intensity. The image hasderof ⌧21t B335_SOFIA_HAWC_plus
no.
Stephens et al. 2017) in the same wavelength range are best ex- been convolved with a Gaussian profile with the SOFIA/HAWC+ band
has little is00plained by the process of self-scattering, whether this process E FWHM of 18.2 .
cally, the pd Fig. 7). We consider the same dust mixture as in Section
3.3, havebut an impact in the case of Bok globules remains tothe core b
could
also
6 with oblate grains with an axis ratio a/b = 2. Here, the radius of
depth e↵ec
investigated.
t an oblate dust grain is given by the radius of anbe
scattering
asanexplanationfortheobserved
polarization
patternequal-volume
the polariz
To answer
thisquestionisimportantbecauseonly
thepod sphere. The optical properties were calculated with a codeand polarization degree, and conclude that the measured polarby N.
this e↵ectlarization
by dichroic emission or absorption of magnetically ization is caused by the emission of nonspherical magneticallys Voshchinnikov5 based on Voshchinnikov & Farafonov(1993).
whether or; See Fig. 8 (right) for the resulting (beam-convolved)aligneddust
grainsallowsustodraw
conclusions
about
the
unpolarizaaligned dust grains.
mation ofderlying magnetic field structure, while polarization by scattertion map based on our best-fit model.
gin needs
ing
primarily
providesinformation
about
the
dust
grain
size
and
As the derived optical depth for our model is only on the orbeen obse
Fig. 7. Intensity
map of B335
withpolarization vectors rotated by 90
3.5.
Magnetic
field
structureandstrength
distribution(e.g., Brunngräber & Wolf 2019). In this context,
der of ⌧214µm ⇡ 0.1 towards the center, the dichroic
extinction
to indicate the projected local magnetic field direction. The length of IRAS4A (itMoreis important
to note that the polarization pattern alone is not As discussed
has little
impactpolarizationon the netmap
polarization
the vectors
is arbitrary
constant anddoes not indicate
the strength
in this sim
in Sect.
3.4, theandmeasuredpolarization
is causedbyFig.
6. Syntheticof B335 atofa B335.
wavelengthofspecifi214 µm
00auniqueindicatorof the underlying polarization mechanism be- the emission
cally,
thepolarization
degreedecreases
from
the
outer
regions
toof the ofmagnetic
field. Themagneticallybeam size of 18.2
is indicatedin the lower following
derived under the assumption of self-scattering. Total intensity is shown
nonsphericalaligneddust grains.
left.
the
core by onlya factor light)of ⇠ with
0.95 overlaid(see Fig.polarization
8 left).
The
opticalcause
polarizationby scattering can generate patterns that mimic In this case,(re-emission
and scatteredvectorsinthe polarization vectors rotated by 90 indicate thee red.
depthPolarizatione↵ect can due
therefore
be excluded
to explain
theof patterns produced by a toroidal or poloidal magto self-scattering
causesa centrosymmetricpothe decrease
polarization
projected
magnetic
field direction of B335 (Fig. 7). As the pos larization• 142 field
46strengthμG（チャンドラセカールフェルミ法）
