星形成ゼミ 論文速報
May 9, 2014
小屋松 進

Age	
  Gradients	
  in	
  the	  Stellar	
  Populations	
  of	
  Massive	  Star	  Forming	
  Regions	
  Based	  on	
  a	
  New	  Stellar	
  Chronometer	  	
  Konstantin	
  V.	
  Getman,	
  Eric	
  D.	
  Feigelson,	
  Michael	
  A.	
  Kuhn,	  Patrick	
  S.	
  Broos,	
  Leisa	
  K.	
  Townsley,	
  Tim	  Naylor,	
  Matthew	
  S.	
  Povich,	
  Kevin	
  L.	
  Luhman,	  Gordon	
  P.	
  Garmire	
  
【背景】	
  
大質量星形成領域(MSFRs)での星形成の歴史を探る上で信頼できる”時計”が存在しない	  近赤外線とX線の測光を用いた新たな手法(AgeJX)により星の年齢推定をおこなう	
  
【手法】	
  
おうし座分子雲におけるLX–M関係を用いてX線光度から星質量を導出	  J	
  バンド光度を前主系列星進化モデル(質量に依存)と比較して年齢を推定	
  	
  AgeJXは分子雲に埋もれた円盤をまとった星から広い進化段階に対応	
  	
  MYStIXの31784のメンバー候補のうちX線で明るい5525の低質量星にAgeJXを適用	
  15のMSFRsの100を超えるサブクラスターについて年齢の中間値を求めた	
  
MYStIX	
  (Massive	  Young	  Star-­‐Forming	
  Complex	  Study	
  in	
  Infrared	
  and	  X-­‐ray)
	
  
X線,	
  中間赤外線,	
  近赤外線カタログを用いて20のOB型星が支配的な大質量星形成領域の性質を調査 (Feigelson	
  et	
  al.	  2013)	
  	  【結果】	  
MYStIXの領域にわたって年齢の勾配が存在	
  MSFRが広い年齢分布を示す→異なる年齢が空間的に離れて分布している？	
  – 60 –
分子雲で覆われた場所で若く,	
  星団で中間的,	
  広がった分布では古い
近赤外の色J-­‐Hは減光の測定量の代用で,	
  若い埋もれた星団の年齢の指標として役立つ	  – 61 –Median(AgeJX) [Myr]Orion Nebula
3.5
A;1.4
B;1.1
C;1.5
D;2.7
U;2.0

3.0
2.5
2.0
1.5
1.0
0.5
0.8

0.9

1.0
1.1
Median(J-H) [mag]

1.2

1.3

Figure	  7

	  Figure	  6Fig. 7.—オリオン星雲における解析の例	
  AgeJX analysis of the Orion Nebula. Upper left: The full MPCM sample with subclus	
   Age distributions for the MYStIX massive star forming regions
Fig. 6.— Boxplots of the	  JXter membership coded by color and the elliptical contours (yellow) showing the core radii of the
各大質量星形成領域での年齢の分布	
  他の14の大質量星形成領域についても同様の解析をおこなっている	
  shown in order of youngest (top) to oldest (bottom) median ages. The median age is shown
isothermal structures from Kuhn et al. (2014). These symbols are superimposed on the 8.0 µm

as a dark bar. See text for details.

Spitzer-IRAC image. In this and following figures, the unclustered/ambiguous stars are shown inmagenta color, and the Chandra field of view is outlined by the green polygon. Upper right: The

A&A proofs: manuscript no. ms

Table 1. Observation parameters. The projection center for the all the observations is α20
of HD 34078.

Molecule Transition Frequency Instr. Config.
Beam
PA
Vel.
Dense	
  molecular	  “globulettes”	
  and	  the	
  dust	  arc	
  towards	
  
the	
  
GHz)
( )
(deg) ( km
CO J = (1–0) 115.27120 PdBI
5D
5.1 × 3.8
139
0.
runaway	
  O	
  star	
  AE	
  Aur	  (HD	
  34078)⋆	  	
  Molecule Transition Frequency
Instr.
#pol
F
B
Res.
R
′′

12

eﬀ

P.	
  Gratier,	
  J.	
  Pety	  P.	
  Boissé,	
  S.	
  Cabrit,	
  P.	
  Lesaﬀre,	
  M.	
  Gerin,	
  CO
and	
  JG= .	
  (2–1)
Pineau	  des	
  30m/E2
Forêts	
  2	
  
12

12

CO

J = (1–0)( GHz)
230.53800
115.27120

30m/E0

2

0.91
0.95

eﬀ

0.59
0.78

(′′ )
11.2
22.5

( km

Runaway	
  Star【背景】	
  一部の“Runaway	
  Stars”は星間物質との相互作用によりarc-­‐likeなIR構造を示す	  その性質や形成過程はよくわかっていない	
  
