#41 Stellar energetic particle ionization in protoplanetary disks around
T Tauri stars
Ch. Rab,1 , M. Güdel1 , M. Padovani2 , I. Kamp3 , W.-F. Thi4 , P. Woitke5 , and G. Aresu61.07522v1 [astro-ph.SR] 24 Feb 2017

University of Vienna, Dept. of Astrophysics, Türkenschanzstr. 17, 1180 Wien, Austria
e-mail: christian.rab@univie.ac.at
Stellar energetic
particle(SP) がdisk
chemistryに与える影響を調べた
2
INAF-Ossevatorio
Astrofisico
di Arcetri, Largo
E. Fermi, 5 - 50125 Firenze, Italy
3
Kapteyn Astronomical Institute, University of Groningen, P.O. Box 800, 9700 AV Groningen, The Netherlands
4+とNGiessenbachstrasse+
Max-Planck-Institut für extraterrestrische Physik,
1, 85748 Garching, Germany
ディスクの電離トレーサーであるHCO
2H に焦点5
SUPA, School of Physics & Astronomy, University7of St. Andrews, North-2Haugh,
St. Andrews KY16 9SS, UK※ 典型的なT
Tauri型星のSP
fp ~ via
10dellaphotons
cmSelargius,s-1 @1auを仮定
6
INAF, Osservatorio
Astronomicoﬂux
di Cagliari,
Scienza 5, 09047
Italy※ SP電離度はH柱密度と半径の関数として算出（粒子は半径方向のみに動くとする）
Received 14 December 2016 / Accepted 24 February 2017
※ その他の電離源（X線・宇宙線）の変化も考慮。

ABSTRACT

Context. Anomalies in the abundance measurements of short lived radionuclides in meteorites indicate that the protosolar nebulae=> 他の電離源との相互作用を調べ、SP電離のトレーサーを同定する
Ch. Rab et al.: Stellar energetic particle ionization.
Ch. Rab et al.: Stellar energetic particle ionization.

was irradiated by a high amount of energetic particles (E & 10 MeV). The particle flux of the contemporary Sun cannot explain theseanomalies. However, similar to T Tauri stars the young Sun was more active and probably produced enough high energy particles to
explain those anomalies.
Aims. We want to study the interaction of stellar energetic particles with the gas component
CI of
XNthe
SP disk (i.e. ionization of molecular CL XH SP
hydrogen) and identify possible observational tracers of this interaction.
① SPsはディスクを突き抜けることができない
Methods. We use a 2D radiation thermo-chemical protoplanetary disk code to model a disk representative for T Tauri stars. We use aparticleenergydistribution derived from solar flare observations and an enhanced stellar particle flux proposed for T Tauri stars. For
25 cm
-2：
N<H> > 10
this particle spectrum we calculate the stellar particle ionization rate throughout the disk with an accurate particle transport model.
We study the impact of stellar particles for models with varying X-ray and cosmic-ray ionization rates.
たいていのエネルギー粒子は止まり、Results. We find that stellar particle ionization has a significant
on⇣ astheabundances
theN common
ionization
Fig. 6.impact
Ionization rates
a functionof vertical columnof
density
at radii of 1 disk
and 100 au(solid and dashedtracerslines respectively). The
maximum values for N
at the midplane of⇡ 4 ⇥ 10 cm and N⇡ 2 ⇥ 10 cm at 1 au and 100 au, respectively.
19the disk,
1 are N
HCO+ and N2 H+ , especially in models with low cosmic-ray ionization
rates
(e.g.
10
s
for
molecularhydrogen).
In
contrast
SP電離度は急速に落ちる
Red lines are for X-rays, blue lines are for SPs and the black lines are for CRs. Left panel: model CI_XN_SP with ISM CRs andto
normal X-rays;
right panel: model CL_XH_SP with low CRs and high X-rays.
cosmic rays and X-rays, stellar particles cannot reach the midplaneof the disk. Therefore molecular ions residing in the disk surface
and dashed lines respectively). The
25 Fig. 6. Ionization rates ⇣ as a function of vertical column density N <H>,ver at25 radii2 of 1 and 100 au (solid
232
Tタウリ型星の半径方向柱密度はN>> 10than<H>ionization
layers are more a↵ected by stellar particle
molecular
tracing
theofcoldlayers/midplane
the
disk.
maximum
values forions
N
atSP
the midplanethe disk,
are N
N <H>,ver
⇡ 2CI
⇥ 10
<H>,ver
CI
XN
CI<H>,ver
XH ⇡SP4 ⇥ 10 cmofand
T cm at 1 au and 100 au, respectively.
Red lines are for X-rays, blue lines are for SPs and the black lines are for CRs. Left panel: model CI_XN_SP with ISM CRs and normal X-rays;
Conclusions. Spatially resolved -1
observations of molecularions tracing di↵erent vertical layers of the disk allow to disentangle theright panel: model CL_XH_SP with low CRs and high X-rays.
cm-2なので、ζ
10-20
s @1aucontributionSPof <<stellarparticle ionizationfrom other competing ionization sources. Modeling such observations with a model like theone presented here allows to constrain the stellar particle flux in disks around T Tauri stars.
<H>,ver<H>,ver<H>,ver

25

2<H>,ver23

2

Key words. Stars: formation - Stars: circumstellar matter - Stars: activity - Radiative transfer - Astrochemistry - Methods: numerical② SPsはH2の主要な電離源となる＠CO ice line
上の暖かい分子層1. Introduction

Our Sun acts as a particle accelerator and produces energetic

CL XN SPCL XH SP

CL T

assumptions T Tauri stars might show on average a continuous
proton flux of fp (Ep
10 MeV) ⇡ 107 protons cm 2 s 1 at a
distance of 1 au from the star.

#41の続き

Ch. Rab et al.: Stellar energetic particle ionization.

CR電離度が低いモデルではさらに重要④ SP電離はTタウリ型星の出す値より１桁小さくなると重要で

HCO+

HCO+

N2 H+

N2 H+

NN2 H+ [cm 2 ]はなくなる
その場合、X線とCRによってH2は電離される

low CR

NHCO+ [cm 2 ]③ SP電離はHCO+やN2H+の柱密度を3-10倍増やす＠r<200au

ISM CR⑤ SPsがディスクの深層まで届かないとき、X線とCRが主要な
H2電離源となる

Fig. 9. Vertical column density profiles for HCO+ and N2 H+ for our model series (Table 5). The left column shows the models with the ISM likeCRs (⇣CR ⇡ 2 ⇥ 10 17 s 1 ) the right column with low CRs (⇣CR ⇡ 2 ⇥ 10 19 s 1 ). The top row shows HCO+ , the bottom row N2 H+ . The blues lines
are for models with, the red lines are for models without SPs. Dashed (solid) lines are for models with high (normal) X-rays. The orange solid line
shows the Turner model. The gray shaded area marks a di↵erence of a factor 3 in the column densities relative to the CI_XN (ISM CR, normal
X-rays) and CL_XN model (low CR, normal X-rays), respectively.

cold molecular layer and the dip in the profile vanishes.
=> HCO+はSPsに対してより影響を受けやすくなる（N2Hthe+よりも）tion rate is too low to produce a significant amount of HCO+ in

N2 H+ traces the distribution of gas phase CO as it is efficiently destroyed by CO (Qi et al. 2013a,b, 2015). This is also
seen in our model. The sharp transition in the N2 H+ column density at r ⇡ 30 au traces the onset of CO freeze-out (Fig. 9). We
note however, that the actual midplane CO ice-line is at ⇡ 12 au(see Appendix C for details). In the ISM CR models N N2 H+ isdominated by CR ionization as N2 H+ mainly resides in the cold
molecular layer. In the warm N2 H+ layer X-ray ionization dominates. However, due to the lower densities in the warm molecular layer, this layer only contributes significantly to the column
density within the radial CO ice line, and the impact of X-rays
is only visible there (compare models CI_XH and CI_XN in
Fig. 9). The high X-ray luminosity decreases the contrast between the peak of N N2 H+ close to the radial CO ice line and N N2 H+
inside the CO ice line by about a factor of five.