pattern forour model.
Thetheresulting
ofofpolarization
the polarization
degreetowards
dense degrees
regions
B335.
As et al. 2016; Brunngräber & Wolf 2019).
netic
field (Yangof a142±46µGpattern,
at 214 µm.
calculated
magnetic
larization vectors
show
uniform
the Theresultingmagnetic
d forself-scattering
are inverylow (seevectorinthethisoure↵ectcannot be model
excluded
generalforreference
otherWhile
Bok
globules,previous numerical studies of self-scattering were mostly field structure
field strengthof
B335 isoninthe
the scales
same range
as derived
for
is
also
uniform
resolvedbyourob-other- upperrightcornerofthepolarization
map).
The
isocontour
lines
mark
/#0
/1
whether or not the optical depth could have an impact
on theonfor-circumstellar disks (e.g., Yang et al. 2017; Brun-• 𝜌focussedg cm、𝑣=2003).
0.14
cmof
s/# 、
4 5.41×10µG,
seethee.g.,
Wolf
et field
al.
For
B335,
*.$ =globules
2(34
40%,
60%, and 80%of themaximum
intensity.
image orihas
servation. To get(⇠100-300
an estimate
formagnetic
strengthu, 20%,
mationof polarizationholes
or canbe excluded
asngräber
aThepossible&band
Wolf 2020), its role in the case of Bok globules has B335, weWolf
et al.
report a magnetic field strength
134+46
convolved
a Gaussianespecially
profile withasthethis
SOFIA/HAWC+- beenapply
the(2003)
Davis-Chandrasekhar-Fermi
methodof(Davis
39 µG
gin needs
to bewithinvestigated,e↵ect has already9#+.:barelybeen
addressed.
However,acompleteanalysisofselfFWHM of 18.200 .
at
850
µm.
As
the
polarization
pattern
at
214
µm
is
more
uniforms E
1951;& Fermi 1953):
𝜎5 =Chandrasekhar
26.8
been observed in the case of the young stellar object NGC1333
than at/6.8
850 µm (see Fig. 1 in Wolf et al. 2003), the magnetic field
scatteringin Bok globules is beyond the scope of this study. Inof IRAS4A •(Ko 偏光パターンが観測結果と異なる
et al. 2020). The degree of polarizationthat is seenrstrength,which
was calculated
the basis of the far-IR data, isstead, we
4⇡
von
h in
this simple
is too high.
This
aspect polarizationis consideredin limit
the ourselves to examining whether self-scattering
scattering
as anmodel
explanation
for theobserved
patternturb
Fig. 8. Cor
slightly higher.B =
⇢
.
[CGS]
(7)
gas
plays
a
role
in
B335
at
a
wavelength
of
214
µm.
For
our
anal-r following
section.
and polarization
degree, and conclude that the measured polarsity (left) a3¯ strength, it is possible to cal•
光学的に薄く、ダストサイズが小さいので偏光度が小さい
✓
Based
on
the
derivedmagneticfield
y ization is caused by the emission of nonsphericalysis,
we simulate self-scattering by spherical dust grains basedmagneticallyvectors for
culatethemass-to-flux
ratio
M/
.
This
ratio
indicates
whether- aligned dust grains.
on the temperature profile of the dust and consider the best-fit Here, ⇢gas describes the gas density, vturb the turbulence velocity,aligned dus
the core is in collapse or equilibrium. Following Crutcher et al. 60% and 80- 23.4model parameter values for the dust and structure of B335 found and ✓¯ the¯standard
deviation
the meanpolarization
cd(2004),
we calculate
the ofobserved