【観測】	
  
Runaway	
  O-­‐Star	
  AE	
  Aur	  (HD	
  34078)の周りの12CO(1–0)を	
  IRAMの30	
  m鏡およびPdBIで観測	
  
	
  
	
  
	
  	
  

Notes. (a) 	
  listed as on-source time/telescope time (b) evaluated at the mosaic phase center
correction for primary beam attenuation, see top panel of Fig. 3)

天球上で大きな固有運動をもつ恒星	
  Figure	  1 24 µm image towards HD 34078, from (France et al.
Fig.
1. Spitzer	  2007).
Contours are equally spaced every 50 MJy/sr, starting from100
are measured
from
the 	
  star position, displayed asHD	
  MJy/sr.
34078	
  OﬀsetsのSpitzer	
  
24	
  μm	
  連
続波	
  
a black star. The red arrow represents the star on-sky proper motion for(Gray	
  Scale	  &	
  Black	
  Contours)	
  the next 1000 years. The solid open white curve displays a parabola
黒星印in (矢印の起点)：星の位置	
  pointing
this direction, with its focus at the star and a projected starapex
distance Robs = 15′′ (see Appendix
A for a description of pro赤矢印：天球上の固有運動	  (長さは1000年分)	
  jection eﬀects in a paraboloid). The two dashed curves correspond to
白い歪んだ円：今回のPdBI観測領域	
  variations of Robs by ±5′′ . The closed white contour corresponds to theregion mapped with the PdBI interferometer. The right ascension axis
increases towards the left.
highly excited H2 , indicating the unexpected presence of dense
4−3

steady-state w
Boissé et al. (2star and the IR
tance of HD 3
1999), is large
bow shock forsis (Martins etwitnessing thepressure on gr
natively, the w
mated (as pro
IR bow shockphysical proce
Distinguis
to better under
tween HD 340
absorbing gas
aﬀect the win
sizes in othermapped the bowith the Platewith IRAM-30
at ≃ 4′′ resolu
times better th
The observ
The resulting
scribed in Sec
summarize an

2. Observat

Table 1 summ

【結果】	
  単一鏡フラックスの1/3が2つの小さく明るい”Globuletts”	  	
  (Figure	  2	
  right	
  の	
  1,	
  2)	
  に起因する	  星方向のスペクトルはダブルピークを示していて,	
  2つのGlobulettesが視線方向に存在	
  他波長での観測を説明できる	  CO	
  Globulettes	
  はほぼ放物面に沿って存在	  重力的に束縛されていない	  圧力により支えられているかも	
  星と広がったガスの相互作用の結果だろう	  (例えばBow	  Shockでの熱的不安定性?)	
  
	
  
	
  
P. Gratier et al.: Dense molecular “globulettes” and the dust arc towards the runaway O star AE Aur (HD 34078)
	  Figure	
  5	
  (↑)
	
  
星方向のスペクトル	
  上：12CO	
  (1–0)	
  (黒)	
  と12CO	
  (2–1)	
  (青)	  (30	  m)	  下：12CO	
  (1–0)	
  (30	  m	  +	  PdBI)	
  	
  Figure	
  2	
  (←)

12CO	
  (1–0)積分強度図	
  

左が30	
  m鏡で右が30	
  m+PdBI	  Fig. 2. Spatial distributions of the 12 CO(1–0) line integrated intensities. The left/right panels show the data from IRAM-30m only (at an angularresolution of 22.5′′ ), and the hybrid synthesis data from PdBI + IRAM-30m (at an angular resolution of 5.1′′ × 3.8′′ at a position angle of 139◦ ),respectively. The beam shape is drawn as an ellipse in the lower right corner of each panel. Black contours are drawn at 5, 10, 15, 20, 25 and30 σ. Oﬀsets are measured from the star position, denoted by the black star with the red arrow representing the star on-sky motion for the next

【結論】	  O型星のこんな近くに”Globulettes”が存在するのは大きな固有運動と関係しているだろう	  CO	  Globulletes	
  と	
  IR	
  arc	
  にはよい相関が見られ,	
  星からの放射・風と星周ガスとの相互作用の結果	
  (相互作用の詳細はよくわからない)	
  星風による質量損失率はUV放射から推定されるよりも非常に大きい	
  これは大きなIR	
  arcには都合がよいが低速なCO	
  Globulettesの説明は難しい	
  ダスト粒子への輻射圧の効果も問題	
  
	
  
	
  

The	
  tiny	  globulettes	
  in	
  the	  Carina	
  Nebula	
  ⋆	
  	
  
T.	
  Grenman	
  and	  G.	
  F.	
  Gahm	  Figure	
  1	
  (↓)
	
  
Carina	  Nebula	
  の中心領域	
  (1.3	
  x	
  1.5	
  deg)	
  
