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• PMS同定: 既知のPMSと比較し検出率を評価。205PMSが検出• 赤化: 3.6mumとVbandデータから空間分布も導出• 距離: ZAMSのﬁtting=>3.5 0.3kpc(ペルセウスアーム)• 年齢: HR diagramより

reference

BD +33 1025
10.31
0.26
-0.73
Hoag et al. (1961)
O7V(n)z
Sota et al. (2011)
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) O4.5V(n)((fc)) Sota et al. (2011)
HDEtable
242908-0.72
Johnsonthe PMSmembers
withLeft-hand(V9.35I)0 6
1.4,
the -0.66
relation
between& of
HDE
0.32
JohnsonMorgan(1955) stars
O7Vz
Sotacontour
et al. (2011)
Figure 5. Reddening
mapof242926
NGC 1893.
side
panel:
spatial distribution
the observedin the V band andmap
HDE
2429359.43(= 0.20
-0.73
Hiltner(1956)
O6.5V((f))zSota
et al. 2(2011)
I)ande↵ectivetemperatureofBesselletal.
(1998)
0densityof the di↵erential(Vsurface, where)isthedi↵erenceinthesurfacedensityofstars
withina2⇥2arcminarea
V[3.6]
HDEThe2430180.21
et inal.the(1961)
B0V of Fig. 1. Dash-dotted
Negueruela
et al. (2007)
wasadopted,andbetween(V boxHoag
I)
0 andin the 3.6 µm andV bands.
field anotherof view10.94
isrelation
the sameas the-0.77
solid
shown
left-hand side panel(blue),TYC 2394-162-1
10.79
0.25
et al.
temperature(Bessell1995)
wasto0.14the= remaining
dashed (green), solid(orange),and dotted(red)linesapplied
represent
5,Hoag10, 15,and(1961)
20 per area, respectively. Middle panel: correlation
TYC
2394-450-1
11.47
0.17
-0.66
Masseyetal. (1995)
B1III members within
Masseyet al.
(1995)
betweenandcoolmeanPMS
reddeninghE(B For
V )i,
where
theMS
mean
reddeningwe
wasobtainedfrom the early-typea given2⇥2
members.
theothermembers,
averTYC
2394-503-1
11.08 the0.58
0.03early-typeCu↵eymembers(1973a) within a given area. The vertical bar represents
arcmin2 area. The
sizetheof circles
is proportional
numberaged
temperature
derivedtofrom
spectralof type-e↵ective
TYC
11.61
0.22 statistics
0.12 is Hoaget al. (1961)
the standard deviationof the2394-539-1
reddening,
while
the Poisson
assumedthe error in the. The solid line represents antemperature
relation andcolour-temperature
relationswithettoal.be(1995)
TYC
2394-629-1
12.00+ 0.418.
0.16 Right-hand-0.62 side
MasseyB1.5V
Masseyet al.side(1995)
adopted relation,an
E(B
V
)
=
0.012(±0.002)
panel:
spatialvariationofreddening.
in theleft-handappropriate
weight. The
corrections
for etallal. (1961)
11.04bolometric-0.46
B0.5V solid (orange),
Negueruela
et al. (2007)
panel was converted to TYC
E(B 2394-1141-1
V ) using the relation
in0.44the middle
panel.Hoag
Dash-dotted(blue), dashed (green),
dotted (red)the members
were carriedout using table
5 in Paper
2394-1214-1-0.49
Hoag et0.al. (1961)
O8
Hoag & Applequist (1965lines represent E(B VTYC
) = 0.48,
0.54, 0.60,10.15
and 0.65 0.51
mag, respectively.
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
B0.2V
Massey et al. (1995)
2394-1502-1
MasseyB1.5III
Massey et al. (1995)(1995).TYC
Negueruela
et al.11.68(2007) 0.19
presented-0.55
the morereli-et al. (1995)
arcsec) in width. The TYCstars 2394-1594-1
in the V and I11.86
band images were
835
PMS
members in NGC 1893 were selected using H↵ phoCu↵ey(1973a)able spectral
type and luminosity0.37
class for0.22
2 B-type
giant
little a↵ected by the column,
while their profile was
YSO catalogue (Caramazza et al. 2008), and
TYC 2394-1623-1
0.26signifi- We tometry,
Høg
et al.the(2000)
stars (HDE
243018 and 10.86
TYC 2394-1141-1).
adoptedcantly distorted in theTYC
U bandimage. Thus,
we replaced
the
the
X-ray
sourcelist (Caramazza
For thePMS
2394-1691-1
11.34
0.26
-0.64
Masseyetal.
(1995)
B0.5V et al. 2012).
Masseyet al.
(1995)
their classifications in this work (see Table 1). The lumiphotometric data of thestar
ID 2600 withthe meanmembers,
compared theB0.5V
detection efficiencies
TYC
2394-1738-1
11.19
0.25 values-0.67
Hoag et al.we(1961)
Masseyofetthe
al. in(1995)
nosity class of two B-type stars, TYC 2394-450-1 (B1.0III,
obtained in previous studies(Massey et al.
1995; Sharma
etdividualmembershipselection criteria within the completeTYC 2394-1744-1
11.78
0.23
-0.45
Hoag et al.
(1961)
Mbol =
4.988, log Te↵ 11.23= 4.395) and TYC
2394-1502-1
TYC
-0.62
Hoaglimit
et al.of(1961)
B0.5V
Masseyet al.
(1995)
al. 2007). In the caseof
the2394-1912-1
relatively faint star ID0.28
2595, theness
1 M . A total of
205 PMS members
with
mass