mass-to-flux
ratioangle
to the ✓criti- with a Gaus
in Section 3.3.
of the polarization
vectors.
One should keep in mind that, in thecal mass-to-flux
ratio:• polarization
光学的に薄い（absorptionの影響は小さい
t,r 3.5. Magnetic field structure and strength-! ~0.1 ）ので、dichroic
The resulting
pattern due𝜏+#,to self-scattering
best case, the magnetic field should be estimated as a function of
5,t P[%]µG
is causedby2 ) increasesshowsa centrosymmetricpattern (see Fig. 6) and is qualitatively distance (M/
21 N(HG- As discussed in Sect. 3.4, the measured polarization
to the)observed
center =because
density
towards the (8)7.6 · 10the·.
the emission of nonspherical magnetically aligned
dust grains.
• observed
モデル：ベストフィットのモデル＋完全に整列したダスト（磁場の方向は一定）(M/ due)crit to an insufficient numberBcm
mc 22.6
di↵erent
from thepattern in Fig. 1. Furthermore, the core. However,of2 polarization vec3.7. Influe
rotated by 90
d- In this case, the polarization vectors
lowindicate
opticalthe
depth (⌧214µm ⇡ 0.1, see Sect. 3.3) and the small grain tors, we estimate the magnetic field strength based on the mean
1.0
radia
Here,
N(H
)
describes
thecolumndensityandBthemagnetic
0.0
2
projected
magnetic
field
direction
of
B335
(Fig.
7).
As
the
poI/Imaxs• （左）外側から内側にかけて、偏光度が~0.95しか下がらない
size (compared to the observing
wavelength) result in both low value, considering
theWegasusedensityof forourN(H
best-fitmodel (averfield strength.
the value
to our2 ) corresponding
Fig.
8. Correlation
polarization
degreenormalized
intenlarization
vectors between
show a uniform
pattern,theandresultingmagnetic
18
3
scattering efficiency and a low degree of scattered light polariza- aged overthe 3model
region)⇢gas =3.3).
5.41We· 10
g cm. The value
y field
best-fit(see Sect.
considerthe averagevalue of The derivesity (left)
and synthetic(right)withresolved
overlaid polarization
structureis also intensity
uniform mapon thescalesby our ob+12.9
tion.
Consequently,
the polarization
degree is several magnitudes for ✓¯ = the•（右）モデルをスムージングしたintensity
map＋偏光ベクトル-- vectors
for4 .aTo
Bokglobulewith Mfor= 4.6
M
at = 214µm.
Perfectly
26.8
deg
is
calculated
using
the
95%
confidence
incolumn
densityinsideacirclearound
the
peak
value
with a significant
servationgetanestimatethemagnetic
field
strength
of
6.5
22
2
smallerthantheobserved
one(Kataokaetal.
2015).
Therefore,
sr aligned
dust
grainsareassumed.
The
isocontour
lines
mark
20%,
40%,
radiusof
3
andobtain
N(H
)
⇡
6.28
·
10
cm
.
This
yields panded be
tervalforthestandard
deviation.
We
apply
a
turbulenceveloc2
we apply the Davis-Chandrasekhar-Fermi method (Davisl.- B335,
60%
and
80%
of
the
maximum
intensity.
The
image
has
been
convolved
1
wecan
excludeasignificant
contribution
of
the
e↵ect
of
selfa mass-to-flux
ratio ofet⇠ al.
3.4,1987)
whichandindicates
the magnetic nificant coity of 0.14
cm s (Frerking
derivethat
a magnetic
Chandrasekhar & Fermi 1953):-- 1951;
with a Gaussian with the SOFIA/HAWC+ band E FWHM of 18.200 .
field is too weak to support against gravity by magnetic pressure tion in B3