四角で囲った部分がGlobulettesが含まれるHST	
  ﬁeld	  【背景】	
  中心やや上にη	
  Carina	
  4つの矢印は星団	
  T. Grenman and G. F. Gahm: The tiny globulettes
in the Carina Nebula
生まれたての星団を囲むHII領域の多くには小さな分子雲が豊富に存在する	  ized gas (e.g. Dyson 1968; Kahn 1969; Tenorio-Tagle 1977).
Consequently, the objects are expected to develop bright rims on
こうした“Globulettes”は可視光で明るい星雲
the side facing the cluster because of the interaction with stellar
を背景に暗いシルエットとして見える	
  light. In addition, the models predict that dusty tails emerge from
theN
cloud
cores.
is therefore puzzling that most globulettes lack
Carina	  ebula	
  のItGlobulettesのカタログを作り,	
  any trace of bright rims in H↵, and that most are round, or only
他のHII領域との比較をおこなう	  slightly elongated, without any trace of tails.
In a recent study by Gahm et al. (2013; hereafter Paper 2),【手法】	
  based on NIR imaging and radio molecular line observations ofglobulettes
the Rosette Nebula,
it was found that the objects
Hubble	  Space	
  inTelescope	
  のHα画像から
contain dense cores, which strengthens the suggestion that many
Globulettesを特定してサイズ・質量・密度を
objects might collapse to form planetary-mass objects or browndwarfs that are accelerated outwards from the nebular complex.
求める	
  The whole system of globulettes and trunks expands outwards
半径：半長軸と半短軸の幾何平均	  from the central cluster with velocities of about 22 km s 1 . In thecase where more compact objects are formed inside some glob質量：6563	
  Å	
  willにおけるダスト減光を測定	
  ulettes, theyescape and become free-floating objects in thegalaxy. In both the optical and radio/NIR surveys (Papers 1 and	  2) it was concluded that the density is relatively high even close

to the surface layers, which could explain why the objects lackextensive bright rims in H↵. Some of the optically completely
dark objects were discovered to have thin rims manifested in
P and H2 emission. In a follow-up study of the NIR images,Mäkelä et al. (2014) found that some smaller globulettes are also
crowned by thin bright rims that are not seen in H↵.
The present study is an inventory of globulettes in the Carina
Nebula (NGC 3372) based on images taken from the Hubble
Space Telescope (HST) through a narrow-band H↵ filter. Basicparameters, like size and mass, are derived and we compare theFig. 1. Image of the central region of the Carina Nebula, where
the HST fields containing globulettes are marked. The locations

【結果】	
  今までサーベイされた領域で最も多い288のGlobulettesの位置を確認	  よくまとまった高密度クランプで,	
  だいたいは薄い”おおい”や”尾”をもたない	  	
  縁の明るいものは少数ではあるが他の領域よりは多い	
  いくつかは若い星団の方向を向いて伸びた構造をしている	  
多くは孤立していて隣の分子構造とは1.5	
  pc以上離れている	  Figure	  2Nebula	  (↓)
T. Grenman and G.
F. Gahm:
Figure	  3	
  (←) The tiny globulettes in the Carina
	
  
	
  
Carina	  Nebulaで見つかったGlobulettesの例	
  Carina	  Nebula	
  の Globulettes	
  の性質	
  
上から4行が最も典型的な暗いもの	
  上：半径	
  (arcsec)	
  次の2行が明るいハローをまとったもの	  中：半径をkAUに直したもの	
  最後の2行が伸びた尾をもつもの	  下：質量	
  (MJ)	
  

T. Grenman and G. F. Gahm: The tiny globulettes in the Carina Nebula

Table 1. HST archive data used.

