THE STAR FORMATIONNEWSLETTER
No. 259
塚越崇 (茨城大)
22 Aug. 2014

Sejong Open Cluster Survey(SOS). III. The Young OpenCluster NGC1893 in the HII region W8
B. Lim et al.
• HII領域近傍のクラスターNGC1893の可視光観測• SOSプロジェクトの一環(Paper0を踏襲)• NGC1893は良く知られたクラスター領域。銀河外縁部の星形成9を知る上で重要な領域• AZT-22 1.5m望遠鏡

Figure 1. Finder chart (left) and an H↵ image (right) for NGC 1893. The size of the circles in the finder chart is proportional to tbrightness of the stars. The position of stars is relative to the O-type star HDE 242935 (↵ = 05h 22m 46.s 5, = +33 250 1100 , J200
A square (solid line) outlines the observed region in this work. Another two boxes (dashed line) display the Chandra ACIS fields of vi(Caramazza et al. 2008). The right-hand side panel shows the position of 5 O-type stars and two emission nebulae.

• Maidanak Astronomical Observatory in Uzbekistan

Table 1. Photometric data and spectral type for 24 bright stars• UVBI-bands, Hα

ID

V

B

V

U

B

reference

Spectral typereference

O7V(n)z

Sota et al. (2011)

O4.5V(n)((fc))
O7Vz
O6.5V((f))z
B0V

Sota et al. (2011)
Sota et al. (2011)
Sota et al. (2011)
Negueruela et al. (2007)

B1III

Massey et al. (1995)

B1.5V
B0.5V
O8

Massey et al. (1995)
Negueruela et al. (2007)
Hoag & Applequist (1965• 距離: ZAMSのﬁtting=>3.5 0.3kpc(ペルセウスアーム)

B0.2V
B1.5III

Massey et al. (1995)
Massey et al. (1995)• 年齢: HR diagramより

B0.5V
B0.5V

Massey et al. (1995)
Massey et al. (1995)

B0.5V

Massey et al. (1995)

BD +33 1025
10.31
0.26
-0.73
Hoag et al. (1961)
BD +33 1028
9.15
1.21
0.95
Hoag et al. (1961)
obtainedfromthe spectral
type-e↵ective
temperature
relaHDE242855
10.65
0.06
-0.35
Cu↵ey
(1973a)tion from
5 in Paper9.05
0. For 0.27
the temperaturescale of& Morgan (1955)
HDEtable
242908-0.72
Johnsonthe PMSmembers
1.4, the -0.66
relation Johnson
between& Morgan (1955)
HDE
242926 with (V9.35I)0 60.32
HDE
2429359.43
0.20-0.73et al.
Hiltner(1956)(VI)
e↵ective temperatureof Bessell(1998)
0 and
HDE 243018
10.94
0.21
-0.77
et al. (1961)
was adopted,
and anotherrelation
between(V HoagI)0 andTYC 2394-162-1
10.79
0.25
Hoag et al. (1961)
temperature(Bessell 1995)
was applied
to0.14the remaining
TYC
2394-450-1
11.47
0.17
-0.66
Masseycool PMS members. For the other MS members, weaver-et al. (1995)
TYC 2394-503-1
11.08
0.58
0.03
Cu↵ey (1973a)
aged the temperature derived from spectral type-e↵ective
TYC 2394-539-1
11.61
0.22
0.12
Hoag et al. (1961)
temperature
relation and 12.00
colour-temperature
relationswith
TYC 2394-629-1
0.16
-0.62
Massey et al. (1995)
an appropriate
weight.
Thebolometriccorrectionsforallal. (1961)
TYC 2394-1141-1
11.04
0.44
-0.46
Hoag etthe members
were carriedout using
5 in Paper
TYC 2394-1214-1
10.15
0.51 table-0.49
Hoag et0.al. (1961)
We present
the HRD of NGC
in Fig. 1.17
8. We notethat
TYC 2394-1469-1
11.07 18931.28
Cu↵ey(1973a)4 B-type
were classified
giants by
Masseyet etal.al. (1961)
TYCstars2394-1500-1
10.70 as 0.30
-0.62
Hoag
2394-1502-1
Massey(1995).TYC
Negueruela
et al.11.68(2007) 0.19
presented-0.55
the morereli-et al. (1995)
TYC 2394-1594-1
11.86
0.37
Cu↵ey
(1973a)able spectral
type and luminosity
class for0.22
2 B-type
giant
TYC 2394-1623-1
0.26
et al. (2000)
stars (HDE
243018 and 10.86
TYC 2394-1141-1).
We HøgadoptedTYC 2394-1691-1
11.34
0.26 Table-0.64
Masseytheir classificationsin thiswork (see1). Thelumi-et al. (1995)
TYC 2394-1738-1
11.19
0.25
-0.67
Hoag et al. (1961)
nosity class of two B-type stars, TYC 2394-450-1 (B1.0III,
TYC 2394-1744-1
11.78
0.23
-0.45
Hoag et al. (1961)
Mbol =
4.988, log T= 4.395)0.28
and TYC
2394-1502-1
TYC 2394-1912-1 e↵ 11.23-0.62
Hoag et al. (1961)• photometric diagramから基本量を導出• PMS同定: 既知のPMSと比較し検出率を評価。205PMSが検出された。