• photometric diagramから基本量を導出された。

Spectral type

(B1.5III, Mbol =4.460, log Te↵ = 4.336), is needs to be• IMFはSalpeter/Kroupa IMFとコン

ns were selected
centred on the2011), and thee vicinity of theregions is aboutt the reddeninghe regions. The
5 Myr isochrone
be larger in thes in the regions
as shown in thear range [E(Bential reddeningt in the Region
he Region 2 areening correction.
almost parallel↵erence between
ciated with theireddening.
a given age bin,he age distribuand 2 are shown
ograms are norion. To avoid an solid line) was
yr. The peak oft 1.5 Myr, whiled a younger age
e two groups is
im 129 and 130
the centre of theesent the cumun each region incated in Region
1. According toy that the PMS
is about 0.1 per
he age di↵erence

システント
19
=>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

Figure 14. The cumulative age distribution of the stars in Region
1 and 2. The solid line and dashed line correspond to Region 1
and 2, respectively. See the main text for details.

•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

17
Figure 10. The IMF of NGC 1893. To avoid the binning e↵ect weshifted the mass bin by 0.1 and rederived the IMF (open circle)
using the same procedure. The IMF of NGC 2264 (dashed line –Sung & Bessell 2010 ) is overplotted. Arrows are drawn for the
IMF below completeness limit. See the main text for details.

ideal laboratories to study the IMF of young open clusters
in the outer Galaxy. In this work, we can derive reliable IMF
over a wide mass range, from low-mass PMS stars to massive
O-type stars.
The mass of MS members is estimated by comparing
their position in the HRD to the evolutionary tracks of Ekström et al. (2012) with various initial masses. For PMS
members, the PMS evolutionary tracks of Siess et al. (2000)
were used to determine the individual mass of each star.
The sum of the masses of the members yields a total massof 1,300 M for the cluster. This is a lower limit because alarge number of sub-solar mass PMS stars below the completeness threshold are not taken into account. In order to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•  'PSTUFSJUFNJDSPO _, JSPO

ModelB••

9 .PEFM"iron-rich, ~1µm, ~40K

7 .PEFM#  'PSTUFSJUFNJDSPO _, JSPO

ModelC•iron-free, ~10µm,~100K

7 .PEFM$
Fig. 6. Top: calculated opacitiesof 1 µm forsterite grains at 50 K and GPSTUFSJUFNJDSPOJSPO
295 K. Bottom: calculated opacities
for 1 and 10 µm_,forsterite
grains(T = 150 K).
;POFTa

HSBJOHSPXUI• ベンチマークテスト
でパラメータ依存性を調べている。

USBOTJUJPOBMEJTL!