3.3.3. Impact of SPs

The column densities N HCO+ and N N2 H+ for models with SPs areshown in Fig. 9 (blue solid and dashed lines). The solid orange
line in Fig. 9 shows the results for the Turner model. We discuss
the Turner model separately in Sect. 3.3.4. We define a change
in the column densities by at least a factor of three compared to
the reference model as significant. This is indicated by the grayarea around the column density profiles of the reference models
CI_XN and CL_XN.
To better quantify the impact of SPs compared to the competing H2 ionization sources , X-rays and CRs, we introduce theweighted column density
Z 1
N⇣ (r) =n(r, z) ⇥ f⇣ (r, z) dz [cm 2 ].
(5)

⑥ ディスクの様々な鉛直層をトレースする分子イオンの強度プロファイルのモデリングにより、競合する高エネルギー電離源の寄与を解くことができる=> そのような観測は、ディスクでのSP ﬂuxに制約与える

0

N⇣ is the weighted column density for a particular H2 ionization
source ⇣, n is the number density of a particular molecule in unitsof cm 3 and⇣
この方法はモデル依存なので、ディスクの鉛直化学構造を制約する補助観測が必要
f =.
⇣ +⇣ +⇣⇣=> SPはディスク化学において重要

For the low CR case N N2 H+ drops by more than an order ofmagnitude compared to the ISM CR case. Such a strong impact of CR ionization on N N2 H+ is also reported by Aikawa et al.
(2015) and Cleeves et al. (2014). Higher X-ray luminosities cancompensate for low CR ionization only to some extent. In thehigh X-ray model CL_XH, N N2 H+ is lower by a factor of fivecompared to the ISM CR models.

X

SP

CR

(6)

N⇣ represents the fraction of the column density dominated by aparticular ionization source ⇣.
In Fig. 10 we show N⇣ for HCO+ and N2 H+ normalized tothe total column density of the respective molecule as a functionof radius. The gray shaded area in each plot roughly marks theregion where more than 50% of the total column density of the
Article number, page 9 of 20

SPがどの程度ディスクに届くのか解明するには、磁場のようなSP輸送や、
非恒星起源のjetなどの高エネルギー粒子を考える必要がある

10”$

The Dusty Silhouette Jet HH 1019 in the Carina Nebula
1?
2
2
3
Megan
Reiter,Megan
M.
Kiminki,Nathan
Smith
,
and
John
Bally
1
University of Michigan, 1085 S. University Ave., Ann Arbor, MI 48109, USA

2 Steward Observatory, 933 N. Cherry Ave., Tucson, AZ 85721, USA
3 Center for Astrophysics and Space Astronomy, University of Colorado, 389 UCB, Boulder, CO 80309, USA2

Tr 14星団付近のHST画像から、新たなHH 1019を発見

Reiter et al.

HαAccepted 2017 February 10. Received 2017 February 10; in original form 2016 November 41702.03893v1 [astro-ph.SR] 13 Feb 2017$1019$

#42

Tr14%

Hα
[O$III]

dark dust pillarの先端からbipolar ﬂow
We report the discovery in Hubble Space Telescope (HST) images of the new
dark%pillar%
7つほどのknotsがシルエットとして見える
Herbig-Haro jet, HH 1019,
located near the Tr 14 cluster in the Carina Nebula. Like

[O$III][S$II]
dusty%knots% [S$II]
ABSTRACT
star%

(サイズ<
other HH jets in the region, this bipolar
collimated0.5ʼʼ)flow emerges from the head of adark dust pillar. However, HH 1019 is unique because – unlike all other HH jets known
to date – it is identified by a linear chain of dark, dusty knots that are seen primarily
unseen%in silhouetteagainst the background screen of the H ii region. Propermotions
confirm
4 Reiter
et al.
driving%
that
these dark condensations move along the jet axis at high speed. [S ii] emission
source%traces a highly collimated jet that is spatially coincident with these dust
A knots. The
high extinction in the body of the jet suggests that this outflow has lifted a large
HH$1019$amount of dust directly from the disk, although it is possible that it has entrained
dustemission%from%from its surrounding protostellar envelope before exiting the dust pillar.
If dustB
η%Car%
in HH
1019
originatesfromthecircumstellar
disk,
this
providesfurther
evidence
for aweakly%irradiated%10”$
jet launchedgas%in%the%jet?%from a range of radii in the disk, including those outside the dust sublimation
radius. HH 1019 may be the prototype for a new subclass of dusty HH objects
Figure 1. Left: HST /ACS H↵ image of HH 1019 that shows the chain of dusty knots that make up the jet seen in silhouette against theseen
primarily
intheextinction
bright background of the H ii region. Middle: [S ii] emission clearly traces a narrow,collimatedjet even thoughdetailed structure isagainst the background screen of a bright H ii region.
unresolved in the seeing-limited ground-based image. Right: Cartoon showing the structure of HH 1019. Dark, dusty knots (black) tracejets
may
beweaklycommon,difficult to observe because they require the special
out a collimated bipolar jet emerging from an opaque dust pillar. Faint H↵ and [SSuchii] emission(gray)
trace theirradiated gas but
inthe jet.
condition of a very bright background in order to be seen in silhouette.
C[S II]禁制線： dust knotsと一致