• 赤化: 3.6mumとVbandデータから空間分布も導出(B1.5III, Mbol =4.460, log Te↵ = 4.336), is needs to be• IMFはSalpeter/Kroupa IMFとコン
システント=>outer gal.の星形成は近傍と同じ• Accretion rate•82天体でUV超過を検出•

Macc=1.6E-8 Msun/yr (M<2M⦿)•

数M⦿以上で低質量に比べて10倍になる？

• Triggered star formation•

Sim129,130: headとtailを持つ•O型星からの距離↗で星年齢↘の傾向•

O型星のexpanding timeで説明出来る

VLA observations of ammonia in high-mass star formation
regions
X. Lu et al.

12

Lu et al.

• 大質量形成領域に対するVLAを使ったNH3輝線サーベイ• NH3: 高密度領域のトレーサーで、温度や光学的厚みも導出可能• これまでは単一鏡で分解能が悪い。
高分解能観測を行うことが重要• 3”x3”分解能を達成• 分布は3つのカテゴリに分けられる• ホットコア• NH3散逸• フィラメント
Fig. 1.— Rotational temperature and IRAC three-color images of three representative sources: from top to bottom, a hot core source, a• 温度分布とIRの対応•一致するケース=>原始星輻射起源、power-law分布を示す。
主にホットコア天体で見られる。

•

非一致のケース=>外部熱源からの輻射16

Lu et al.

dv

TFig. 8.— Continued
Fig. 8.— Continued• Dense Clumpについて
NNH3

Mvir/Mcore

Mcore

•

clumpﬁndによる同定•

密度や温度はホットコアが一番大きく、散逸天体が最も小さい•

ビリアル比はホットコアで<1となる=>ホットコアが外の二つのクラスより進化しているからAmmonia in High-mass Star Formation Regions

T-dv• クランプ物理量の相関
•

•

17

Σ-dv

Ammonia in High-mass Star Formation Regions

35

T-dvの弱い相関は、温度は原始星からの輻射、線幅は乱流と思うと説明出来る(Mstar/Mcore=const.を仮定)
R-dv

LIR-T

Σ-dvに良い相関。3つのカテゴリが同じ線上になる。

• フィラメントのcollapse••

フィラメントに等間隔のクランプあり=>フィラメント内の重力分裂での間隔と思われる
IRAS19074+0752のケース•

長さ1.3pc•

クランプ間隔0.3pc

Fig. 7.— Correlations between core properties. Classes of cores are shown in di↵erent symbols and colors. a). The FWHM linewidth vs.
the rotational temperature. b). The FWHM linewidth vs. the NH3 column density. c). The FWHM linewidth vs. the core radius. d).
The cluster luminosity vs. the rotational temperature.