• 強い69µmかつ弱い

SBEJBMNJYJOH

23µmを作るには•

grain growth•transitional disk•

optical depth eﬀect

PQUJDBMEFQUIGM Aur(1.06 ± 0.02) ⇥ 10
(2.1 ± 0.07) ⇥ 10
Sz Cha(2.6 ± 0.1) ⇥ 10 12< 4.7 ⇥ 10 14
TW Hya(2.39 ± 0.04) ⇥ 10 11 (4.52 ± 0.07) ⇥ 10 12
CS Cha(5.23 ± 0.09) ⇥ 10 12(4.7 ± 1.1) ⇥ 10 13
CHXR22E(1.1 ± 0.1) ⇥ 10 14(4.4 ± 0.6) ⇥ 10 15
13
Sz18(3.8 ± 0.2) ⇥ 10
(2.5 ± 0.6) ⇥ 10 14
Sz27(1.83 ± 0.08) ⇥ 10 12(7.8 ± 0.6) ⇥ 10 14
12
Sz45(2.83 ± 0.02) ⇥ 10
(4.6 ± 0.1) ⇥ 10 13
Sz84(7.0 ± 0.1) ⇥ 10 13(5.59 ± 0.09) ⇥ 10 14
12RX J1615(2.24 ± 0.05) ⇥ 10
(2.5 ± 0.3) ⇥ 10 13
Oph22(1.9 ± 0.09) ⇥ 10 13(5.4 ± 0.1) ⇥ 10 14
13
Oph24(3.2 ± 0.2) ⇥ 10
(1.08 ± 0.03) ⇥ 10 13
a
SR 21
...
...
ISO Oph196 (3.54 ± 0.01) ⇥ 10 13(8.8 ± 0.6) ⇥ 10 14
DoAr 44(9.4 ± 0.1) ⇥ 10 12(2.66 ± 0.08) ⇥ 10 12
14
Fig. 3. Best(1.3fit of ±early
K-typetransitional
disks.<
Colors
Ser29
0.1)
⇥ 10
1.7 as⇥ in10Fig.151
13
Ser34(1.02 ± 0.05) ⇥ 10
(7.8 ± 0.7) ⇥ 10 15
RX J1842.9(3.41 ± 0.08) ⇥ 10 12(4.6 ± 0.6) ⇥ 10 13
12RX J1852.3(5.67 ± 0.09) ⇥ 10
(1.01 ± 0.07) ⇥ 10 12

On the gas content of transitional disks: a VLT/X-Shooter
study of accretion and winds
M
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(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(8.8 ± 0
K (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G

Fig. 4. Best fit of G-type transitional disks. Colors as in Fig. 1.

フィットする(Manara+2013b)
Fig. 3. Best fit of early K-type transitional disks. Colors as in Fig. 1• 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with a gaussian profile on the dereddend spectrum the LVC ofthe [OI] 630 nm line and we integrate the flux of the best fit
to derive the line flux. The error on the flux is derived from thestandard deviation of the continuum estimated around the line.
The derived flux, error, peak velocity (v0 ), and FWHM of thegaussian fit is reported in Table 5. The lines and their best fits
are shown in Fig. 7. In all the objects with detected [OI] 630
nm line, with the exception of LkH↵330, the line is slightly
blueshifted, with values of v0 ranging from ⇠ -2 km/s to ⇠ -8
km/s in most cases, and only two objects (Oph24 and ISO-Oph196) with v0 < -10 km
ject is still uncertain alength calibration desline is systematicallyin some kind of win[OI] 630 nm line der
km/s. We note, howeshould be considered
the nominal resolutio

Fig. 1. Best fit for transitional disks with spectral type in the M class. The red l
the photospheric template, light blue line the slab model, and blue line the best
8

10

Fig. 4. Best fit of G-type transitional disks. Colors as in Fig. 1.