Key words: stars: formation — stars: pre-main-sequence — ISM: Herbig-Haro ob6天体で固有運動の測定（2005年と2015年のHα画像）:
jects — ISM: jets and outflows

cooling zones. However, few comparable examples of dust in
These sources are associated with CO outflows and young
driving sources, leading Chini et al. (2001) to suggest thatthe jet body exist.
Ddust is heated in the interaction between the jet and theIn this paper, we present the discovery of another memenvironment. Extended (sub)-mm continuum emission that
ber of this rarefied group of jets with direct evidence forEtraces the bipolar outflow lobes has been detected in thedust. Unlike other Herbig-Haro (HH) jets, HH 1019 in the
L 1157 molecular outflow (Shirley et al. 2000; Gueth et al.
Carina Nebula appears to be composed of a chain of dark,2003). Many molecular lines contribute to the flux observed
dusty knots seen in silhouette against the background screen
HH 1019
F
at 1.3 mm and 850 µm, however Gueth et al. (2003) concludeof the H ii region (see Figure 1). UV photons from nearby
200 km s-1
that line emission can only
for . 50% of the continTr14 and Tr16 illuminate the edges of dusty knots. This is sources, higher luminosities may evaporate grains near the
1 account
INTRODUCTION
5"
uum flux. In addition, the spectral index of the 1.3 mm anddistinct from the illumination of the body observed in most
850 µm emission is consistent with thermal dust emission.
of the other 40 HH jets discovered in the same H↵ imaging star so that material launched from the inner disk may be
Protostellar
jets
accompany
activedisk
accretion
onto
formHowever, these observations have an angular resolution ofsurvey of Carina (Smith et al. 2010a). The high mass-loss dust-free (e.g. Kama et al. 2009). Figure 3. H↵ image of HH 1019 with boxes marking the jet knots
00
used to measure proper motions. Vectors indicate the magnitude
several arcsec (in the caseof Lstars.
1157, 12Whilecorresponding
torate physics
inferred for thosejetslaunch
implies intermediate-massing
the detailedof jetand colli- driving
and direction
the motionof the5280 AU a distance of 440 pc, Gueth et al. 2003), making
sources, suggesting that the HH jets in Carina o↵er a uniqueDustinprotostellar
jets,
especially
inofthebodyofsilhouette
the knots.
mationremainuncertain,
accretionenergythein the forit unclear whether the dustin the outflow
is entrained
from
view
of the underlying
jets mustthat maypower
be common

ジェット軸に沿って高速移動（~50-100 km s-1）
molecular outﬂowよりも速い=> knotsはjetをトレース

wide-angle outﬂow

narrowband
is a forbidd
increased ab
associatedthe knots aview, we idof the pillar
jet. These t
H↵ image.
limb of the(features I[S ii] e
Figure 1). [Sseen in silh
the seeing-lof the seeinof the HST
extends at lseen in silho
from the pi
irradiated j
the [S ii] mH↵, both tr
and H↵ ma
Tr16, perhaWith tmeasure theknots clearlcontrast wi
and cannot
proper mot

HH 1019 D
HH 1019 E
HH 1019 F
HH 1019 G
HH 1019 H
HH 1019 I
HH 1019 J#42の続き-1 (3)-15 (3)-26 (3)
...
...
...
...

40 (3)
74 (3)
102 (3)
...
...
...
...

43 (3)
83 (3)
115 (3)
...
...
...
...

2912 (812)1842 (463)1841 (328)
...
...
...
...

0.1
0.1
0.1
0.1
0.05
0.1
0.1

-5.1
-4.6
-4.1
-4.8
-5.4
-4.6
-5.2

The Silhouette Jet HH
H_[O III]

1.0
0.9

J

1.1
1.0

normalized intensity

0.9

Table 1. HH 1019 knot proper motions

C

0.9

C

1.1

0.9

D

D
E

1.1
1.0

A

0.9

E

1.1

1.1

1.0
0.9

1.0

F
G

1.0

B

0.9

F

H

1.1

0.9

Knot
name

1.0

G
B

1.0

1.1
1.0

J
IA

I

1.1

0.9

H_
[O III]

1.1
1.0[O III]0.9

H-0.6 -0.4 -0.2 0.0 0.2 0.4 0.6
position [arcsec]

normalized intensity

1.1

HH 1019 A
HH 1019 B
HH 1019 C
HH 1019 D
HH 1019 E
HH 1019 F
HH 1019 G
HH 1019 H
HH 1019 I
HH 1019 J

x
mas

y
mas

38 (3)0 (3)3 (3)-1 (3)-15 (3)-26 (3)
...
...
...
...

-86 (3)-45 (3)90 (3)
40 (3)
74 (3)
102 (3)
...
...
...
...

age
[yr]

⌧

log(Mknot )M

103 (3)50 (3)
98 (4)
43 (3)
83 (3)
115 (3)
...
...
...
...

1754 (377)2030 (753)532 (423)
2912 (812)1842 (463)1841 (328)
...
...
...
...

0.2
0.03
0.2
0.1
0.1
0.1
0.1
0.05
0.1
0.1

-4.4
-6.4
-4.8
-5.1
-4.6
-4.1
-4.8
-5.4
-4.6
-5.2

-0.6 -0.4 -0.2 0.0 0.2 0.4 0.6
position [arcsec]

Figure 4. An HST /ACS [O iii] image (center) shows dusty jets knots with lines and labels indicating the locationof theH_intensity
1.1
tracings plotted to the right and left. The H↵ intensity is plotted with a solid line while [O iii] is shown with a dashed line. Absorption[O III]
dips are deeper in [O iii] than H↵ due to emission from the ionization front seen at the edge of the jet (negative positions in the plots).
1.0[O III]禁制線の強度からdust質量を算出

Indirect evidence for dust in jets points to an origin in a
dusty disk wind. For example, Agra-Amboage et al. (2011)
infer the presence of dust in the DG Tau jet from Fe depletion. In this case, depletion is higher in the lower velocity
components that surround the fast spine of the jet (see also
Podio et al. 2011). For a disk wind, the outflow terminal
velocity is related to the Keplerian velocity at the launch
point in the disk, suggesting that slower jet material originates from larger radii. Disk rotation studies also place theorigin point of these slower, outer jet layers at larger disk
radii (see, e.g. Ferreira et al. 2006) that are more likely toreach outside the dust sublimation radius. Evidence for asimilar, onion-like velocity structure has been observed in a
few jets in Carina (Reiter et al. 2015b,a), as expected fora jet launched from a range of disk radii. However, unlike

vT
1]

[km s

DG Tau, HH 1019 emerges from a deeply embedded proto0.9
star. It is therefore possible that the jet entrained dust fromJ
the circumstellar envelope (see, e.g. Shadmehri & Downes
2008; O↵ner & Arce 2014).
1.1
Bright gas-phase [Fe ii], as seen in many of the jets
in Carina (Reiter et al. 2016), may indicate
little depletion
1.0
onto grains (e.g. Ellerbroek et al. 2014). However, HH 1019
provides evidence that dust survives in at least one jet in Carina. Multi-wavelength observations of0.9
other jets in Carina
I and ionizationsuggest the coexistence of multiple excitation
states, similar to the jets in Orion where
Smith et al. (2005b)
1.1
detect thermal dust emission. The physical conditions in thejet must allow dust in HH 1019 to survive processing throughshocks and external photoevaporation1.0
in the H ii region.
To estimate the dust mass of the silhouette knots, we

H_
[O II

J
IA

C

MNRAS 000, 1–?? (2016)

normalized intensity※前景放射の補正をしていないので下限値
・星周ディスクに対するOʼdell & Wen 94の手法でτを柱密度にG
・典型的なknot質量 ~10-5 Msun
B
0.9

C

D

D
E

A

E

1.1

HH 1019はHH天体の新しいサブクラスの典型になるかも
1.0
0.9

このようなjetは共通して起こるのかもしれないが、
G

1.1

非常に明るい背景が必要なのでシルエットで見るのは難しいだろう[O III]