TABLE 1
Summary of the VLA observations.
Index

Project

Date

Config.

uv Distance(k )NumberPhase Calibrators

1
2
3
4

AZ0103AS0708AS0747AC07331997 Dec 13
2001 Jul 23
2003 Jan 16
2004 Jun 14/15

D
C
DnC
DnC

2 - 772 - 268
2 - 167
2 - 167

1
1
5
42

5
6

AC0733AW0666

2004 Jun 212005 Nov 12DD

2 - 81
2 - 81

12
1

1923+210
1829 1061829 106, 1849+005
0530+135, 1743 038,
1851+005, 1925+211
1851+005, 2322+509
1832 105=> 等温ガスシリンダーの収縮で説明可能。(非圧縮だと間隔が狭い)
Fig. 8.— Continued

RMS Level(mJy beam 1 per channel)
8.5
7.0
2.2
2.4
2.2
2.2

Location and sizes of forsterite grains in protoplanetarydisks - Interpretation from the Herschel DIGIT programme
K. M. Maaskant et al.: Location and sizes of forsterite grains in protoplanetary disks

et al. 2008) are generally iron poor. Iron fractions (by number)in olivine crystals found in the Itokawa asteroid and in ordinarychondrites can be as high as ⇠ 30% (Nakamura et al. 2011).
This may reflect the lack of large-scale mixing (Shu et al. 1994;
Bockelée-Morvan et al. 2002; Gail 2004; Jacquet 2014) and indicates the importance of crystal formation (i.e. local heating
events) associated with the inner and outer disk region (Urey
1967; Huss et al. 2001; Harker & Desch 2002; Desch et al. 2005;
Ábrahám et al. 2009; Morlok et al. 2010; Tanaka et al. 2010).
The transition from amorphous to crystalline silicate dust
leads to strong changes in the optical response function of thematerial (Hallenbeck & Nuth 1998). Crystalline silicates haveseveral sharp solid-state features. A detailed analysis of thesefeatures provides a wealth of information on the properties of
the crystal grains (Koike et al. 2003; Suto et al. 2006). Because
these properties can potentially be linked to the specific circumstances required to crystallise the dust grains, crystalline silicatesallow us to probe the physical processes that change the chemical compositions in the disk.
The width and peak position of the 69 µm feature of
forsterite is sensitive to the temperature (Bowey et al. 2002;
Suto et al. 2006; Koike et al. 2006). In addition, the peak position of the feature becomes red-shifted for higher fractions ofiron (Koike et al. 1993, 2003). Herschel observations using the
Photodetector Array Camera and Spectrometer (PACS) have allowed us to study the 69 µm feature with high spectral resolution (Sturm et al. 2010; Mulders et al. 2011; de Vries et al.
2012; Sturm et al. 2013). A detailed analysis of the 69 µm feature shapes in protoplanetary disks is presented in Sturm et al.
K.
and sizes
sizes ofof forsterite
forsterite grains
grains ininprotoplanetary
protoplanetarydisks
disks
K.M.
M. Maaskant
Maaskant et
et al.:
al.: Location
Location and(2013). These authors found the following results. Most of theforsterite grains that give rise to the 69 µm bands are found tobe ⇠ 100 200 K and iron-poor (less than ⇠2% iron), which
supports the hypothesis that the forsterite grains form through Fig. 1. Spectra of all sources with 69 µm detections. The continuuman equilibrium condensation process at high temperatures. The subtracted spectra are scaled relative to the strongest band. The objects
large width of the emission band in some sources has been inter- are sorted by the F30 /F13.5 ratio.
preted to indicate the presence of forsterite reservoirs at different
temperatures. It was also concluded that any model that can explain the PACS and the Spitzer IRS observations must take the 2. Observations
effects of a wavelength-dependent optical depth into account.
In this section, we analyse the short- and long-wavelength
Radiative transfer modelling has only been performed for forsterite features originating from protoplanetary disks. We
HD 100546 (Mulders et al. 2011) to study the consistency be- compare the feature strength ratio I23 /I69 with the forsterite
HD141569
tween the relative mid-infrared band strengths and the informa- 69 µm shape and check for their consistency.
In addition, we
64tion contained in the23
shape ofthe 69 µm band. It was found that present trends between the I23 /I69 ratios and the millimetre (mm)the best solution is given by a model with a high abundance of luminosities of the disks. Note that throughout this paper, the inforsterite located in the inner edge of the outer disk. Because
tegrated intensities are given by I [erg cm 2 s 1 ], the monochroHD141569
protoplanetary disks are highly optically thick media, radiative matic flux by F [erg cm 2 s 1 µm 1 ], and the luminosity bytransfer modelling is essential for investigating optical-depth ef- L [erg s 1 ].
fects and for simultaneously studying the effects of grain growth,
protoplanetary gaps and, radial mixing.