• 630nm [OI]による検証

with a gaussian profile on the dereddend spectrum the LVC ofthe [OI] 630 nm line and we integrate the flux of the best fit
to derive the line flux. The error on the flux is derived from thestandard deviation of the continuum estimated around the line.
The derived flux, error, peak velocity (v0 ), and FWHM of thegaussian fit is reported in Table 5. The lines and their best fits
are shown in Fig. 7. In all the objects with detected [OI] 630• ウィンド起源の低速成分を17天体で検出

196) with v0 < -10 km/s. Even if the exact value of v0 in each object is still uncertain also after the procedure to correct the wavelength calibration described above, we see that the [OI] 630 nm
line is systematically blueshifted, meaning that it is originatedin some kind of wind. The mean value of the FWHM of the[OI] 630 nm line derived from the spectra of our targets is ⇠ 40
km/s. We note, however, that the values of FWHM . 30 km/s

Name
F
Accretion
properties[erg
s cm ]•

v0,[OI] 630 FWHM[OI] 630[km/s][km/s]
14
LkH↵330
(6.1
±
1.9)
⇥
10
7.7
24
Manara et al.: X-Shooter study of accretion and winds
in Transitional DisksDM Tau(9.1 ± 1.8) ⇥ 10 15
-4.3
38
14
LkCa 15(4.6 ± 1.9)
⇥
10
-8.7
43
Table 6. Properties of [NeII] 12.8µm
line from the literature
GM Aur(3.9 ± 1.0) ⇥ 10 14-2.2
42
Sz Cha(2.4 ± 0.9)Name⇥ 10 14-4.7
30
F[NeII] hires FWHM
v0
F[NeII] Spitzer Ref
TW Hya(9.0 ± 0.8)LkH↵330⇥ 10 14-4.8...
...24...
0.38 ± 0.19 G10
...37
...
0.55
G10
CS Cha(2.8 ± 1.5)DM⇥ 10Tau14-4.4...
LkCa......
...
0.28 ± 0.02 S12
CHXR22E< 8.0 ⇥ 10 16 15...< 0.5......
...
1.2 ± 0.06
G10
Sz18
< 2.6 ⇥ GM
10 15Aur... ...
Sz Cha14
...
...
...
1.62 ± 0.20 E13
Sz27(5.4 ± 0.3)TW⇥ 10Hya-7.6
49 -6.2
3.8 ± 0.3
14.6
5.9 ± 1.1
P09,G1014
Sz45(1.8 ± 0.5)CS
⇥ 10
-4.8
50 -3.3 3.63 ± 0.07 P09,E13
Cha
2.3
±
0.2
27
Manaraetal.:
X-Shooter
studyofaccretion
andwinds
in Transitional Disks
Sz84(2.4 ± 0.2)CHXR22E⇥ 10 15-6.5...
...44
...
...
...
14
...49
...
...
...
RX J1615(1.7 ± 0.8)Sz18⇥ 10-8.7...
......
...
0.63 ± 0.07 E13
Oph22< 7.3 ⇥ Sz27
10 15... ...
Sz45 14
...
...
...
...
...
Oph24
(3.0 ± 1.3)Sz84⇥ 10-13.0...
67
...
...
...
...
14
SR 21
< 6.4 ⇥ RX
10 J1615... ± 0.2
...
1.4
20.5
-7.5 2.76± 0.46S12
Fig.
8.
Logarithmof the mass accretion rate vs inner hole size
ISO Oph196 (1.4 ± 0.4)Oph22⇥ 10 15-20.4...
...68...
... our sample.
...
forDi↵erent
symbols are used to distinguish the
...50
...
...
...
DoAr 44(2.6 ± 1.2)Oph24⇥ 10 14-7.7...
methodsusedintheliterature
to derive the size of the inner
16
0.5
15.1
-8.3
< 3.0
S12,G10Ser29< 4.2 ⇥ SR
10 21... ± 0.1
...
hole.
...47
...
... Blue...squares are adopted when this has been derived using
Ser34(4.1 ± 0.8)ISO-Oph196⇥ 10 15-7.3...
Fig. 11. Logarithmic luminosity of the low-velocity component
DoAr 44
< 0.3
...
...
0.68
± 0.33 fitting,
S12,G10while red circles when the values are derived
IR-SED14± 0.8)Ser29⇥ 10-5.5...
43
...
...
...
...
of the [OI] 630 nm line vs the logarithmRX
of J1842.9
the accretion(5.0lumifrom
resolved
mm-interferometry observations. Downward arRXandJ1852.3(2.6 ± 1.3)Ser34⇥ 10 14-3.8...
...36
...
...
...
nosity of our objects (red filled symbols)for two samples
rows
are
upper
RX J1842.9
< 0.2
...
...
0.43 ± 0.13 S12,G10 limits. The two lowest points are, from leftof classical TTauri stars (blue empty circles from Rigliaco et al.
Notes.
Fluxesarereported
intheformat(flux±
err)
multipliedbytheto ±right,
and Ser29. The object at Rin =25 AU andRX
J1852.3
...
...
...
0.72
0.04 CHXR22E
G102013, green empty symbols from Natta et al. 2014). Downward
orderof
magnitude.
log
Ṁ
=
9.8
is
Ser34.
acc
arrows are upper limits.
References. P11: Pascucci et al. (2011); P09: Pascucci & Sterzik[OI] 630
1
2•