F
H

F

1.0
0.9

B

H

-0.6 -0.4 -0.2 0.0 0.2 0.4 0.6
position [arcsec]-0.6 -0.4 -0
po#43 Vertical structure of the transition zone from infalling rotating envelope
to disc in the Class 0 protostar, IRAS 04368+2557

Nami Sakai,1‹ Yoko Oya,2 Aya E. Higuchi,1 Yuri Aikawa,3 Tomoyuki Hanawa,4Cecilia Ceccarelli,5 Bertrand Lefloch,5 Ana López-Sepulcre,2,6 Yoshimasa Watanabe,2Takeshi Sakai,7 Tomoya Hirota,8 Emmanuel Caux,9,10 Charlotte Vastel,9,10
Claudine Kahane5 and Satoshi Yamamoto2
1 The Institute of Physical and Chemical Research (RIKEN), 2-1 Hirosawa, Wako-shi, Saitama 351-0198, Japan

ALMA望遠鏡を用いた高分解能観測により、

2 Department of Physics, The University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan

3 Center for Computational Science, University of Tsukuba, Tsukuba, Ibaraki 305-8577, Japan
4 Center for Frontier Science, Chiba University, Chiba 263-8522, Japan. Sakai et al.

L78

L1527のClass 0天体の垂直方向構造をはじめて明らかにした5 Laboratoire d’Astrophysique de Grenoble, Universite Joseph-Fourier, BP 53X 5 rue des Moutonnées, Grenoble Cedex 9 F-38042, France

N. Sakai et al.

6 Institut de Radioastronomie Millimétrique, 300 rue de la Piscine, Domaine Universitaire, F-38406 Saint Martin d’Héres, France

7 Department of Communication Engineering and Informatics, Graduate School of Informatics and Engineering, The University of Electro-Communications,

Chofugaoka, Chofu, Tokyo 182-8585, Japan

8 National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan

天体：IRAS 04368+2557 in L1527

9 Universite de Toulouse, UPS-OMP, IRAP, Toulouse, France

10 CNRS, IRAP, 9 Av. Colonel Roche, BP 44346, F-31028 Toulouse Cedex 4, France

ほぼedge onの円盤状構造（左図：0.8 mm dust continuum map）

Accepted 2017 January 4. Received 2017 January 3; in original form 2016 October 4

ABSTRACT

We have resolved for the first time the radial and vertical structures of the almost edge-on
envelope/disc system of the low-mass Class 0 protostar L1527. For that, we have used Atacama
Large Millimetre/submillimetre Array (ALMA) observations with a spatial resolution of 0.25
× 0.13 arcsec2 and 0.37 × 0.23 arcsec2 at 0.8 and 1.2 mm, respectively. The L1527 dust(c) CCH積分強度分布
continuum emission has a deconvolved size of 78 × 21 au2 , and shows a flared disc-likestructure. A thin infalling-rotating envelope is seen in the CCH emission outward of about
150 au, and its thickness is increased by a factor of 2 inward of it. This radius lies between
the centrifugal radius (200 au) and the centrifugal barrier of the infalling-rotating envelope(100 au). The gas stagnates in front of the centrifugal barrier and moves towards verticaldirections. SO emission is concentrated around and inside the centrifugal barrier. The rotation
speed of the SO emitting gas is found to be decelerated around the centrifugal barrier. A partof the angular momentum could be extracted by the gas that moves away from the mid-plane
Figure 1. (a) 0.8 mm dust continuum map of L1527. Contours are 10σ , 20σ , 40σ , 80σ , 160σ , 320σ and 640σ , where 1σ is 0.09 mJy beam−1 ; (b) deconvolved
aroundfit theIf thisis the(c,d)
case, the centrifugal barrier would be related to theisc thickness as observed in the dust continuum map. The widths are obtained by the Gaussianof thecentrifugal
intensity profilebarrier.
along east–west
direction.
ntegrated intensity distributions of CCH (N = 4 − 3, J = 9/2 − 5/2, F = 5 − 4 and 4 − 3:colour) and SOmechanism(JN = 87 − 76 : colour)
superposed on theoutflows,
0.8 mmlaunchingof low-velocity
such as disc winds.

infalling-rotating envelope (IRE)が150auより外側で見え、内側では２倍ほどの厚さになっている
遠心力バリアの内側ではCCH emissionがない

ust continuum map (contours: 10σ , 40σ , 160σ , 320σ and 640σ ). The IRE traced by CCH is broadened inward of the radius of about 150 au.

Fig. 2(a) shows the velocity channel maps of CCH. In the "V =.8–1.1 and 1.2–1.5 km s−1 panels, the CCH distribution essentiallyhows a ‘ Y’-shaped feature, although significant east–west asymmetry is seen. The branching point of the ‘ Y’ shape corresponds tohe centrifugal barrier. The "V = 1.2–1.5
s−1 component
1 I NkmT RO
D U C T ItracesON−1

Key words: methods: observational – stars: formation – stars: low-mass – ISM: individual ob-observed PV diagram. Note that the two hyperfine components arejects: L1527– ISM:
kinematics
and dynamics
partly blended,
which makes
the PV
diagram complicated.
We also – ISM: molecules.
observed the N = 3–2 transition of CCH and the J = 5–4 transitionof CS with ALMA. Their behaviour is essentially similar to that ofanalysing
the N = 4–3 transition of CCH, although the angularresolutionthe
is IRE kinematics with those lines (Sakai et al. 2014a,b).

#43の続き(d) SOの積分強度図
IREがない領域で光っている
Resolving the centrifugal barrier in L1527(c) mid-planeに沿ったSOの位置速度図

Figure 1. (a) 0.8 mm dust continuum map of L1527. Contours are 10σ , 20σ , 40σ , 80σ , 160σ , 320σ and 640σ , where 1σ is 0.09 mJy beam−1 ; (b) ddisc thickness as observed in the dust continuum map. The widths are obtained by the Gaussian fit of the intensity profile along east–west dire
Integrated intensity distributions of CCH (N = 4 − 3, J = 9/2 − 5/2, F = 5 − 4 and 4 − 3: colour) and SO (JN = 87 − 76 : colour) superposed on
dust continuum map (contours: 10σ , 40σ , 160σ , 320σ and 640σ ). The IRE traced by CCH is broadened inward of the radius of about 150 au.