K. M. Maaskant et al.

• フォルステライトがどこで形成されているのか？

• 結晶化には>1000K必要だが、円盤外側でも検出有り• radiative transferによる赤外スペク
トル(23,64mum)の計算• 観測データ• 63µm: Hershcel/PACS,
23µm:Spitzer/IRC• ミリ波が弱いやつで63µm検出率が高い。穴の有無に相関無し(GroupI - II)• 69µmの検出はGr.I>Gr.II。I /I では
1、2桁の違いがある• Radiative transfer, dust model•

シリケイト化合物80%、アモルファスカーボン20%•

実験室optical constantsでオパシティへ•

放射輸送コードMCMaxを使用•

色々な円盤モデルに対してベンチマークテスト

K. M. Maaskant et al.: Location and sizes of fo

10
mm

co
tio
is
pla
by
is
mo
na
Rin
du
ios
the

• HD141569に対するモデリング•

ガスリッチ円盤でデブリ円盤直前段階•

強い62µm featureを示していたが24µmは見えない•

文献とSEDから円盤構造を推測、3つのダストモデルでプロファイルを計算••

K. M. Maaskant et al.: Location and sizes of forsterite grains in protoplanetary disks
SEDを再現する為に円盤を4つに分けた

良くフィットするのは•ModelB••iron-rich, ~1µm, ~40K

ModelC•iron-free, ~10µm,~100K

Fig. 6. Top: calculated opacities of 1 µm forsterite grains at 50 K and
295 K. Bottom: calculated opacities for 1 and 10 µm forsterite grains(T = 150 K).

Ta• ベンチマークテスト
でパラメータ依存性を調べている。

!

• 強い69µmかつ弱い23µmを作るには•

grain growth•transitional disk•

optical depth eﬀect

GM AurSz ChaTW HyaCS ChaCHXR22E
Sz18
Sz27
Sz45
Sz84
RX J1615
Oph22
Oph24
SR 21aISO Oph196
DoAr 44
Ser29
Ser34
RX J1842.9
RX J1852.3(1.06 ± 0.02) ⇥ 10
(2.6 ± 0.1) ⇥ 10 12(2.39 ± 0.04) ⇥ 10 11(5.23 ± 0.09) ⇥ 10 12(1.1 ± 0.1) ⇥ 10 14(3.8 ± 0.2) ⇥ 10 13(1.83 ± 0.08) ⇥ 10 12(2.83 ± 0.02) ⇥ 10 12(7.0 ± 0.1) ⇥ 10 13(2.24 ± 0.05) ⇥ 10 12(1.9 ± 0.09) ⇥ 10 13(3.2 ± 0.2) ⇥ 10 13...
(3.54 ± 0.01) ⇥ 10 13(9.4 ± 0.1) ⇥ 10 12(1.3 ± 0.1) ⇥ 10 14(1.02 ± 0.05) ⇥ 10 13(3.41 ± 0.08) ⇥ 10 12(5.67 ± 0.09) ⇥ 10 12(2.1 ± 0.07) ⇥ 10
< 4.7 ⇥ 10 14(4.52 ± 0.07) ⇥ 10 12(4.7 ± 1.1) ⇥ 10 13(4.4 ± 0.6) ⇥ 10 15(2.5 ± 0.6) ⇥ 10 14(7.8 ± 0.6) ⇥ 10 14(4.6 ± 0.1) ⇥ 10 13(5.59 ± 0.09) ⇥ 10 14(2.5 ± 0.3) ⇥ 10 13(5.4 ± 0.1) ⇥ 10 14(1.08 ± 0.03) ⇥ 10 13...
(8.8 ± 0.6) ⇥ 10 14(2.66 ± 0.08) ⇥ 10 12< 1.7 ⇥ 10 15(7.8 ± 0.7) ⇥ 10 15(4.6 ± 0.6) ⇥ 10 13(1.01 ± 0.07) ⇥ 10 12