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•[OI]強度とRinの依存性無しでほぼ一定•

IR classによる違いも無し。[OI]生成プロセスはCTTSと同様と思われる。

• TD v.s. pTD(2009); S12: Sacco et al. (2012); E13: Espaillat et al. (2013); G10:
Güdel et al. (2010)

witha similar methodology. For these reasons we se2012). High resolution observations analyzedof TW Hya,
in particular,
12that results di↵ering from those found in previous
Notes.
are reported
units
of 10 14eterg; v0 suggest
and FWHM
lect
as a comparison
sampletheFluxes
objectsstudiedinby
Alcaláal.s cmThey
showed that most (&80%) of the [NeII]
emissionarises from thein units of km/s.
works
may
arise
from di↵erent sampleManara
selection
and/or di↵erent
These
are locatedet al.: X-Shooter study of accret
region where the disk is still optically(2014).
thick, but
still within⇠10 in the Lupus I and III clouds and havemethods
toestimate
Ṁ
.
To
avoid
this
possiblemethodologicalacc
AU from the central star (Pascucci et agesal., 2011).
⇠ 3 Myr, similar to the objects in our sample. The analysis
T
Among the objects in our sample,
been wasobservedof 13
thathavesamplecarried out with the same methodology as bias, we have shown here only the comparison between our sample
of
TDs
and
the
sampleofcTTs
inLupus,
whichisanalyzedawith MIR spectroscopy and all of them
havea
[NeII]
line
detecFig.
9. Logarithmof theWemass
accretion
vs logarithmic
logarithm ofrelathe Fig. 10. Logarithmic luminosity of the low-velocity component
the onewe used here.
show
in Fig.rate
9 thetion, as we report in Table 6. In all theseobjects
wedetected
also
in
the
same
way
as
our
objects.
We
stress
here
again
that
our
TDs
stellarmass forourandsampleof TDsa sampleclassition between
Ṁacc
M⇤ fortheseandtwoforsamples.
Ourofdata
are of the [OI] 630 nm line vs inner hole size for the TDs in our
the [OI] 630 nm line, with the onlycal
exception
have been
selected tobe mostly
strongaccretors, thus our sample
TTauribeingstarsSR21.
from Alcalá et al. (2014). Our targets are shown sample.
Downward
arrowsare upper
limits.

