THE STAR FORMATIONNEWSLETTER
No. 273 (#41-45)
塚越崇(茨城大)
09 Oct. 2015

1. Detection of two power-law tails in the probability distributionfunctions of massive GMCs
2. Star-forming regions a the periphery of the supershell surrounding
the Cyg OB1 association. I. The star cluster vdB130 and its ambient
gas and dust medium
3. The migration of gas giant planets in gravitationally unstable discs
4. Grain size limits derived from 3.6 micron and 4.5 micron coreshine
5. The Dense Filamentary Giant Molecular Cloud G23.0-0.4:
Birthplace of Ongoing Massive Star Formation

Detection of two power-law tails in the probability distribution functions of massive GMCs
N. Schneider et al.

• 3つの大質量星形成領域のprobability
distribution function (PDF)で2つのpowar-law slopeを発見• MonR2, NGC6334, CepOB3• Herschelによる連続波観測データ• slope成分に空間的な違いが見られる• slopeの解釈 (s∝-α-1(モデル依存), ρ~r-α)• 1st powar-law: free-fall的• 2nd powar-law: free-fallだけでは無理。collapseを遅くさせる必要あり。

• メカニズムについては追研究が必要• (1)回転構造, (2)光学的厚み, (3)磁場, (4)幾何学的構造, (5) stellar feedback

2 theSitnikCyg
et al. OB1 association. I. The
Star-forming regions a the periphery of the supershell surrounding
star cluster vdB130 and its ambient gas and dust medium
J. Steinacker et al.

• open cluster vdB130の多波長データ解析•

Cyg OB1 super-shellのNE part

複雑な空間構造を持つ•クラスターに付随する星の再定義、クラスターの星周環境との関連を調べる• 観測データ•

4

Sitnik et al.

SAO RAS 6m望遠鏡 2.3 Archival data
spectral obs. & narrow-band
image
schel space telescopes.
The(Ha,[SII])
region studied was observed
This work is partially based on data from Spitzer and Her-•by Spitzer space telescope within the framework of the‘A Spitzer Legacy Survey of the Cygnus-X Complex’ program (Hora et al. 2011). The images at 3.6 µm, 4.5 µm,
5.8 µm, 8.0 µm, and 24 µm were downloaded from the project
web-page1 at the NASA/IPAC Infrared Science Archive.
The Herschel archival data obtained within the frameworkof the ‘Hi-GAL: The Herschel Infrared Galactic Plane Survey’2 program (Molinari et al. 2010) were downloaded fromthe Herschel Science Archive3 .
In this study we also used images from the 2MASS
IR survey (Skrutskie et al. 2006) and stellar proper motions
data from the UCAC4 catalogue (Zacharias et al. 2013).

SAI MSU 125cm望遠鏡

長期Ha観測から視線速度の導出に使用•••

Spitzer & Herschel data archive
2MASS image

UCAC4 proper motion
catalog
3 OPENSTAR CLUSTER VDB 130
Initially designation ‘vdB 130’ has been applied to a reflection nebula associated with the BD+38 3993 star (HD228789, vdB 130a) of spectral class B1-B2 (van den Bergh
1966; Racine 1968). The reflection nebula vdB 130 can beseen on the optical maps shown in Figures 1a and 2. Racine• Proper motionで36天体を選定

E-DSS2 image?

Figure 1. (a) A north-western part of the supershell surrounding the Cyg OB1
map. Black rectangles show the cluster region (top) and the region studied inrectangle indicates the region for which the PV diagram is constructed (see Se
stellar winds: O7e HIP100173 (1), O8 HD193595 (2), and Of HD 228841 (3). (bSpitzer image. Overlaid are contours of 13 CO J=1-0 emission. The box limits thetriggered star formation within walls of this shell can help toreveal the spatially resolved star formation history in the entire region. In this series of papers we plan to present resultson the stellar content, young stellar objects, and the diﬀuse
matter in the area, putting them into a general context ofthe star formation in giant molecular cloud complexes.
In this work we analyze a small region in the northwestern wall of Cyg OB1 supershell. It contains the
young open cluster vdB 130 characterized by an anomalously strong and non-uniform reddening (Racine 1974;
Moﬀat & Schmidt-Kaler 1976; Le Duigou & Knödlseder2002), the reflection nebula vdB 130 of the same
name (van den Bergh 1966), also referred to as GN
20.16.0 (Magakian 2003), multiple compact and extended
infrared (IR) sources, and a molecular cloud of a cometary
shape seen in CO emission (cloud A in Schneider et al.
Figure2. Star cluster vdB 130. Stars from the original Racine
2007).
(1974) paper are marked with indices from 1r to 14r. Stars seThe goal of this study is to revise the stellar content