T. Grenman and G. F. Gahm: The tiny globulettes in the Carina Nebula

Field/Target R.A.
Dec.
ImagesFig. 3. Left: Distribution of average radii as measured for all globulettes found in the Carina Nebula expressed in arcsec. Middle:(J2000.0)(J2000.0)
The corresponding distribution of average radii expressed in kAU and adopting a distance to the complex of 2.9 kpc. The verticalarrow marks the peak in the corresponding accumulated size distribution for objects in seven H ii regions (De Marco et al. 2006).
Right:of masses-59:47:42
for Carina globulettes
1 / PosDistribution30
10:41:27
J9dk09010 less massive than 10 M J , expressed in Jupiter masses.
2 / Pos 30
10:41:38-59:46:17
J9dk09020
3
/
Pos
30
10:41:40-59:44:41
J9dka9010
iny globulettes in the Carina Nebula
4 / Pos 19
10:42:23-59:20:59
J9dk12010
5 / Pos 19
10:42:48-59:19:44
J9dk32010
6 / Tr 14
10:43:07-59:29:34
J900c1010
tion of average radii as measured for all globulettes found in the Carina Nebula expressed in arcsec. Middle:
7 / Tr 14
10:43:23-59:32:06
J900b1010
istribution of average radii expressed in kAU and adopting a distance to the complex of 2.9 kpc. The vertical
8 / Tr 14
10:43:24-59:27:55
J900c2010
k in the corresponding accumulated size distribution for objects in seven H ii regions (De Marco et al. 2006).
9 / Tr 14
10:43:39-59:34:37
J900a1010
f masses for Carina globulettes less massive than 10 M J , expressed in Jupiter masses.
10 / Tr14
10:43:41-59:30:17
J900b201011 / Tr 14
10:43:47-59:35:53
J90001020
12 / Tr 14
10:43:55-59:37:09
J90001010
13 / HH 666 10:43:58-59:54:39
J900a9010
14 / Tr 14
10:43:59-59:32:38
J900a2010
15
/
Tr
14
10:44:00-59:28:05
J900b3010
obulettes found in the Carina Nebula expressed in arcsec. Middle:
Fig. 2. Examples of globulettes found in the Carina Nebula num16 /toHH
666
10:44:01-59:58:42
J900b9010
U and adopting a distancethecomplexof 2.9 kpc.
The vertical
17 / Tr H
16ii regions
10:44:05-59:40:16
J900c5010
bered according to Tables A1-A5 in Appendix A. The most typstribution for objects in seven(De Marco
et al. 2006).

not is a bit arbitrary, but this is less important in our statisticalanalysis. Figure 5 shows some examples of cloudlets with peculiar shapes, and their locations can be found in Figs. B1-B6,
except for j which is in an HST field void of globulettes. This object, like object c, was noted and labelled in Smith et al. (2003).
Objects B and D could be similar to standard globulettes, butsince B appears to be surrounded by strong H↵ emission, and D
is in the background behind strong foreground emission, theseobjects could not be measured for extinction. The other objects
have peculiar dusty tails, and object A even shows three suchoutgrowths. These tails can be the result of erosion of elongatedobjects with peculiar density structure, or they may be examples of dust-enshrouded jets from embedded sources. Objects
CN 219 and 241 were recognized as HH-objetcs in Smith et al.
(2003), but their nature remains unclear. The objects CN 219 has
a bright rim with a detached bright spot just to the north and thethin, twisted dust-tail in CN 241 is remarkable as is the surrounding, faint, and very elongated bright halo. Some of the peculiar
objects in Fig. 5 deserve a closer inspection.
6

-19

3

gr/cm )

10

8

(10

【結果】	
  
他のHII領域に比べて	
  
・小さい	
  
・低質量(惑星質量程度)	
  ・平均密度は高い(10 	
  cm 	
  を超えるものも)	
  小さいものほど密度が大きい	
  3.3. Peculiar objects
Globulettesは分子ガスのより大きな塊から分離し
The Carina Nebula hosts a large number of clouds of very irregular shape like dark worm-like filaments and larger fragments,
たものかもしれない	
  such as the ”Defiant Fingers” described by Smith et al. (2004).
Some of these fragments are accompanied by smaller cloudlets
・薄い周囲が蒸発し高密度なコアだけが残った	
  with irregular shapes that most likely have eroded from the larger
ones. We did not include these objects in the list of globulettes.
・残ったコアが時間とともにより圧縮	  In Figs. B1-B6 we have marked some larger fragments, since
they might be important birth-places of globulettes (see Sect. 4).
	  The border-line between what we define as a globulette orareas containing globulettes. This object and j are denoted in
Smith et al. (2003). The images span: A 500 x 500 , B 400 x 400 , C
2000 x 3000 , D 300 x 300 , c 800 x 800 , E 700 x 1000 , F 2000 x 1000 , CN241
700 x 500 , G 3500 x 1800 , CN219 700 x 700 , j 1500 x 1500 .

!

the interaction with photons coming from the bright stars in thiscluster. It should be noted that there are also a large number ofround, dark objects in these areas. In area III, elongated objects
are more randomly oriented, particularly in the central part of theimage. An example is CN 93 with 60 M J in Field 18, an isolated
globulette seen in projection against the cluster Tr 14. However,both the direction of the tail and the bright rim indicates that theglobulette is influenced by some object east of the cluster core.
Globulettes just above and along the western extension of the
V-shaped dust lane (in the right part of this panel), as well asa group to the left in the panel point in the general direction of
the bright nebulosity surrounding ⌘ Carina and Tr 16. Several
O stars are spread over this nebulosity, and it is likely that their
combined radiation caused the shaping and orientation of theseelongated globulettes. In addition, we found the same predominant direction
in regions south of area III (not shown
5 of objects-­‐3here).