Constraining the magnetic ﬁeld properties of Bok globule B335 using SOFIA/HAWC+
N. Zielinski: Magnetic field properties of the Bok globule B335

Niko Zielinski, Sebastian Wolf, Robert Brunngräber
ciencytorqueQ , definedasat large scales (⇡ 104 - 105 au). For the comparisonof the mag3=>
Radiative(RAT)
ここでは、𝛼= −3に固定（平均値）して、𝑄;field< を観測から求める。
netic
direction,
we
infer
for
the
near-IR
results
that the8ref（各パラメーターはダストのサイズ、材質に依存）>>if 1.8 apolarization vectors are parallel to the magnetic field (dichroic
< Q
⇣
⌘↵Q
(9)
RAT eﬃciency Q = >ref
A&A
proofs:manuscript
no.
B335_SOFIA_HAWC_plus>
absorption). Thus, the magnetic field direction is in a compa: Q · 1.8 aotherwise,
rable range to（interstellar
our findings (✓¯radiation
Fig. 7). The magnetic
B ⇡ 140 , ﬁeld）
field
strengthinthenear-IR(12-40µG,
Bertrang
et al. 2014) is
is particularly important (Draine & Weingartner 1996; Hoang
refsmallerbyafactor
of
5-15
than
in
the
far-IR(⇠
142 µG). As& Lazarian 2014). The required parameters Q and ↵Q dependon grain shape and material, but are not well constrained. The the gas density is significantly lower on larger scales where theparameter ↵Q is in the range of -2.6 to -4 (Lazarian & Hoang magnetic field is traced through observations at near-IR wave2007; Herranen et al. 2019; Reissl et al. 2020). We set ↵Q to an lengths, this finding can be attributed to the dependence of theaverage value of -3, while we adjust the parameter Qref to fit the magnetic field strength on the gas density in the flux-freezing
limit (Heitsch 2005).
observations. The values for Qref are assumed to vary between
0.01 and 0.4 (Lazarian & Hoang 2007; Reissl et al. 2020).
In the following, we consider the radiation fields of subse- 870 µm and 1.3 mm:
quent sources separately:
Maury et al. (2018) report that the magnetic field structure at– a) Central star only,
1.3 mm shows an ordered topology within the inner envelope,– b) ISRF only and,with a transition from a large-scale poloidal magnetic field in the– c) central star and ISRF combined.
outflow direction to a strongly pinched one in the equatorial direction (Fig. 10 bottom-right). Consistent with the 1.3 mm obserWe apply our best-fit model derived in Sect. 3.3. The tempera- vations, Yen et al. (2020) find that the magnetic field at 870 µm
ture map of B335 is calculated self-consistently on the basis of (Fig. 10 bottom-left) changes from ordered to more complex andthe aforementioned radiation source(s) and the dust distribution. asymmetric structures within the inner 100 au of B335. The variSubsequently, on basis of the temperature maps, the re-emission ation of the polarization angles is significantly higher at 870 µm
maps are computed. In the final step, equivolume oblate dust and 1.3 mm than at 214 µm. However, one has to keep in mind
grains with an axis ratio of 2:1 are considered. Figure 9 shows that a possibly more complex structure at 214 µm is not resolved
the correlation between polarization degree and normalized in- due to the lower resolution of SOFIA/HAWC+. The observed
tensity for four di↵erent values of Qref 2 (0.1, 0.2, 0.3, 0.4).
polarization degrees at wavelengths of 870 µm and 1.3 mm are
We find that the ISRF ensures alignment of the dust grains in the range of ⇠ 1-11%, while for the 214 µm (HAWC+) casein theregionsbetweenof thepolarization
globule, degree
but not
the innerregionFig.outer
9. Correlation
andin
normalized
intensity
for threeof radiation sources
the SOFIA/HAWC+
thedi↵erent
degreecasesof polarization
variesandbetween⇠ 0.5 and 5 %. The deobservationof B335 at 214µm.notGreen:
Radiationdeep
field based
only onThis
the ISRF;
Radiation field based only on the star; Red: Radiation field
becausethe radiation
does
penetrateenough.
is Purple:
rived
magnetic
field
strength
at
1.3
mm
is
higher (300-3000 µG)
the star highand ISRF;
Blue: depth
Best-fit forof ourtheobservation
Fig. 3).
due based
to anonoverly
opticalstellar (seecontribution
than our result. In analogy to the possible role of flux-freezing
to the ISRF ( . 5 µm, ⌧1µm,core ⇠ 120), which is important for discussed above, this finding may be explained by a strong in𝑄<3=> = 0.1 − 0.2、ISRF+Starモデルで、観測結果を比較的よく再現できるが、中心に向かって偏光度が下がるのは再現できない。
the alignmentof thegrains.
ISRF
cannotiii) Our model
consistsof aHowever,
mixture of thesilicate
andalonegraphitedust ex- shown
thisofe↵ectresult infielda decreasein the polarizawardthat
drag
the maymagneticlines (Mauryet al. 2018).
grainshomogeneously
distributed
over
the
Bok
globule.
For
tion
degreewithincreasing
incident
radiation
field, explainplain the drop in the degree of polarization because it produces