r付はRacine(1974)で観測されたもの
lected in this study based on their proper motions are markedof our reon the dproper m
for a nu
the strucinterstelpossibleulation a
the natu
the neigschel arc
main res10• カラーexcessと距離Sitnik et al.
24

クラスター年齢5-10Myr
15

23

10

0
50

50

22

22

N2

50
−150

−50

0

N4

0

50

150

100

−1
00

−50

100

0
−100

−50

−150−5

DEC (2000)

21

20

−100

0

DEC (2000)

0• クラスターは分子雲の’head’に付随

N3

−1
5

20

19

N1

−150

headでは速度にズレ=>stellar wind/super nova

21

0

−100

optical (gray)
13CO (contour)

19

•

23

0

100

• 可視光データ-3.3kpc, 赤外データ1.8kpc=>’通常の’extinction lawではない

150

•

24

(a)

39° 18′39° 18′Sitnik et al.
•10 領域の星間ガスは電離ガスとshockedガス

Figure 7. VdB 130 region with contours of the ‘head’of the CO
cloud (Schne
s
s

h
m
m
s
s
20
00s the 8.017
and18ontoµm50image (b). 40
RA (2000)

0.0

0
100

10

0
50

50

22

50
−150

−50

150

100
−1
00

N3

−1
50

−100

19

N1

−150

39° 18′

0

50

−50

20

100

0
−100

−50

DEC (2000)

N2

N4−150−5−10000

DEC (2000)

0

21

Ha - ionized gas
20h 18m 00s

17m 50s40s
RA (2000)

0.3[SII]/Ha - shock tracer

30s

0.4

0.5

0.00

evidence in favour of the physical connection between themolecular cloud and ionized hydrogen in the (b)cluster(c)out23skirt. The data exhibit single-component Hα profiles toward23
vdB 130 with peak line-of-sight velocities ranging from 0
to 15 km s−1 . These velocities characterize gas whose mo22
22
tions are mostly due to Galactic rotation and streaming
motions caused by spiral density waves. To analyse char5r
acteristic systematic
motions of ionized hydrogen in the
1r 1r
21
21
11r 11r 6r
6r
studied‘diagram of occurrence’of line5r region, we use
12r a12r
of-sight velocities (see, e.g. Arkhipova et al. 2013 and ref20
erences therein). Characteristic line-of-sight velocities of the20
ionized gas VLSR in the region are defined as average values at half-maximum of the corresponding diagram together
19
4rwith root-mean-squareerrors of their estimates. The ob19
tained interval of4rthe most commonly occurring ionized gas
velocities toward vdB 130 is ∆VLSR ∼ 3 − 9 km s−1 (based
39° 18′on 24 velocity measurements out of 26) and is shifted in
39° 18′
h
m
s
m
s
s
s
positive
relativeto velocities
of30COemission
h
mthe
s
s
s
20201850direction30
1800
00s 17m 5517
50s
4540
40s
35s
25s
−1
RA
(2000)VLSR ∼ −1 ÷ 2.5 km s RA
in(2000)
the head of the cloud (see Fig. 2

40

20s

23

0

22

Vlsr, km s-1

150

15

0.2

2

DEC (2000)

(a)

23

0.1

20

24

DEC (2000)

24

30

21

-40
20

19

-80
39° 18′h

20 18
Figurea• ダスト連続波放射22'広がったフィラメント - COに対応する構造あり

23'

•

21'
二つのBlob- COの’head’に対応

22'• Blob(を含む何箇所か)でF8/F24<1

2

Blo b W
Blo b EHa lp h a Blo b

21'

De c . (J 2 0 0 0 )

•

2

20'

•

9 °2 0 'HII+ 3regionを示す
+ 39°

19'

42s

•

20h17m 36s

8µm(+500µm)