白点線：IRE modelFig. 2(a) shows thevelocity channel maps of CCH. In the "V
PV diagram. Note that the two hyperfine comp
黒実線：Keplerianmotion（0.18
M= sunobserved）の速度場
0.8–1.1 and 1.2–1.5 km s panels, the CCH distribution essentiallypartly blended, which makes the PV diagram complicated−1

shows a ‘ Y’-shaped feature, although significant east–west asymobserved the N = 3–2 transition of CCH and the J = 5–4
metry is seen. The branching point of the ‘ Y’ shape corresponds toof CS with ALMA. Their behaviour is essentially similar
the centrifugal barrier. The "V = 1.2–1.5 km s−1 component traces
the N = 4–3 transition of CCH, although the angular res
the envelope closer to the protostar than the "V = 0.8–1.1 km s−1
slightly lower.
component. On the other hand, the "V = 1.6–1.9 km s−1 compoHere, we note that the distribution of the CCH emiss
nent, which corresponds to the rotation velocity at the centrifugal
affected by self- and dust-absorptions. On the other hanbarrier (1.8 km s−1 ), shows a ‘ U’-shaped feature within a radius ofself-absorption feature due to foreground gas is seen at the
150 au. No significant emission is seen for "V of 2.0 km s−1 orvelocity (Fig. 2b), the absorption feature towards the pro
higher, which exceeds the rotation velocity at the centrifugal barsition in the moment 0 map (Fig. 1c) is likely due to this frier. This is consistent with the distribution of the CCH emission,
gas too. However, the absence of CCH emission inside thwhich is absent inside the centrifugal barrier. In these velocity changal barrier is not ascribed to the self-absorption from the fnel maps as well as the moment 0 map (Fig. 1), it looks as if a partgas. If CCH were present in the mid-plane inside the cof the CCH emission is distributed in the surface area on the discbarrier, its emission should be Doppler-shifted due to th
like structure traced by the dust continuum emission. However, thismotion around the protostar to escape from the self-absor
would not be the case, considering the above velocity structure.
dust opacity at 345 GHz is evaluated to be less than 0.1
This component most likely represents a projection of the ring-likeassuming the temperature model for this source reported
broadened distribution of the IRE around the centrifugal barrier.
et al. (2013). Hence, the absence of the CCH emission
Inside the centrifugal barrier, the density is 108 cm−3 or higher.
to the high opacity of the dust continuum emission, eithe
Since CCH is chemically very reactive, it could be lost throughthe SO line emission distribution is seen even in the close
reactions.panels)
Alternatively, CCH could be depleted on to the(∼10 au) of the protostar, as described later in Section 3.
CH (N = 4 − 3 J = 9/2 − 5/2 F = 5 − 4: upper panels; N = 4 − 3 J = 9/2 − 5/2 F = 4gas-phase− 3: lower
dust grains. Besides, the production of CCH is ineffective in highocity range described at bottom in the panel. (b) PV diagram of the CCH lines (colour) superposedto thedensity gas, becauseof aIRE
low ionization degree (i.e. a low abundance
3.3 SO+).
ofCe every 20 per cent of the peak intensity. (c) PV diagram of the SO line (JN = 87 − 76 : colour) superposed
Fig. 2(b) shows the position velocity (PV) diagram of CCH along
Fig. 1(d) shows the moment 0 map of the SO line. The SO
he Keplerian model (black contours). Note that the small undulation of the lowest contourof the overlaid
Keplerian
the mid-planeto the IRE model (Sakai et al. 2014b; Oyashows complex clumpy structures. The total flux is 2.0
etwithal. 2015).
IRE ofmodel
assumes
It is panels;
20 per Ncent
inFthe0 obtefact due to a limited mesh size of the model. In addition to the componentFigure
associatedthe Thecentrifugal
2. (a) Integrated
intensity
maps
the CCH(N =a4constant− 3 J = molecular
9/2 − 5/2 abunF = 5 − 4: upper= 4 lower− 3 J =than
9/2 that− 5/2= 4Cycle− 3: lower
dance,
anddoes
not consider
excitation
effects.
Nevertheless,theresolving-out
effect could
thus be(2.5
Jy km
s−1 ). The
line.
Each
map
is
obtainedbyintegratingthevelocityrange
described
atbottominthepanel.
(b)PV
diagram
of
the
CCH
lines
(colour)superposedtoeen (pink dotted lines). Yellow dotted arrows indicate the voids of the emission. (d) Schematic
illustrationofmodel well
the characteristic
behaviourthenot diagram
serious. of the SO line (JN = 87 − 76 : colour) supmodel (dotted contours).
Thereproduces
model contours
are every 20kinematic
per cent of
the peakofintensity.
(c)is PV・100auの遠心力バリア付近に低速度成分が付随
この回転速度はIREモデルの回転速度よりも遅い
・高速度成分はKeplerian motionと一致し、
inner disc起源と思われる

IREが遠心力バリアの手前で停滞・弱い降着ショックを生じ、温度が上昇（~190 K）することでIREが膨張し、一部ガスは垂直方向へ流れだす

遠心力バリア付近のSOの速度がIREの回転速度より遅い（図c）
MNRASL 467, L76–L80 (2017)=> 衝撃により回転エネルギー → 熱に変わるとともに、

7 envelope/disc system.

to the IRE model (dotted white contours) and the Keplerian model (black contours). Note that the small undulation of the lowest contour of the Kmodel within 0.2 arcsec from the centre is an artefact due to a limited mesh size of the model. In addition to the component associated with the cenbarrier, a faint high-velocity component can be seen (pink dotted lines). Yellow dotted arrows indicate the voids of the emission. (d) Schematic illustr
the physical and chemical structures of the L1527 envelope/disc system.

absent in the IRE exshown by the black contours in Fig. 2(c). Here, the outer radius
角運動量を垂直方向に放出している
ier. The SO emission isis arbitrarily set to be 160 au. Althoughestimated
It istheconfirmedthat the velocity
SO emissionisis absent in the IRE except
forthevicinityofthecentrifugal
barrier. The SO emission is
CCH emission. In orderstill higher than the observed one in some parts, the observed PVseen inside the broadened region of the CCH emission. In ordershown by the black contours in Fig. 2(c). Here, the outeris arbitrarily set to be 160 au. Although the estimated velo
still higher than the observed one in some parts, the observ

#44

A dearth of short-period massive binaries in the young massive
star forming region M17:
Evidence for a large orbital separation at birth?
H. Sana1 , M. Ramírez-Tannus2 , A. de Koter1, 2 , L. Kaper2 , F. Tramper3 , and A. Bik4
1e-mail: hugues.sana@kuleuven.be・大質量連星の誕生時の特性はよくわかっていない。
Anton Pannekoek Institute for Astronomy, University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands
e-mail: m.c.ramireztannus@uva.nl=> 非常に若い大質量星形成領域M17に注目
European Space Astronomy Centre (ESAC), Camino bajo del Castillo, s/n Urbanizacion Villafranca del Castillo, Villanueva de laInstitute of astrophysics, KU Leuven, Celestijnlaan 200D, 3001 Leuven, Belgium

2.02153v1 [astro-ph.SR] 7 Feb 2017

2
3
4

Canada, E-28 692 Madrid, Spain
Department of Astronomy, Stockholm University, Oskar Klein Center, SE-106 91 Stockholm, Sweden

Received September 15, 1996; accepted March
16, 1997 Msun）の視線速度を調べた
・大質量前主系列星11天体（M~6-20&A proofs: manuscript no. m17_PaperII_v2

al. 2012; Kobulthus particularly
one of the largest
ere massive stars

roscopic binaries
1
rad < 20 km s ).
ies of small andce massive stars
013a; Kiminki &nstall et al. 2015;
d established thatof one month or
km s 1 ). Here wef the massive starat quantifying its

ndidate mYSOs1
which were conB stars observednce (B111, B164
253). The known
well as the CEN
ecause they wereTable 1. Sample stars and journal of the observations. Sources classified as PMS in Paper I appear in bold. Sp.Types are from Paper I.