Constraining the magnetic ﬁeld properties of Bok globule B335 using SOFIA/HAWC+
Niko Zielinski, Sebastian Wolf, Robert Brunngräber

B335のその他の偏光観測
N. Zielinski: Magnetic field properties of the Bok globule B335
214 μm1.24 μm
10 000 au

1.24 μm1000 au
1000 au• 偏光ベクトルと磁場の向きが平行（dichroic absorption）• 磁場の強さは12-40 μGで、今回の結果（~142 μG）よりも小さい• ガス密度が低いため？
• 磁束凍結限界における磁束とガス密度の関係は要議論

870 μm1300 μm

870 μmと1.3 mm
200 au• ポロイダル磁場を観測• 214 μmでは分解能が低いため細かい構造は分解できない• 偏光度は1-11 %（ 214 μmでは0.5-5 %）
• 磁場強度は300-3000 μG（1.3 mm）→磁束が内側へ引き摺られたため？

Fig. 10. Polarization maps and corresponding magnetic field structure of B335. Top-left: Polarization map of B335 in the near-IR (Js band,
1.24 µm, overlaid on a DSS intensity map; Bertrang et al. 2014). Top-right: Polarization map of B335 at 214 µm (overlaid on the SOFIA/HAWC+
intensity map; this publication). Bottom-right: Polarization map of B335 at 1300 µm (overlaid on the ALMA polarized intensity map; reuse of
Figure 1 in Maury et al. 2018). Shown is the derived magnetic field direction. The arrows indicate the outflow direction. Bottom-left: Polarization
map of B335 at 870 µm (overlaid on the ALMA contour map; Yen et al. 2020, ©AAS. Reproduced with permission). Shown is the derived
magnetic field direction.

3D morphology. Such projection eﬀects can explain the variety of recently observed field configurations, in particular within asingle MC. Finally, we do not find a correlation between the observed orientation of B and the N -PDF.

Atlas of CO-Line Shells and Cavities around Galactic Supernova Remnants with FUGIN
Yoshiaki Sofue, Mikito Kohno, Tomofumi Umemoto

Atlas of CO-Line Shells and Cavities around Galactic Supernova Remnants with FUGIN
Yoshiaki Sofue, Mikito Kohno, Tomofumi Umemoto A morphological search for molecular shells and cavities was

performed around 63 Galactic supernova remnants (SNR) at 10  l  50 , |b|  1 using the FUGIN (FOREST
Unbiased Galactic Imaging survey with the Nobeyama 45-m telescope) CO line data at high angular (2000 ) and velocity(1.3 km s 1 ) resolutions. The results are presented as supplementary data for general purpose for investigations of
the interaction between SNRs and interstellar matter in the form of an atlas of CO-line maps superposed on radio
continuum maps at 20 cm along with a list of their kinematic distances determined from CO-line radial velocities.
(Full atlas including all figures is available in this URL: https://nro-fugin.github.io/2020-apjs-CO-Shell-Atlas-SNRFUGIN-IX.pdf)

An ALMA view of the Galactic super star cluster RCW38 at 270-AU resolution

Kazufumi Torii, Kazuki Tokuda, Kengo Tachihara, Toshikazu Onishi, Yasuo Fukui We report millimeter andsubmillimeter continuum and molecular line observations of the Galactic super star cluster RCW 38, obtained fromthe Atacama Large Millimeter/Submillimeter Array with a minimum angular resolution of 000 .17 ⇥ 000 .15 (' 289 AU ⇥255 AU). The C18 O image reveal many massive condensations embedded within filamentary structures extending
along the northwest-southeast direction in the center of cluster. The condensations have sizes of 0.01-0.02 pc, H2column densities of 1023 -1024 cm 2 , and H2 masses of 10-130 M . In addition, the 233-GHz continuum image reveals
two dense, small millimeter-sources with radii of 460 and 200 AU (Source A and Source B). Source A is embedded
within the most massive C18 O condensation, whereas no counterpart is seen for Source B. The masses of Source A
and Source B are estimated as 13 and 3 M at the optically-thin limit, respectively. The C18 O emission shows avelocity gradient of 2 km s 1 at the central 2000 AU of Source A, which could be interpreted as a Keplerian rotation
with a central mass of a few M or infall motion of gas. Further, the ALMA 12 CO data reveal that Source A andSource B are associated with molecular outflows exhibiting maximum velocities of ⇠30-70 km s 1 . The outflows haveshort dynamical timescales of <1000 yr and high mass outflow rates of ⇠ 10 4 -10 3 M yr 1 . These observationalsignatures suggest an early evolutionary phase of the massive star formation in Source A and Source B.

SEDIGISM-ATLASGAL: Dense Gas Fraction and Star Formation Eﬃciency Across the
Galactic Disk

J. S. Urquhart, C. Figura, J. R. Cross, M. R. A. Wells, T. J. T. Moore, D. J. Eden, S. E. Ragan, A. R. Pettitt,h"ps://arxiv.org/pdf/2011.11916v1.pdf
A. Duarte-Cabral, D. Colombo, F. Schuller, T. Csengeri, M. Mattern, H. Beuther, K. M.
Menten, F. Wyrowski,

The asymmetric inner disk of the Herbig Ae star HD 163296 in the eyes of VLTI/MATISSE:
evidence for a vortex?