埋もれたB型星に励起源だろう
18m 00s
57s
54s
51s
20h17m 48s+ 3 9 °1 8 'R.A. (J 2 0 0 0 )0)0)

Big Fila m e n t

De c . (J 2 0 0 0 )Ha lp h a Blo ba

24'

De c . (J 2 0 0 0 )ila m e n t

18m 00s

54s

48s

42s

20h17m 36s

1

c

2

R.A. (J 2 0 0 0 )

c

d

24'

24'

1r
11r

21'

12r

5r

20'

22'

2

21'

2

2

19'
4r

F8/F24+ 3 9 °1 8 '20h17m 36s

2

20'

19'

42s

23'

De c . (J 2 0 0 0 )22'De c . (J 2 0 0 0 )De c . (J 2 0 0 0 )23'

18m 00s

54s

48s

R.A. (J 2 0 0 0 )42s

20h17m 36s

24µm(+160µm)+ 3 9 °1 8 '+ 3 9 °1
18m 00s

54s

48s

R.A. (J 2 0 0 0 )42s

20h17m 36s

16 π σSB |r − rp |The migration of gas giant planets in gravitationally unstable discs
where Lp and rp , are the luminosity and position
ofoftheThe survival
gas giant planets
Dimitris Stamatellos
protoplanet,
respectively.gap Thisluminosity
isasgivenis more diﬃcultand takes twice
much timeby
tohappen (4 × 10 yr in comparison with 2 × 10 yr for thenon-radiative protoplanet). The gap not only takes more
• 軌道長半径の大きい惑星の形成メカニズム
time to open but it is also shallower and less wide (Fig4

•

GM Ṁ
,
L =f
R

4

p
p protoplaneture 3a), as the
moves closer to the centralp star (where both the disc scale height and Hill radius of
重力不安定円盤での形成が考えられるが、migrationの
the protoplanet
accare smaller). The protoplanet grows inmass slower, and the inward migration continues. The
protoplanet in this case migrates closer to the central
効果が大きく外側で保持出来ない問題がある
star (Figure 2a) but its mass growth is considerably supp
p the orbit of thepressed (Figure 2b). At the same time
protoplanet is circularized (Figure 2c).
accThe final outcome of disc-planet interactions is distinctly diﬀerent for the two cases presented here (Figure 1). In the case without radiative feedback the protoplanet eventually becomes a brown dwarf with a massof 28 MJ , on a eccentric orbit (e ≈ 0.17), and a semimajor axis of 50 AU. When irradiation from the accreting protoplanet is included in the model, the protoplanet reaches a mass of only 14 MJ , i .e. around the
3D SPH (SEREN) deuterium burning limit and within theplanetary-mass
regime (< 11 − 16.3 MJ ; Spiegel et al. 2011), while it rea relatively wide (semi-major axis of 15 AU),power-law Σ(R) & mains
T(R)onorbit,
circularavoiding excessive inward migration.
The role of the radiative feedback from the protoplanetis critical. Initially, irradiation from the protoplanet
3kyr後にprotoplanetを入れる
heats the disc making the opening up of the gap morediﬃcult, while thermal pressure delays large gas accretion onto the protoplanet. Therefore, the protoplanet
continues its inward migration without accreting excesset tosively. Its migration is eventually slowed acc
down once the ⊙"!
gap is opened −3/4
up but the protoplanet still continues toheat andRstabilize the disc. The eﬀective viscosity due to⋆total pseudo-ambient
10 andK , the inward(3)TA (r) = 250theK disc T:
self-gravity is+small
flow of gas isAUrelatively
slow. Most of the gas within
the protoplanet
4
4
4⋆ orbital radius is accreted
planetonto the central star creating aset
to AR is the distance
where
from
star
measured
onwhenthe radiative
hole inthe
the disc.
the other hand
A inner
A On
feedbacktheis notthetakenprointo account,"−3/4! fromdisc midplane. The contribution
to protoplanet
TA (r) fromthediscremains
unstable;theToomreparameter
Q atR
toplanet isT ⋆ (r) = 250 K
+
10
K
,
(3)
the
edges
of
the
gap
is
below
1
(see
Figure3b).
ThereA
AUgas flows rather "fore, thefast towards the outer edge of!
1/4
the gap, ontothecircumplanetary
disc and eventually
0.5
1
1.5
L
p
planetonto
the protoplanet,as the eﬀectivedue [時間依存あり]
towhere R isT the distancefrom
the star measured
onviscosity
the
(r)
=
,
(4)
-3
A
log
[columndensity
(g
cm
)]
2
the
disc
self-gravity
islarge.
The
gravitationallyunsta16
π
σ
|r
−
r
|
disc midplane. The contribution
SB to TAp(r) from the proble gap edges may also be responsible for the outward(5)