ABSTRACT

σ1D = 5.6

0.2 km s-1

Aims. The formation of massive stars remains poorly understood and little is known about their birth multiplicity properties. Here,↓（若い星の集団でよく見られる）短周期連星がIdentifierwe aim
Sp.TypeMJD
vradd(3)the strikingly low radial-velocity dispersion measured for a sample of 11 massive pre- and nearto quantitatively
investigate
(2)
1
1
B92
CEN(OI)
000( (kms )(pc)
main-sequence2 450
stars
1D = 5.6 ± 0.2 km s ) in the very young massive star forming region M17, in order to obtain first constraints
B111
2 (337)
O4.5V 6490.116 3.4 ± 0.8 0.695
properties
B163. . on. the multiplicity
kA55455.982. . . of young
0.154 massive stellar objects.
compute thedispersion of synthetic populations of massive stars for various multiplicity properties and
B164
25Methods.
O6VWe 6490.123
1.8 radial-velocity± 2.3 0.372
B215. . we. compareO9-B1Vthe
6115.249
12.3 1D± 1.0
0.856
obtained
distributions
to the observed value. We specifically investigate two scenarios: a low binary fraction and
B243
51a dearthB8V
6114.188 19.5
± 1.8systems.
0.607
ofshort-periodbinary...
...
...
6490.140 20.2 ± 1.5 . . .
+0.16
Simulated
populations
B253
26Results.
B3-B5III
6490.197
9.5 ± 0.8 with
0.573low binary fractions ( fbin = 0.12 0.09 ) or with truncated period distributions (Pcuto↵ > 9 months)
to reproduce
observed within their 68%-confidence intervals. Furthermore, parent populations with fbin > 0.42
B268
49are able
B9-A0
6114.236 the1.4low± 1.0 1D0.686...
. . or
. Pcuto↵. .<. 476114.259± 1.1 at . the. . 5%-significance level. Both constraints are in stark contrast with the highd can be 17.1
rejected
1D binary fraction and
...
...
...
6490.169 4.3 ± 1.3...
plethoraofshort-period
systemsinfew
Myr-old,
well
characterized
OB-type
populations.
To
explainthedi↵erencein the context of
B275
24
B7III 5054.125 11.2 ± 1.4 0.698
scenario
would requireof the outcome of the massive star formation process. In the context of the second scenario,
B289
31the firstO9.7V
5456.024
2.0 ± 2.5a variation
1.390
...
. . compact.
. .binaries.
6114.290
2.3 ± 2.0
. . and the cut-o↵ period may be related to physical length-scales representative of the bloatedmust form
later .on,
B311
16 (258)
O8.5Vz 6490.097
4.2radii± 0.4 or1.601
pre-main-sequencestellarof their accretion disks.
B331
92
late-B 6115.276 13.9 ± 1.5 1.194
Conclusions.
Iftheobtained
constraints
for the M17’s massive-star population are representative of the multiplicity properties ofB337
93
late-B 6489.185 7.5 ± 0.6 1.171
massiveyoung
stellarobjects,
ourresults
may provide support to a massive star formation process in which binaries are initially
...
...
...
5.6 ± 0.2
...
1D
(1)

M17のmYSOsにもあるとすると、あり得ない数値
参考：他のOB型星の集団 σ

２つの可能性：~25 - 35 km s-1↓(i) 連星率が低い(ii) 短周期連星を持たない
Key words. binaries: spectroscopic — Stars: early-type— Stars: formation — open clusters and associations: individual: M17

formed at larger separations, then harden or migrate to produce the typical (untruncated) power-law period distribution observed in

Notes. (1) Bumgardner (1992). (2) Alternative IDs: CEN (Chini et al.
few Myr-old OB binaries.
1980) and OI (Ogura & Ishida 1976). (3) Projected distance of the stars
to the center of the cluster (↵ = 275.124574 , = 16.179031 ).

the candidate mYSOs or to the confirmed PMS stars in our sample (Table 1) yield similar values of 5.9±0.2 and 6.9±0.4 km s 1 ,#44の続きFig. 2. Same as Fig. 1 for di↵erent cuto↵ periods (Pcuto↵ ). The distribution
corresponding to Pcuto↵⇡ 9 mois intermediate
between those of
・観測結果とMonteCarlopopulation
synthesisから得られる結果の比較
Pcuto↵ ⇡ 8 yr and 1 month and is not shown for clarity.
初期星団の速度分散σcl = 2 km s-1
連星率fbin： 0.00から1.00まで0.01刻みで変化
Sana et al.: Binarity in M17

rameter distributions, the observed velocity dispersion in M17 isbest reproduced by a low binary fraction=> fbin = 0.12+0.16
連星率は低い0.09 .
Similarly, no simulation with fbin
0.34 (resp. 0.42) can
しかし、
reproduce M17’s 1D within their
0.90 (resp. 0.95) percentiles
and the corresponding hypotheses
can thus be rejected at the 10 IC 1805）~ 40% over
M17と似た星団質量の天体（e.g.
and 5%-significance levels, respectively.
Fig. 1. Simulated 1D distributions for di↵erent parent binary fractions.
The vertical dashed line indicates the observed 1D for our M17 sample stars. The median, 0.16 and 0.84 percentiles of the simulated distributions are indicated on the upper part of the graph. The distribution
corresponding to fbin = 0.03 is not shown for clarity; it is almost entirely dominated by the velocity dispersion of the cluster and thus peaks
strongly at 2 km s 1 (⇡ cl ).

Fig. 2. Same as Fig. 1 for di↵erent cuto↵ periods (Pcuto↵ ). The distribution corresponding to Pcuto↵ ⇡ 9 mo is intermediate between those of
Pcuto↵ ⇡ 8 yr and 1 month and is not shown for clarity.

rameter distributions, the observed velocity dispersion in M17 isbest reproduced by a low binary fraction fbin = 0.12+0.16
0.09 .
Similarly, no simulation with fbin
0.34 (resp. 0.42) canreproduce M17’s 1D within their 0.90 (resp. 0.95) percentiles
and the corresponding hypotheses can thus be rejected at the 10
and 5%-significance levels, respectively.

・別の案：Runaway star - 以下の条件が必要
the multiplicity properties derived by Sana et al. (2013a) for
Galactic young open clusters as the baseline in our population
synthesis. These include fbin = 0.7 and fPorb / (log10 Porb ) 0.5 ,with 0.15  log10 (Porb /day)  3.5 and a flat mass-ratio distribution. Adopting an Öpik law all the way down to orbital periodsof 1 day would not change our conclusions.
The cluster velocity dispersion ( cl ) is computed from thevirial theorem, assuming a total mass for the M17 cluster of
1.4 ⇥ 103 M (Stutzki & Guesten 1990) and a typical 1 pc cluster
radius. This yields cl = 2 km s 1 . Other published mass estimates range from a few times 103 to 6⇥104 M (Pérez-Beaupuits
et al. 2015; Povich et al. 2009), corresponding to cl = 1.4 to
5.2 km s 1 . Adopting the former value has no impact on our
results while picking the latter one would only strengthen ourconclusions because it leaves almost no room for binary contributions to the observed 1D value. In the following, we adopt
1
cl = 2 km s .
Each Monte Carlo run is formed by 105 realisations and uses
the sample size, number of epochs, time sampling and measurement accuracy of the observational sample. The density distributions of the simulated 1D for given parent multiplicity properties are then constructed and used for tests of hypotheses thatallow us to accept or reject the specific parent population at various significance levels.