6

4

2

0
0

1

2

3

4

5

6

7

8

Radius (kAU)

Figure	
  6	
  (↑)

Fig.	
   6. Average density versus radius for all dark globulettes in
the Carina complex.

Carina	  Nebula	
  の Globulettes	
  のサイズと平均密度の関係	
  4. Discussion
None of the 288 globulettes coincides in position with any of the
YSO candidates listed in Povich et al. (2011) and Gaczkowski
et al. (2013). There are stars seen inside the boundaries of two
globulettes in optical images, namely CN 138 (Fig. 2) and theobject designated j in Fig. 5 in Smith et al. (2003). These stars
are likely to be foreground stars, since they show no sign of IR
excess judging from the Two Micron All Sky Survey (Skrutskie
et al. 2006) or existing Spitzer images. There are no obvious
proplyd candidates in our sample, except possibly for CN 219
with a jet-like extension (see Figure 5) and listed as a HerbigHaro object (HH 1011) in Smith et al. (2010a).
The fraction of objects with bright rims and halos is 39%,
which is large compared to findings from other H ii regions (DeMarco et al. 2006; Paper 1). In the central parts of the Carina
Nebula objects with bright rims even dominate over those with-perature using the task ‘combine’ in MAPPING, by multiplying by
11.35 K (Jy beam−1 )−1 .
To perform spectral analysis, single spectra have been extracted
from individual pixels (0.8 arcsec × 0.8 arcsec) within the cube.
These data have then been smoothed using a Gaussian weighting(full width at half-maximum, FWHM = the major axis of the synthesized beam ∼4 arcsec). Pixel spacings are equivalent to 0.5 ×
the major axis of the synthesized beam (∼2 arcsec). The typical
spectral rms noise in each pixel is 0.1 K. The analysis is restricted
to the isolated hyperfine component (F1 , F = 0,1 → 1,2) of N2 H+(1–0) unless otherwise stated. This is because the isolated component is expected to be optically thin (see Section 3.2 for moredetails).
The 3.2 mm continuum data have also been CLEANed using the
Hogbom algorithm. The synthesized beam is 4.2 arcsec × 3.1 arcsec, with a position angle of 17.◦ 3. The typical map rms noise is0.07 mJy beam−1 , estimated from emission-free regions. In thiswork, the continuum data are used for comparison only. A full
discussion and analysis of the continuum data will be provided inHenshaw et al. (in preparation).
Utilized throughout this paper is the 8 µm extinction-derived,
2 arcsec resolution, mass surface density map of BT12, as
modified by KT13 to include corrections for the presence ofthe near-infrared extinction-derived IRDC envelope. When direct+

3 R E S U LT S
3.1 Intensity distribution and moment analysis
Fig. 1 displays the results of moment analysis covering the
PdBI map. Moment analysis has been performed between 42 and
48 km s−1 (to incorporate the majority of emission in the average
spectrum), and above 0.3 K (the 3σ level). The left-hand panel ofFig. 1 compares the spatial distribution of the N2 H+ (1–0) integrated
intensity (black contours; zeroth-order moment) with the mass surface density, as derived in KT13 (colour scale). To highlight thedensest portion of the cloud, the (solid) contours are plotted from10σ (the dotted contour refers to the 5σ level). The central panelof Fig. 1 displays the VLSR map (first-order moment), and is shown
here between 44.5 and 47.0 km s−1 (this narrower velocity range has
been chosen to pick out the variation in velocity from the brightest
emission). The right-hand panel displays the velocity dispersion of
the N2 H+ (1–0) emission, or second-order moment, between 0.0
and 1.5 km s−1 .
It is evident from the left-hand panel of Fig. 1 that the N2 H+
is extended over a large portion of the cloud. This confirms theresult from Paper IV that the dense gas is extended over parsec scales in G035.39−00.33. The emission traces the morphology of the mass surface density very closely. The white crossindicates the position of H6, as determined in BT12 (position:The	
  dynamical	
  properties	
  of	
  dense	  ﬁlaments	
  in	
  the	
  infrared	
  dark	  cloud	
  G035.39−00.33⋆	  	
  Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on May 1, 2014

J.	
  D.	
  Henshaw,	
  P.	
  Caselli,	
  F.	
  Fontani,	
  I.	
  Jime	  ́nez-­‐Serra,	
  and	  J.	
  C.	
  Tan	  【背景】	  Infrared	  Dark	  Cloud	
  (IRDC)	
  
大質量星・星団形成の初期条件を研究するための”実験室”	
  