where M is the mass of the protoplanet, Ṁ is the accretion •rate
on to it, and R the accretion radius. f = 0.75
protoplanetへのガス降着とradiativeis the fraction of the accretion energy that is radiatedaway atfeedbackを入れて計算
the surface of the protoplanet, rather than being expended driving jets and/or winds (Machida et al.
2006). •As we assume that this planet has formed bygravitational instabilities in the disc, the accretion hap• the second hydrostatic core. The radius ofpens onto
the second• core is uncertain; it is estimated to be ∼ 1 −20 R⊙ (Masunaga & Inutsuka 2000; Tomida et al. 2013;
heating
Vaytetet al. 2013).
Here we shall assume R= 1R .
feedback• stellar
pseudo-ambient temperature is
r The
orbitaltotal•
et of ini#$#$
AU from
T (r) = T (r) + T
(r) .
(6)
oheating
evolve.

e as theorbitalWe
note that the radiative feedback from the star is fixedding
thet of inioplanet’s
whereasthe radiative feedback from the protoplanet isAU
fromentricityo variable
evolve.
with time and depends on the accretion• of gas
gas
frome as the23

• FBありモデルでの違い•

gapを開けるのに時間がかかる•

migrationが続くが原始惑星の成長がゆっくりなのでgapは浅く狭くなる

•

円軌道

•

最終の惑星質量が小さめ• 原始惑星を考慮するとself-regulatedな
効果を生む
•

•

4

原始惑星への質量降着

•

惑星質量↑

•

gapを開けやすくする

原始惑星からのradiative feedback•

周辺の円盤温度↑•

降着を妨げgapを開けづらくする=>inward migration+低惑星質量となる。
gapが開いた後はmigrationが止まり原
始惑星が生き残るFig. 2.— The evolution of (a) the semi-major axis α, (b) the
mass Mp , and (c) the eccentricity e of the protoplanet. The protoplanet initially migrates fast inwards, but it grows in mass and it
is able to open up a gap in the disc. The migration then is essentially halted. In the non-radiative case the protoplanet becomesa brown dwarf, whereas in the case with radiative feedback mass
growth is suppressed and the protoplanet’s mass remains within
the planetary-mass regime (marked by the horizontal dashed linesin (b)). Radiative feedback keeps the protoplanet on an circularorbit (e ≈ 0).

position is 0.17RH and 0.13RH for the simulations with
106 and 2 × 106 particles, respectively. We find that theplanet migrates inwards as expected with a migration
timescale of (1.7 − 2.6) × 104 yr when radiative transfer
is included (compared to 1.5 × 104 yr in the literature;
Ayliﬀe & Bate (2009)), and (4 − 7) × 104 yr for the local
isothermal case, i.e. where the disc temperature depends
only on the distance from the central star (compared to(9 − 11) × 104 yr and (6 − 12) × 104 yr in the literature;
Ayliﬀe & Bate (2009); Bate et al. (2003)). Therefore our
calculations are in very good agreement with previous estimates of migration timescales.
In the case of the massive, wide-orbit planets that we
examine here, the planet’s Hill radius is much larger andtherefore easier to resolve. For example, when the Jovian protoplanet is at 50 AU away from the central starits Hill radius is RH ∼ 3.5 AU. We use 106 SPH particles so that the smoothing length at the initial po-

Fig. 3.
the prot
amined hradial ce
the star,cell of t
shalloweplanet isaveragedc is thevelocity,star. Ra(Q > 1).

pletely2011) aThis
cretionpowereregulatstudies
Earth-lin fact(Benı́tegestedof the p
as it m2013).
planetradiusbeen suGrain size limits derived from 3.6 micron and 4.5 micron coreshine
J. Steinacker et al.
• コアでのMIR散乱は弱いと思われていたが、近年の観測でMIR散乱(coreshine)が見つかっている
Steinacker et al.: Grain size limits derived from coreshine•