•

降着、ウィンドどちらも違い無し

reported as blue circles, while data from Alcalá et al. (2014) asasbluediamonds.
circles, whilegreen
The data
solidfromline theon literature
the plot isare
thereported
best fit withrela5.3. Accretion and wind properties green
inobjects
with
inner
disk
diamonds.
The
lines
are
the
best
fit
to
the
data
reported theintion from Alcalá et al. (2014), and the dashed lines represent
emission
Alcaláetal.
(2014)(solid
line)
and
the
0.4
dex
spreadreported
dispersion of 0.4 dex around this best fit relation. The typicalinthatstudy.
Downward
arrowsareasupper
limits.
The
twoerrorsthequantities
shown
a black
cross.
We
seelowthat
Following the analysis described in Sect.
3, on
we
dividethe sam- areestpoints
are,
from
left
to
right,CHXR22E
andSer29.
Typical⇠80%of theTDs
ple in two classes of objects: we refer
to objects
withno havenear- values of Ṁacc consistent with the valerrors
are
shown
withetthe
cross.
infrared color excess as TD, while thoseexcessreferredueswithfoundby areAlcaláal.black(2014)
for Lupus objects of the sameto as PTD, as reported in Table 3. TheMmorphologicaldi↵erence.
Therefore,
for
these
objects
we
do not see any di↵erence in⇤between these two classes is the presenceof warm properties
dust in the with respect to those in cTTs. We also
the accretion
inner region of the disk of PTD, which
be a )small
ring constantoflinecould(L[OI]630
appears
regardless the size of the dust deperform
on these
two samples a Kolmogorov-Smirnof
statisticaldust at few tenths of AU from the starpleted(e.g., Benisty
et al.,
cavity inthe2010;
disk with values between ⇠ 10 6 and ⇠ 10 4
test (K-S
test). diskWhen
restricting to objects in the same M⇤ range
Espaillat et al., 2010). This di↵erence
the implies
innermorL . inThis
that
the properties of the wind traced by thethe probability
that
the two samples are drawn from the same diphology could be due to di↵erent evolutionary
stagesofthese[OI] 630 nm line - that can be a disk wind, an accretion-driven

TDのinner holeはガスリッ
チであり、またinner hot
dustの有無にも影響していない

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 rate
pleof Rigliaco et al. (2013, and references therein) comprisesof cTTs.
mostly strong accretors with low to intermediate stellar mass,
The two main outliers in Fig. 8-9 are the object with an upwhile Natta et al. (2014) has a sample of low- and very lowper limit on Lacc , Ser29, and CHXR22E. The lower intensity ofmass YSOs with lower accretion rates. Here we compare only
accretion for these targets implies that the gas density in the inthe luminosity of the LVC of the [OI] 630 nm line derived inner disk is substantially depleted with respect to the one of cTTs.
our work and in the comparison samples. The best way to comThese objects do not have any peculiar property reported in thepare these values is to analyze the L[OI]630 -Lacc relation, which isliterature. From Fig. 9 we note that in the same range of M⇤ ofwell characterized in the literature (see e.g., Rigliaco et al., 2013,
these objects there are other TDs with values of Ṁacc comparaand references therein). This is shown in Fig. 11, where
we plot
ble 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– 18 –

Fig. 1.— (a) Multi-channel intensity maps of thfeatures 2 and 3, and (c) a velocity map of the CH3

IRAS 18151−1208 MM2. The origin of these map−12◦ 08′ 07′′ with the errors of ∼ 20′′ .