① M17ではすでにcore
colapseが起きている
3.3. A truncatedperiod distribution② Runaway starの放出速度は遅い
We now assume that the binary fraction is consistent with con③ 初期大質量星連星率が低い or 短周期連星系がない
3.3. A truncated period distribution

We now assume that the binary fraction is consistent with constraints from OB stars in young open clusters, that is, fbin = 0.7
for Porb < 3500 d (Sana et al. 2012), but that the binary
population is composed of longer-period binaries, which typically results in a lower 1D (Gieles et al. 2010; Hénault-Brunet
et al. 2012). In practice, we adopt a truncated period distribution
where the orbital periods of all binaries with Porb < Pcuto↵ areiteratively re-drawn from the parent distribution until they are
equal to or larger than Pcuto↵ .
Resulting distributions for representative values of Pcuto↵ areshown in Fig. 2. Distributions with Pcuto↵ ⇡ 8 yr have a medianthat best matches the M17’s observed 1D while distributions
with Pcuto↵ & 283 d (⇡ 9 months) are all able to reproduce theobserved 1D within their 68% confidence range. Parent populations with Pcuto↵ < 131 d (resp. 47 d) can be rejected at the 10
and 5%-significance levels, respectively.

straints from OB stars in young open clusters, that is, fbin = 0.7
for Porb < 3500 d (Sana et al. 2012), but that the binary
populationis composed of longer-period binaries, which typi・短周期連星系がない可能性
cally results観測事実：
in a lowerOB型連星の20％が短周期連星
1D (Gieles et al. 2010; Hénault-Brunet<= でも形成メカニズム不明
et al. 2012). In practice, we adopt a truncated period distribution
where the『大質量連星はもともと広い間隔で生まれる』ことが考えられる
orbital periods of all binaries with Porb < Pcuto↵ areiteratively re-drawn from the parent distribution until they are
3.2. A low binary fraction

We first investigate the possibility that the low RV dispersion observed in M17 results from a low binary fraction. In our MC experiment, we keep the distribution of the orbital properties fixed
and we vary the binary fraction from 0.00 to 1.00 in steps of
0.01. The obtained 1D distributions for representative values of

4. Discussion

4.1. A low binary fraction?

Under the assumption that the distribution of orbital parametersin M17 is not di↵erent from that of other well characterized OBstar populations, a low RV dispersion can only be explained by a
low binary fraction. This, however, would make M17 quite atypical. Populations of massive stars in clusters across the entire
mass range have all revealed a large binary fraction, from mod-

71v2 [astro-ph.GA] 22 Apr 2017#45

YOUNG STELLAR OBJECTS IN THE MASSIVE STAR-FORMING REGIONS W51 AND W43

G. S ARAL 1,2,3 , J. L. H ORA 2 , M. AUDARD 1 , X. P. KOENIG 4 , J. R. M ARTÍNEZ -G ALARZA 2 , F. M OTTE 5,6 , Q. N GUYEN -L UONG 7,8 A. T.
S AYGAC 9,10 , H. A. S MITH 2
20 April 2017; Published in ApJ

ABSTRACT(left) and W43 (right). Red dots: class I; green dots: class II; blue dots:
星形成過程の理解には、大質量星形成を理解する必要がある
Figure 6. [4.5]-[24.0] vs. [3.6]-[4.5] IRAC color-color diagram for W51

transitioncomplexes
disk candidates;
black
dots:
classtwo
III and
red inasterisk:
We present the results of our investigation of the star-forming
W51
and
W43,
ofphotospheres;the brightest
embedded
protostars.
the2MASSとSpitzerのデータを使って、W51とW43でYSO候補天体を同定した
first Galactic quadrant. In order to determine the young stellar object (YSO) populations in W51 and W43=>
we used color-magnitude relations based on Spitzer mid-infrared and 2MASS/UKIDSS near-infrared data. We
identified 302 Class I YSOs and 1178 Class II/transition disk candidates in W51, and 917 Class I YSOs andTable
5 YSOs in both regions
5187 Class II/transition disk candidates in W43. We also identified tensSourceof groups
ofClassification Summary
using the Minimal Spanning Tree (MST) method. We found similar cluster densities in both regions, eventhough Spitzer was notby
ableGutermuth+08,09
to probe the densest part of W43. By using the Class II/INumber
ratios,of wetraced the
Objects
IRACカラーで選択
relative ages within the regions, and based on the morphology of the clusters, we argue that several sites ofType
W51 We used
W43 spectral
star formation are independent of one another in terms of their ages and physicalconditions.
energy distribution-fitting to identify the massive YSO (MYSO) candidates
since they playrole in theClass I
302 a vital917
色等級図や二色図を使い、Figure 4. Color-colorand magnitude-color
diagrams
for the W51them
region.to see if they are related
Class II to any massive 1101
4993 tracers
star formation
process, andthen examinedstar formation
Upper panel shows[4.5]-[5.8]vs.
[3.6]-[4.5]and[4.5]-[8.0]vs.
[3.6]-[5.8]Transition
Disks
77
194and 14 insuch as UCH II regions, masers, or dense fragments. We identified 17 MYSO candidates in W51,
IRAC color-color
diagrams. Lower panel shows [3.6]-[8.0] vs. [3.6] and・系外銀河等のcontaminantを除去
Embedded protostars
56
144
and foundthat groups
YSOs
candidates
are positionallyassociatedwith[3.6]-[4.5] vs.W43,[3.6] respectively,
IRAC magnitude-color
diagrams.
Red dots:ofclass
I; hosting MYSO FaintClass I*
100
303
H・YSOをクラス分け
II regionsin W51,though black:
we doclassnot III
seeandany
MYSO candidates associatedgreen: class II;
blue:transition
disk candidates;
photoFaint Class with
II* previously
1618identified
4549 massive
spheres; gray:dense
AGB starcandidates;orangetriangles:
backgroundcontamifragments in W43.
Faint Transition Disk*
35
46
nants.
Class
III/Photospheres
157,
165
429,
895
Keywords: infrared: stars — stars: early-type — stars: formation —AGB
stars:
pre-main sequencecandidates
180866

1. INTRODUCTION

Near-infrared and mid-infrared observations have revealed
a broad range of stellar densities of YSOs in star-forming regions (Bressert et al. 2010). While we have a detailed pictureof the formation of low-mass stars with advances in observational techniques and improved theoretical knowledge in thelast decade, our understanding of massive star formation isstill lagging. Massive stars form and evolve quickly and they
produce strong winds and outflows which heat, ionize, anddisrupt their natal molecular cloud, which makes the observa-