ダスト成長を示唆• Spitzerの’Hunting for Coreshine’プログラムの観測データを使用•

10天体を選定•

一つ(RhoOph9)はPAHのコンタミ大• モデリングの仮定•

空間的に形状・サイズが不変•

power-lawサイズ分布•

ice-coated silicate & 炭素でMie theoryでFig. 1. Oﬀ-subtracted surface brightness images at 3.6 µm for the sources of the sample. The image sizes have been adjusted toshow the sources at the same distance in order to visualize the source sizes. Most cores have a complex shape, and the diﬀuse
signals diﬀer with respect to the peak brightness and point source contamination. In the case of Rho Oph 9, the image has not been
oﬀ-subtracted and the bright emission in the upper left corner is caused by transiently-heated particles and not by coreshine. The
crosses indicate the approximate location of the core center based on the corresponding 8 µm extinction maps or thermal emission
maps in the FIR/mm.

• 光学的厚みを考慮して放射輸送を解く• 密度構造(単純なもの)、ダストサイズを変化させ観測に合わせる
We have varied the power-law index and discuss its impact
on the findings in Sect. 4.1. While there are studies suggestingthat the spectral index of the distribution may be diﬀerent from
-3.5 as observed in the ISM (Weingartner & Draine 2001) or
deviate from power laws (Weidenschilling & Ruzmaikina 1994;
Ormel et al. 2009), we use this distribution as a typical casewhere the total grain surface area is dominated by the small

Numerical calculations by (Ormel et al. 2009) that are basedon aggregate collisions show that grains in cores go through aphase of compaction so that the spherical approximation might
be better than initially expected from a growth picture thatinvolves fluﬀy aggregates with large geometrical filling factors. Nevertheless, future modeling with a more complex dust
model should include non-spherical grains (as discussed, e.g., in3.6 µm
3.6
µm

A
A

0.12

0.14

surface brightness MJy/sr

0.10
0.08
0.06
0.04

0.10
0.08
0.06
0.04

D

0.5
0.4
0.3
0.2
0.1
0.0

0.6

4.5 µm
4.5
µm

0.5
0.4
0.3
0.2
0.1

Robs

# of pixels

3.0
200

100
1.5

0.5

F

0.10
0.08
0.06
0.04
0.02

3.6 µm
250

0
1.0

1.5

2.0

2.5
Robs

3.0

3.5

4.0

4.5 µm300
pixel #

350

400

5

H

300

1.0

6.2

0.00
200

L260

3.5

2.0

6.4

0.12

E

400

G
G

2.5

6.6

0.0

4.0

C
C

6.0

0.02

3.6 µm3.6µm

8.0 µm8.0µm

6.8

surface brightness MJy/sr

0.6

B
B

0.12

mean optical depth for extinction

mean optical depth for extinction

0.02

7.0

4.5 µm4.5µm

ratio R = Ics(3.6 µm)/Ics(4.5 µm)

surface brightness MJy/sr

0.14

Steinacker et al.: Grain size limits derived from coreshine

surface brightness MJy/sr• 結果例L260

4

I
Rtheo

3

Robs distribution
2

1
0.4

0.6
0.8 1.0 1.2 1.5
maximal grain size a2 [µm]

2.0

2.5

Fig. 5. Data analysis for L260. Warm Spitzer IRAC surface brightness maps (oﬀ-subtracted) at 3.6 and 4.5 µm (Panels A,B) and• 6天体では単一のamaxで各波長の再現に成功•

0.5~1.5µm••

ISMより成長
過去の研究とconsistent

Steinacker et al.: Grain size limits

6. S

L260
L1262 full RT
L1517A[e.g., Andersen+2013]

L1512 full RT• 3天体ではreasonableなamax無し

L1544 full RT
L1506C•

簡略化した仮定が原因?

Thi
mis
ten
and
emi
tiva
leng
ran
larg

ecc806

L1439
ISRF+local RF•

additional componentの存在?