On the gas content of transitional disks: a VLT/X-Shooter
study of accretion and winds
C.F. Manara et al.

• 遷移段階円盤の穴内のガスを調べる• 穴形成メカニズムの切り分けに重要• 広帯域スペクトル観測から、中心星、質量降着、ウィンドを調べる• ESO/VLT X-Shooter: 300-2500nmをカバー• 22天体のTD: Maccの大きいものにバイアス
フィットする(Manara+2013b)• スペクトルはモデルで良く再現可能• 星のテンプレート+等温の水素slab(降着)• J-Kの超過から、TD or pTDを切り分け• 円盤風• ウィンド起源の低速成分を17天体で検出(8.8 ± 0(1.04 ± 0
< 8.0
(5.9 ± 1(2.0 ± 0(2.8 ± 0

Notes. Fluxes are reported in the format (flux ± err) multiplied by the order ois very uncertain, thus we do not report these values for this object.

• SpT, Av, Laccをパラメータにし、スペクトルを• 630nm [OI]による検証(7.9 ± 0< 4.0(2.2 ± 0(2.2 ± 0(9.1 ± 4(7.9 ± 0(3.2 ± 0(2.9 ± 0(2.7 ± 0(6.9 ± 1(2.8 ± 0(7.7 ± 0

Name
FAccretion
properties[erg
s cm ]•[OI] 630
1
2

LkH↵330DM Tau
LkCa 15
GM AurSz ChaTW HyaCS ChaCHXR22E
Sz18
Sz27
Sz45
Sz84
RX J1615
Oph22
Oph24
SR 21
ISO Oph196
DoAr 44
Ser29
Ser34
RX J1842.9
RX J1852.3(6.1 ± 1.9) ⇥ 10 14(9.1 ± 1.8) ⇥ 10 15(4.6 ± 1.9) ⇥ 10 14(3.9 ± 1.0) ⇥ 10 14(2.4 ± 0.9) ⇥ 10 14(9.0 ± 0.8) ⇥ 10 14(2.8 ± 1.5) ⇥ 10 14< 8.0 ⇥ 10 16< 2.6 ⇥ 10 15(5.4 ± 0.3) ⇥ 10 14(1.8 ± 0.5) ⇥ 10 14(2.4 ± 0.2) ⇥ 10 15(1.7 ± 0.8) ⇥ 10 14< 7.3 ⇥ 10 15(3.0 ± 1.3) ⇥ 10 14< 6.4 ⇥ 10 14(1.4 ± 0.4) ⇥ 10 15(2.6 ± 1.2) ⇥ 10 14< 4.2 ⇥ 10 16(4.1 ± 0.8) ⇥ 10 15(5.0 ± 0.8) ⇥ 10 14(2.6 ± 1.3) ⇥ 10 14

v0,[OI] 630[km/s]
7.7
-4.3
-8.7
-2.2
-4.7
-4.8
-4.4
...
...
-7.6
-4.8
-6.5
-8.7
...
-13.0
...
-20.4
-7.7
...
-7.3
-5.5
-3.8