Other **
524
1416
Unclassified
124, 094 198, 409
Galactic star-forming
regions285,
are 252
the ideal
places for this purTotal
641, 732

pose, since they host many active sites of star formation indifferent* These
environments
at various
stagesof evolution.
are the YSO candidates
with lower
confidence.
** Includes PAH emission-dominated sources, H
2We started a detailed study of a sample ofmini-starburst
shock emission dominated sources, broadline AGN
complexes(Motte et al. 2003) to better understand how mascandidates, and PAH galaxy candidates. Unclassified
sive stars
form
affecttheirsourceslackanddetection
in fourbandsenvironments(either HKS , IRAC and the effects
2, or IRAC
1, 2, 3,and 4) or a bright
24 µm
that they1 andhaveon theformationof MIPS
lower-massstars (Saral
detection.
& Hora 2013). Precursor clouds of young massive clusters(YMC) (Bressert et al. 2012; Ginsburg et al. 2012; Walker etof 2.16
for [Ks],
and 3.55,to4.44,
5.73,
and 7.87µm for [3.6],
al. 2016)
are thought
be thepotential
descendantof mini-Note. — Table 7 is published in the electronic edition of the Astrophysical Journal. A portion isa Subgroups identified with d = 5900 are named under the main group with main number and an acb Central coordinates of the YSO groups or subgroups.
c Number of stars with infrared excess. Includes Class I, Class II, deeply embedded protostars, and#45の続きd I; Class I, II; Class II candidates

W51

W43

NumberAND
in parentheses
YOUNG S TELLAR O BJECTS INe W43
W51 indicates uncertainty in ratio.

20 pc
3

20 pc

W51

W51

4a

5a

5

-1.5
2a

6

2e

2d

Declination (Deg)

1

9

2f
3a
1b
1d

1g
1a

1c

14.0

8

4

14.0

Declination (Deg)2c2bDeclination (Deg)

14.5

14.5

2

1e

-2.0

-2.5

1f

7

-3.0

13.5

13.5

・Minimal Spanning Tree (MST) 法を使ってYSOの空間分布からクラスターを同定=> W51で９つ（YSOの60%）
W43で51つ（YSOの73%）

11

f Number of deeply embedded protostars.
g Number of transition disk candidates.

291.0

290.5
Right ascension (Deg)

291.0

290.0

290.5
Right ascension (Deg)

290.0

Figure 11. Left: the 51 MST groups determined using the SLF method with a branch length oFigure 9. Left: the nine MST groups identified using dc = 8200 .5 are shownMSTwith subgroups
group numbers
overlaidwith
on theIRAC length
4.5 µmofimage.
Right:
theThe16 YSOsMST are plotted in codetermineda branch
5900 (1.61pc).
subgroups identified by dc = 5100 are shown with a cluster number and a letter (see Table 6).

・Class II/Iの比較 W51: 2.50、W43: 4.58=> W43の方が進化が進んでいる
Table 8Clusters/Subgroups in W51

Parameter

・大質量YSO（> 103 Lsun）を同定（SED FittingとBayesian法を使用）=> W51: 17天体、W43: 14天体

dc = 5100

SLF
MethodPara

Number of clusters

916

Percent in clusters
Group diameter
Class II/I Ratio

60
9600 .5-117800 (2.5-30 pc)
2.50(0.25)*

32
4000 -25600 (1-7 pc)
2.06(0.27)*

Cutoff distance
YOUNG S TELLARO BJECTS 82
IN00 .5(2.17
W43pc)AND W51 5100 (1.34 pc)

Table 12Clustering Analysis Summary* Number in parentheses indicates uncertainty in ratio.

Region
W51

Number of Clusters

CutoffClusteredcation of the far-IR
bumpdistance(00 ) is known.
Fraction

9168200 .5(2.17 pc)
5100 (1.34 pc)

60
32

NumberCutoff dPercent
Group dClass II/

* Number iClass II/I
Ratio

Method2.50(0.25)
2.06(0.27)

a

We ha

b

Figure 10. K-S statistic distribution of random MST groups, between each of the 1000 random realizations
and the mean random cluster distribution. The
a green
W43
51 W51 region (left),
8100 (2.20
73 region 4.58(0.19)
histogram shows the groups identified with SLF
branch length dc = 8200 .5 in theand dcpc)= 8100 in the W43(right). The blue histogramb
95 and dc = 590059(1.61
3.71(0.19)
shows the groups identified with branch length dc = 5100 in the W51 region (left),
in00theW43pc)
region (right).43
Dashed lines
in both figures represent
c random clusters and the 15
00 (0.90
a boththe K-S statistic value between the mean CDF
of all
CDF of observed48MSTgroups,of histograms. For
G305pc) with the same
76 color code
3.50(0.30)bregions, the inferred probability that the observed MST groups are drawn from32
a random distributionis lesspc)
than 10-3 . 44
3000 (0.60
3.00(0.30)
c

is described in detail by Robitaille etG333
al. (2007). Only objects 11
2
with at least six data points were fitted,c and we used e- /2 45
values as weights. We chose the fitsW49
that satisfy the criterion 3
7of ( 2 - 2best )/ndata < 3 as in Robitaille et al. (2007). Only objects with a minimum of 10 fits satisfying
theW51
defined best-fit 5
Rescaled
criteria were used when calculatingRescaled
the weighted
W43 mean val- 8
ues for the physical parameters. WeRescaled
allowedG305
the cAV to vary 4cbetween 0 and 40 mag, which gives
enoughflexibility
and 2
Rescaled G333
prevents unphysical solutions, and is also consistent with the
AV values for MYSOs that are forming
UCH IIs (e.g., Hana Straight-line
fit method
son et al. 2002). We also let the distanceto vary within 10%
b

a

00

pc) W51, and815.6 kpc4.50(0.40)
ranges of73
5.4(1.30
kpc forfor W43. For sources00bpc)
54 for sources
4.00(0.40)
without a52
2400 (0.90µm detection(except
in the saturateda
96 (5.20 pc)regions), we
52 used an
2.11(0.43)
or high-backgroundupper limit based on00b
40 (2.20 estimatepc)
21µm by 2.00(0.64)
the completenessat 24
Gutermuth and Heyer(2015) in order
to better constrain the SED. The luminositya12000 .94(6.47)
74
1.61(0.21)
and temperature
are directly derived
fromthe SED, anda then9100 .82(4.91 pc)
73
2.79(0.21)
mass and52age
valuesaredetermined
using
the
evolutionary
00 (2.78apc)
73
0.55(0.15)
tracks in78
a 00Hertzprung-Russel
diagram.
The lack of photoa(4.22 pc)
70
0.03(0.04)
metric points at wavelengths longer than 24 µm means thatthe disk and envelope parameters are poorly constrained. The
total dust mass, for example, can only be estimated if the lo-

・W51、W43、その他（G305やG333、W49）の比較- W51もW43も似たYSO星団密度 <=> G305とG333はより密集した星団
The method using the short branch length distance enables us to see subgroups in the region.

c Saral et al. (2015)- W49: W51やW43よりも小さい領域。

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