L1498
0.2

•

0.3

0.4
0.6 0.8 1.0 1.2 1.5
maximal grain size a2 [µm]

2.0

3.0

Fig. 15. Summary of the derived maximum grain size ranges for
短波長でダスト性質を調べるとと the modeled sources. The black bar gives the range where Rtheo
is in the FWHM range of Robs , and the blue (green) bar showsrange for the model maximum coreshine to be within the obもに、現実的な密度分布を入れて theserved maximum coreshine range at 3.6 (4.5) µm, respectively.
For L1439, we also give size ranges for a model that includes aのfull-3D計算を行う必要がある local radiation field. Light green boxes indicate ranges of agreement.

ing
stru
that
tens
are
rati
den
mo
from
ties
a fe
size
exc
the
cal
is v

The Dense Filamentary Giant Molecular Cloud G23.0-0.4: Birthplace of Ongoing Massive Star
Formation
Yang Su et al.

• GMCのCO輝線広域マッピング観測•

大質量星形成の理解へ• G23.0-0.4 (d=4.4kpc)•

2 SNR+複数HII領域が付随•

フィラメント構造• 13.7mミリ波望遠鏡 (中国)•

CO,13CO,C18O (1-0)同時観測•

0’8 resolution– 24 –

Fig. 1.— 12 CO (J=1–0; black), 13 CO (J=1–0; blue, multiplied by a factor
C18 O (J=1–0; red, multiplied by a factor of ten) spectra of the 30′ × 30′ regio• 複雑な速度構造•~77km/s成分が付随する分子雲(PV図からも示唆)

CO, 13CO, C18O

• Tex分布:optically thick 12CO•

平均で20K•

SNR interface regionでは>26K• τ, Σ分布: 13CO/C18O
Fig.
3.— Map C18Oはthinof the excitation temperature with a linear scale in units of K from the• 13COはthick,−1
optically thick 12 CO (J=1–0)
emission
in
the
74–79
km
s
interval. The red rectangle-2-3• dense regionではΣ>5E22 cm (n>6E3 cm 相当)を示
shows the region of GMC G23.0−0.4.
唆=>star forming region• 速度構造にSNRとのinteractionあり

Fig. 3.— Map of the excitation temperature with
optically thick 12 CO (J=1–0) emission in the 74–79
SNR W41側
shows the region of GMC G23.0−0.4.

無関係な成分

G22.7-0.2側– 28 –

• 激しい星形成活動の証左•

少なくとも2つのSNR•

多数のHII領域•

多数の大質量星クラスター•

high density dust clumps• GMCの年齢の推定: 10-20Myr•– 28 –The black boxes, black circles, red circles, blue circles, andblue boxes show the fifteen dust-continuum-identified MC
clumps, the two 6.7 GHz methanol masers, the twelve H ii
regions, the two H ii/SNR complexes, and the two massive
star groups, respectively.

Fig. 8.— Left panel: Map of the abundance ratio X13 CO /XC18 O overlaid with the ATLASGAL 870 µm continuum emission contours (0.2, 1.2, 2.2, and 3.2 Jy/beam). The dense dust
clumps (>1.2 Jy/beam) are highlighted in red contours. Right panel: Map of the opticaldepth of the C18 O (J=1–0) emission. The black boxes, black circles, red circles, blue circles,
and blue boxes show the fifteen dust-continuum-identified MC clumps, the two 6.7 GHz
methanol masers, the twelve H ii regions, the two H ii/SNR complexes, and the two massive
star groups, respectively.

tSNR=tGMCを仮定

大質量GMCの1~3tffに相当
Fig.• 8.—
Left panel: Mapof the abundance ratio X13 CO /XC18 O overlaid with the ATLASGAL18870 µm continuum emission contours (0.2, 1.2, 2.2, and 3.2 Jy/beam). The dense dust
C Oフィラメントのビリアル解析•clumps(>1.2 Jy/beam) are highlighted in red contours. Right panel: Map of the opticaldepth of the C18 O (J=1–0) emission. The black boxes, black circles, red circles, blue circles,=>平衡状態and blue boxes show the fifteen dust-continuum-identified MC clumps, the two 6.7 GHz
methanol masers, the twelve H ii regions, the two H ii/SNR complexes, and the two massive•starクランプは等間隔分布:0.18º~0.26ºgroups, respectively.

=>フィラメント分裂の理論予想とコン
Fig. 9.— Map of the column density from C18 O (J=1–0) emission with a logarithmic scalein units of 1022 cm−2 . The four fragmentation points are marked with black circles.

システント (~17pc=~0.22º at 4.4kpc)