【観測】	
  
G035.39–00.33をPdBIでN2H 	
  (1–0)の高感度・高空間分解能観測	  ←	
  Figure	  1
	
  G035.39–00.33の全体像	
  等高線:	
  N2H+積分強度図	
  (左)	
  質量面密度	  (中)	
  速度場図	  (1st	  moment)	
  (右)	
  速度分散図	  (2nd	  moment)	
  Figure 1. Left: integrated intensity contours (black) of N2 H+ (1–0), overlaid on top of the mass surface density plot from KT13. The spectra have been
integrated between 42 and 48 km s−1 , focusing solely on the isolated (F1 , F = 0,1 → 1,2) hyperfine component. Contours increase from 5σ (dotted contour) insteps of 5σ (solid contours; where σ ∼ 0.1 K km s−1 ). The synthesized PdBI beam is shown as a white ellipse in the bottom-left corner (the filled black circle

The dynamics of filaments in G035.39−00.33

2867

Figure 6. Three-dimensional PPV cube showing the centroid velocity of the three velocity components: (cyan) F2a, (red) F2b, and (green) F3, observed acrossG035.39−00.33. Additional velocity components, C4, C5, and C6, are shown in purple, blue, and orange, respectively (see Appendix D for more details).
Unassigned data points are shown in black. The mass surface density map can be seen at the base of the cube, overlaid with N2 H+ (1–0) integrated intensity
contours. Contour levels increase from 5σ (dotted contour; where σ ∼ 0.1 K km s−1 ) in steps of 5σ (solid contours), as with Fig. 1. The vertical dotted line
corresponds to the position of H6.

at 3σ (where σ = 7 × 10−2 mJy beam−1 ) and increasing in stepsof 2σ . Overlaid are the extended 4.5 µm (green squares; Chambers
et al. 2009), 8 µm (red open circles), and 24 µm emission (red open
triangles; Carey et al. 2009), as well as the low-mass cores (yellow
squares) and high-mass cores (magenta squares) identified using
Herschel (Nguyen Luong et al. 2011). The left-centre, right-centre,
and right panels show the VLSR maps and spatial location of all
three velocity components, as deduced from the Gaussian fitting
routine. Overlaid on top of each map are arrows indicating themagnitude and direction of velocity gradients in two dimensions.
To achieve this, we have followed the analysis of Goodman et al.
(1993). By assuming that the centroid velocities of observed lines
take an approximately linear form:
VLSR = V0 + A"α + B"δ,

	
  

(3)

whereby "α and "δ are offsets in right ascension and declination(in radians), least-squares minimization can be used to estimate

values of A and B, using MPFIT2DFUN (Markwardt 2009). The velocity
gradient, ∇v, can then be calculated for a cloud at distance D, using(Goodman et al. 1993):∇v =(A2 + B 2 )1/2,
Dand its direction, %∇v (direction of increasing velocity, measured
east of north), using%∇v = tan−1

!

"
A
.
B

(5)

This method has been adapted following the procedure outlined
in Caselli et al. (2002b) to calculate multiple gradients within agiven region with a good determination of VLSR (seven contiguous
positions with significant measurements of VLSR ). In the case of the

Figure	  12	
  (Bottom)	
  (←)
	
  
観測された速度勾配の2つの解釈	  Figure 12. Top: a close-up image of the region surrounding H6. The size and direction of the arrows correspond to the magnitude and direction of the velocity
gradient (pointing in the direction of increasing velocity). Dark green and dark red arrows refer to the velocity gradient calculated at each point, using velocities
from the surrounding 27 pixels (note – this is used for a spatial representation only, and not used in the analysis); light green and light red arrows refer to thevelocity gradient calculated from the surrounding 38 pixels, for F2b and F3, respectively (see Section 4.2 for details). The symbols have the same meaning as(4)

MNRAS 440, 2860–2881 (2014)

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on May 1, 2014

Downloaded from http://mnras.oxfordjournals.org/ at National Astronomical Observatory of Japan on May 1, 2014

【結果】	
  
複数の超音速(σNT/cs~1.5)フィラメント(F2a,	
  F2b,	
  F3)をスペクトル方向・空間方向に同定	
  単一鏡のマッピング観測で見えていた大スケールでの速度勾配はサブ構造によるもの	
  各フィラメントでの速度勾配	
  The dynamics of filaments in G035.39−00.33 2871
・大局的には小さい	
  (<0.7	
  km/s/pc)	
  ・局所スケールでは大きい	
  (1.5–2.5	
  km/s/pc)	
  →力学的なプロセス	
  高密度ガスの運動は	
  ・フィラメント内の高密度コア	
  ・すでに形成された星からのOutﬂowフィードバック	  Figure	
  6	
  (↑)
	