The asymmetric inner diskMATISSEof the Herbig
Ae starHD et
163296in the eyes of VLTI/MATISSE:
Collaboration,
J. Vargaal.
evidence for a vortex?
MATISSE Collaboration, J. Varga et al. Context. The inner few au region of planet-forming disks is a complex

environment. High angular resolution observations have a key role in understanding the disk structure and the dynamical processes at work. Aims. In this study we aim to characterize the mid-infrared brightness distribution of
the inner disk of the young intermediate-mass star HD 163296, from VLTI/MATISSE observations. Methods. We
use geometric models to fit the data. Our models include a smoothed ring, a flat disk with inner cavity, and a 2D
Gaussian. The models can account for disk inclination and for azimuthal asymmetries as well. We also perform
numerical hydro-dynamical simulations of the inner edge of the disk. Results. Our modeling reveals a significant
brightness asymmetry in the L-band disk emission. The brightness maximum of the asymmetry is located at theNW part of the disk image, nearly at the position angle of the semimajor axis. The surface brightness ratio in theazimuthal variation is 3.5 ± 0.2. Comparing our result on the location of the asymmetry with other interferometricmeasurements, we confirm that the morphology of the r < 0.3 au disk region is time-variable. We propose that thisasymmetric structure, located in or near the inner rim of the dusty disk, orbits the star. For the physical origin of theasymmetry, we tested a hypothesis where a vortex is created by Rossby wave instability, and we find that a uniquelarge scale vortex may be compatible with our data. The half-light radius of the L-band emitting region is 0.33 ± 0.01
au, the inclination is 52 +57 , and the position angle is 143 ± 3 . Our models predict that a non-negligible fraction of
the L-band disk emission originates inside the dust sublimation radius for µm-sized grains. Refractory grains or large(& 10 µm-sized) grains could be the origin for this emission.

On the pumping of the CS(v = 0) masers in W51 e2eD. J. van der Walt, A. Ginsburg, C. Goddi We present the results of numerically solving the rate equations forthe first 31 rotational states of CS in the ground vibrational state to determine the conditions under which the
J=1-0, J=2-1 and J=3-2 transitions are inverted to produce maser emission. The essence of our results is that theCS(v = 0) masers are collisionally pumped and that, depending on the spectral energy distribution, dust emission can
suppress the masers. Apart from the J=1-0 and J=2-1 masers the calculations also show that the J=3-2 transition
can be inverted to produce maser emission. It is found that beaming is necessary to explain the observed brightness
temperatures of the recently discovered CS masers in W51 e2e. The model calculations suggest that a CS abundanceof a few times 10 5 and CS(v = 0) column densities of the order 1016 cm 2 are required for these masers. The rarityof the CS masers in high mass star forming regions might be the result of a required high CS abundance as well as
due to attenuation of the maser emission inside as well as outside of the hot core.

h"p://arxiv.org/abs/2012.05697v1

the L-band disk emission originates inside the dust sublimation radius for µm-sized grains. Refractory grains or large(& 10 µm-sized) grains could be the origin for this emission.

On the pumping of the CS(v = 0) masers in W51 e2eD. J. van der Walt, A. Ginsburg, C. Goddi

On the pumping of the CS(v = 0) masers in W51 e2eD. J. van der Walt, A. Ginsburg, C. Goddi We present the results of numerically solving the rate equations forthe first 31 rotational states of CS in the ground vibrational state to determine the conditions under which the
J=1-0, J=2-1 and J=3-2 transitions are inverted to produce maser emission. The essence of our results is that theCS(v = 0) masers are collisionally pumped and that, depending on the spectral energy distribution, dust emission can
suppress the masers. Apart from the J=1-0 and J=2-1 masers the calculations also show that the J=3-2 transition
can be inverted to produce maser emission. It is found that beaming is necessary to explain the observed brightness
temperatures of the recently discovered CS masers in W51 e2e. The model calculations suggest that a CS abundanceof a few times 10 5 and CS(v = 0) column densities of the order 1016 cm 2 are required for these masers. The rarityof the CS masers in high mass star forming regions might be the result of a required high CS abundance as well as
due to attenuation of the maser emission inside as well as outside of the hot core.