FWHM[OI] 630[km/s]
24
38
43
42
30
24
37
...
...
49
50
44
49
...
67
...
68
50
...
47
43
36

•

MaccとRinに相関無し=> 星に降着するガスとgapやhole形成メカニズムには依存関係無し•

IR classによる違いも無し。降着に関しては共通•

サンプル選定の影響はあるかも• Wind properties••

Fig. 8. Logarithm of the mass accretion rate vs inner hole sizefor our sample. Di↵erent symbols are used to distinguish themethods used in the literature to derive the size of the innerhole. Blue squares are adopted when this has been derived using[OI]強度とRinの依存性無しでほぼ一定
IR-SED fitting, while red circles when the values are derived
from resolved mm-interferometry observations. Downward arIR classによる違いも無し。[OI]生成プロセスはCTTSrows are upper limits. The two lowest points are, from left
Notes. Fluxes are reported in the format (flux ± err) multiplied by the to right, CHXR22E and Ser29. The object at Rin =25 AU andorder of magnitude.
と同様と思われる。
log Ṁacc = 9.8 is Ser34.

• TD v.s. pTD•

analyzed with a similar methodology. For these reasons we select as a comparison sample the objects studied by Alcalá et al.
(2014). These are located in the Lupus I and III clouds and haveages ⇠ 3 Myr, similar to the objects in our sample. The analysis
降着、ウィンドどちらも違い無し
of that sample was carried out with the same methodology as
the one we used here. We show in Fig. 9 the logarithmic relation between Ṁacc and M⇤ for these two samples. Our data arereported as blue circles, while data from Alcalá et al. (2014) as
green diamonds. The solid line on the plot is the best fit relation from Alcalá et al. (2014), and the dashed lines represent thedispersion of 0.4 dex around this best fit relation. The typicalerrors on the quantities are shown as a black cross. We see that
⇠80% of the TDs have values of Ṁacc consistent with the values found by Alcalá et al. (2014) for Lupus objects of the sameM⇤ . Therefore, for these objects we do not see any di↵erence inthe accretion properties with respect to those in cTTs. We also
perform on these two samples a Kolmogorov-Smirnof statistical
test (K-S test). When restricting to objects in the same M⇤ range
the probability that the two samples are drawn from the same di-

They suggest that results di↵ering from those found in previous
works may arise from di↵erent sample selection and/or di↵erent
methods to estimate Ṁacc . To avoid this possible methodologicalbias, we have shown here only the comparison between our sample of TDs and the sample of cTTs in Lupus, which is analyzed
in the same way as our objects. We stress here again that our TDs
have been selected to be mostly strong accretors, thus our sample
selection is not representative of all the TDs and that we do notderive conclusions for the whole TD population. Nevertheless,our results prove that there are TDs that accrete at the same rateof cTTs.
The two main outliers in Fig. 8-9 are the object with an upper limit on Lacc , Ser29, and CHXR22E. The lower intensity ofaccretion for these targets implies that the gas density in the inner disk is substantially depleted with respect to the one of cTTs.
These objects do not have any peculiar property reported in theliterature. From Fig. 9 we note that in the same range of M⇤ ofthese objects there are other TDs with values of Ṁacc comparable or even higher than cTTs. Therefore, these objects are not

The First Very Long Baseline Interferometry Image of 44
GHz Methanol Maser with the KVN and VERA Array (KaVA)
N. Matsumoto et al.

• KaVA(VERA+KVN)の初期成果論文• 44GHz class Iメタノールメーザー• 大質量星 IRAS18151-1208 MM2• 単一鏡で44GHzメタノールメーザー検出• 最長基線長2300km=>分解能2.7x1.5mas• 広がったメタノールメーザーを検出• コア中心には無い• 十分な分解能でclass Iメタノールメー
ザーの輝度温度を求めた最初の例• サンプルを増やす&proper motion観測へと繋げる