  
色は3つのフィラメント成分に対応(cyan–F2a,	
  red–のいずれかの影響	
  F2b,	
  green–F3)
・infalling	
  material	
  (降着率	  ~	
  7±4	  e-­‐5	
  Msun/yr)	
  ・expanding	
  shells	  (運動量	  ~	
  24±12	
  Msun	
  km/s)	
  
2つのシナリオを区別するにはさらなる分子輝線観測
が必要	
  

AN	
  OPTICAL	  SPECTROSCOPIC	  STUDY	
  OF	
  T	
  TAURI	  STARS.	
  I.	
  PHOTOSPHERIC	  PROPERTIES	
  	
  Gregory	
  J.	
  Herczeg	
  and	
  Lynne	
  A.	
  Hillenbrand	
  

【背景】	
  
星の質量・年齢はHR図上の位置から推定されるが限界がある	
  ・観測の不確かさ	  ・星周現象に関わる不確かさ	
  
【観測】	
  
近傍の281のTタウリ型星を可視分光観測	  正確なスペクトル型と減光量を求める	
  
【特色】	
  
減光量,	
  降着連続波フラックス,	
  スペクトル型を同時に測定できる	
  
【結果】	
  
新たなスペクトル型は過去とだいたい同じだが	
  ・TWAのメンバーについては1–2サブクラスだけ晩期側に	  ・いくつかのよく知られたおうし座分子雲のメンバーについても異なる	  The Astrophysical Journal, 786:97 (32pp), 2014 May 10

Herczeg & Hillenbrand
Table 3Spectral Indices

NameContinuum Range(C)

Band Range(B)

FeatureG-band
R5150

4550–4650
4600–4700

4150–4250
5050–5150

G-bandMgH

TiO 6250
TiO 6800
TiO 7140
TiO 7700
TiO 8465

6430–6465
6600–6660, 6990–7050
7005–7035
8120–8160d
8345–8385

6240–6270
6750–6900
7130–7155
7750–7800
8455–8475

TiO
TiO
TiO
TiO
TiO

x

SpT

Range

Zero-ptrmsa

C/B−25.6 + 29.96x−29.7 + 28.3x

G
K0-M0

G0
K0

∼1
1.0

(M0–M4)K5–M0.5
M0–M4.5
M3–M8
M4–M8

M0
K0
M0
M0
M0

F (5100) Fline (4650) b
F (4650) Fline (5100)
log ( CB − 1)

C/B
log ( CB − 1)
C/B
C/B

3.20–5.43x + 1.73x 2−15.37 + 19.77x
4.36 + 6.33x + 1.57x 2
0.11 + 2.27x−0.74 + 4.21x

...
...
0.42c
0.21
0.18

←	
  Table	  1
	  スペクトル型の導出に用いるスペ
クトル指数	
  波長幅の組(C,	
  B)でのフラックスの中間値からスペクトル指数を計算する	
  Notes. All spectral indices are calculated from the median flux in the spectral range.
a SpT
The
Astrophysical(32pp),
2014 May
Herczeg & Hillenbrandrms calculatedJournal,
from the 786:97
literature
SpT (R5150,
TiO106800), Kirkpatrick SpT (TiO 7140), and Luhman SpT (TiO 7700, 8465).
b The observed flux ratio is divided by the same ratio obtained from a linear fit to the 4650–5300 Å region, see the text.
c 0.3 between M1 and M4, 0.8 earlier than M1.
prominent
TiO bands are present in our red spectra (see←	  FourFigure	  2d The 8120–8160 Å continuum range should be used only for spectra that are corrected
2 and3). Figure 3 compares the spectral types
	   Figurefor telluric
H2 OTable
absorption.

3.1. Quantification of Spectral Indicesand four spectral indices for(Red側)	
  the PMSU, Kirkpatrick, and Luh実際のスペクトルの例	
  mansamples.
For
stars
earlierthan M5,
the Luhman
relationScatter in these quantified relationships
are causedby metal黄色い部分が指数の計算に用いられた波長域	  ship between SpT and TiO depth for young stars was intended