Constraining the magnetic field properties of Bok globule B335 using SOFIA/HAWC+
Niko Zielinski, Sebastian Wolf, Robert Brunngräber Thanks to their well-defined shape and mostly isolated loca-

tions, Bok globules are suitable objects for studying the physics of low-mass star formation. To study the magnetic
field of the prototypical Bok globule B335, we obtained a spatially resolved polarization map with SOFIA/HAWC+ at
a wavelength of 214 µm. For the first time, these observations reveal that polarization holes in Bok globules, that is,the decrease in polarization degree towards their dense centers, also occur in the far-infrared wavelength regime. The
observed polarization pattern is uniform with a mean polarization angle of 48 ±26 and a magnetic field strength of⇠ 142 µG. Moreover, we use complementary polarimetic data for B335 obtained at near-infrared to millimeter wavelengths to analyze and constrain the magnetic field across diﬀerent scales. By applying the 3D Monte-Carlo radiativetransfer code POLARIS (Reissl et al. 2016), we developed a model for the density and magnetic field structure aswell as for the dust properties of this globule. We conclude that the column density towards the center of B335 is too
low to cause the observed polarization hole in B335 via dichroic absorption (Brauer et al. 2016). Furthermore, we
conclude that the eﬀect of self-scattering has no significant impact on the observed polarization. Adopting dust-grainalignment via the radiative torque mechanism, a combination of the interstellar radiation field and the central star asradiation sources is consistent with the decrease in polarization degree at the outer regions of B335 (⇡ 104 au from thecore). However, the model fails to explain the low polarization degree within the inner 5000 au.
h"p://arxiv.org/abs/2011.10420v2

Search for associations
youngcontainingstars (SACY)
An updated census of
Searchcontaining
for associations
young VIII.
stars (SACY)
VIII.
An updated census
of spectroscopic
systems
showingof non-universalmultiplicity
multiplicity among
these associations
spectroscopic
binarysystems binary
showinghints
of hints
non-universal
among
these
associations
S. Zúñiga-Fernández et al.
S. Zúñiga-Fernández, A. Bayo, P. Elliott, C. Zamora, G. Corvalán, X. Haubois, J. M. Corral-Santana, J.
Olofsson, N. Huélamo, M. F. Sterzik, C. A. O. Torres, G. R. Quast, C. H. F. Melo We seek to update thespectroscopy binary fraction of the SACY (Search for Associations Containing Young stars) sample taking in consideration all possible biases in our identification of binary candidates, such as activity and rotation. Using high-resolution
spectroscopic observations we have produced ⇠1300 cross-correlation functions (CCFs) to disentangle the previously
mentioned sources of contamination. The radial velocity values obtained were cross-matched with the literature andwere used to revise and update the spectroscopic binary (SB) fraction in each of the SACY association. In order tobetter describe the CCF profile, we calculated a set of high-order cross-correlation features to determine the originof the variations in radial velocities. We identified 68 SB candidates from our sample of 410 objects. Our results
hint that the youngest associations have a higher SB fraction. Specifically, we found sensitivity-corrected SB fractions+16+9of 22+15
Pictoris, in contrast with the five oldest (⇠ 35 125
11% for ✏ Cha , 31 14% for TW Hya and 32 8% forMyr) associations we have sampled which are ⇠ 10% or lower. This result seems independent of the methodology
used to asses membership to the associations. The new CCF analysis, radial velocity estimates and SB candidates
are particularly relevant for membership revision of targets in young stellar associations. These targets would be ideal
candidates for follow-up campaigns using high-resolution techniques in order to confirm binarity, resolve the orbits,
and ideally calculate dynamical masses. Additionally, if the results on SB fraction in the youngest associations are
confirmed, it could hint of non-universal multiplicity among SACY associations.

h"p://arxiv.org/abs/2010.08575v2