licity and gravity differences between stars. The metallicity
to follow
theassociations
Kirkpatrick iset uniformal. (1991)
results.
However,of nearby
young(e.g.,
Padgett
1996; forIn this subsection, atlases of low-resolution optical spectra
objects
from
M0
to
M3,
the
Luhmanspectraltypes
are ∼0.5
Santos et al. 2008; D’Orazi et al. 2011). Gravity differences
are used to establish a set of quantified spectral indices for
subclasses
later
than
the
medianKirkpatrick
spectraltype
between 1–10 Myr may be significant and are discussed but are(TiOyoung stars. The following descriptions are divided by spectral
7140investigated.
and TiO-6200 spectral indices). For the spectral types
not fully
type, each of which is sensitive to a different spectral index.
later
than
M5, subsections,gravity differences
betweendwarfs andIn the following
we describe
howfieldtheseMspectral
Spectral typing of young stars has typically relied on eyeballpre-main-sequenceMdwarfslead
totheLuhmanspectralindices are used to measure spectral types. Each spectral index istypes
comparisons to a sequence of spectral standards. While thatbeingtoslightly
earlierthantypes
the medianobject (as dissensitive
different
spectraland is Kirkpatrick
discussed separately,
approach can be very accurate, a quantified approach allows forcussedbyLuhman).
ForM
dwarfsearlierthan
M4,
we adopt
beginning with the coolest stars in our sample.
greater consistency between different sets of eyes. A quantified
the spectral type sequence of Kirkpatrick, which may introduce
approach also readily accounts for accretion and extinction bya small offset
betweenourspectraltypes and Luhman spec3.1.1.
SpectralTypesof M Stars
calculating over a grid of values to find a best fit solution.
tral types. For M dwarfs later than M4, we adopt the spectral
The full set of spectral indices discussed in this paper isThemajorityof starsin our sample
M stars. TiO/VO
At opticalbands
type
sequenceof Luhman.
Severalare
additional
listed in Table 3. By design, our focus is on K and M stars.
wavelengths,
M
stars
are
easilyidentified
fromthepresenceof (seeare detected at blue wavelengths and are not well studied
The M-dwarf spectral types rely on the depth of TiO and VO
strongTiO absorption
Kirkpatrick
et al.does(1991,
Figure4). While bands.
our initialapproachnot1993),
considerabsorption bands (hereafter referred to as TiO), which start tohereafter
establishedgridareof calculated
M-dwarf spectral
these Kirkpatrick,bands, the final
spectral atypes
from a best
become detectable at ∼ K5. For K dwarfs, a spectral type index
type fit
standards
from
field
stars.
Reidetal.
(1995),
hereafter
5
to a spectral sequence using the full optical
spectrum.
is developed based on the 5200 Å absorption feature, which is
PMSU,M4–M8.
quantified
relationships
betweenspectral
type
and
the
Objects later than M4 have spectral types assessed
a combination of MgH, Mg b, and Fe i (e.g., Rich 1988). The
depthfromof TiO
at 7100fromÅmoderate
resolution
opticalfrom
the bands
TiO 7700
andÅ8500
bands, witha conversion
spectral typing of BAFG stars relies on a visual comparison ofspectra
based
on
the
Kirkpatricketal.
(1991)
sequence.
spectral index to spectral type calculated from the sequence of
the G band and absorption in H and Ca lines.
Luhmanet al.
hereafter
objects (1999);
providedLuhman
by Luhman.
An(2003),
uncertaintyof 0.2Luhman
subclasses is
The quantification of spectral typing provides an objective(alsoassigned
includes,basede.g., onLuhman2004,2006),
recognized
thatsubclassforthe change in feature strength versus
and repeatable method to measure spectral types with precision.
pre-main-sequencestars, deviation
the depth offeaturesand on the standard
in TiOthe fits
to thedeviates
Luhmanfrom
objects.
The quantified prescriptions of M dwarfs are similar to thosedwarfThisstars
because ofgravity
Gullbring
et al.
uncertainty
is lowerconsistent
with(seethatalso
assignedby Luhman.
Figure 2. Red
spectralsequence
K3–M5.5.
Regions
used forTiO band 1998).
of Slesnick
et al.
(2006)
and from
Riddick
et al.
(2007),
while
Luhman developed a spectral type sequence for young
indices are highlighted in yellow. Selected gravity-sensitive absorption in
5 Thelater
the prescriptions
for earlier spectral types are similar to those
M dwarfs
thanindex
M5 wasbased
on a hybridof fieldgiTiO 7140
developedby Slesnick
et al. dwarf(2006). andOur definition
CaH λ6382, Fe i λ6497, and the Na i λ8189 doublet are highlighted in blue.
uses asinceslightlyTTSs
different
developedby,e.g.,Wortheyetal.
(1994)
and
Covey
et
al.
(2007).
antstars,
arecontinuum
typicallyregion.
luminosity class IV. For stars(A color version of this figure is available in the online journal.)
The spectral indices described here are calculated from theearlier than M5, Luhman relied on the Kirkpatrick grid along
median flux in the relevant spectral range and are tailored to low
with the Allen & Strom (1995) red spectroscopic survey of stanspectral resolution. These spectral indices are then combined
dards. Although the Luhman spectral sequence is well acceptedwith an accurate flux calibration and blue spectra to measure
and widely used, it has no standards or quantified conversions
accretion and extinction simultaneously (see Section 4). In
between spectral index and spectral type. As a consequence,
several cases, the spectral index is changed to a log scale tospectral types based on the Luhman method are likely less pre-