THE STAR FORMATIONNEWSLETTER
No. 263
塚越崇(茨城大)28 Nov. 2014

10

Antoniucci et al.: Accretion evolution in YSOs

POISSON project III. Investigating the evolution of the mass accretion rate
S.Antoniucci et al.

• POISSON (Protostellar Optical-Infrared Survey on NTT)•

Antoniucci et al.: Accret

この論文ではLupus, Serpens領域の結果を紹介• 大きなサーベイで星進化と降着現象の関係を調べる• BrγからLaccを導出=>Maccに変換
Fig. 6. Mass accretion rates of the whole POISSON sample plotted as a function of the stellar mass (left) and age of the targets• 他領域のデータも使用する(right). Upper
limits on Ṁacc are indicated by downward triangles. The typical error bar of a single point is also shown. The dashedlines in the right graph show the expected Ṁacc evolution for a fiducial model of a viscous disk (Hartmann et al. 1998) (see text for
more details)
a 0.5M⊙ star with initial disk masses equal to 0.1M⊙ (lower line) and 0.2M⊙ (upper line). Sample symbols are* aroundacc
the same as in Fig. 4.

• L とL に良い相関
Fig. 4. Accretion luminosity as a function of the stellar luminos• Macc Mß, Macc t-ηを与えると進化傾向が良く見える
ity for the whole POISSON sample. The various symbols referto the five diﬀerent cloud samples (see legend). The solid anddashed lines show the locus of Lacc =L∗ and Lacc =0.1L∗ .

Planet-vortex interaction: How a vortex can shepherd a planetary embryo
S.Ataiee et al.

• 円盤渦による惑星(種)への影響をhydroシミュレーションで調査•

円盤の面密度変化=>渦が形成される(dead zone or planet)• モデル計算•FARGO、ケプラー回転円盤、locally isothermal•

ガウシアン型のpressure bumpを埋め込み、渦の形成を与える•

渦の形成後(?)、惑星を渦周囲や渦内に埋め込む

S. Ataiee et al.: Planet-vortex interaction: How a• 惑星は円盤のある位置にトラップされるが、渦に対する位置関係が
異なる

1.02
2 (t=390)

RWIの1:1 resonanceの位置

1.00

ap

•

4 (t=470)

1.01

0.99

0.98

1 (t=350)5(t=510)3 (t=440)10

3t (orbit)

10

4

Fi
gr
sit

2 (t=390)

4 (t=470)

1.01• パラメータ依存性を確認(右図)
ap

1.00

0.99
1.02

•

1 (t=350)5(t=510)3 (t=440)

4 (t=470)
disk2 (t=390)
viscosity,
disk surface density,1010t (orbit)planet mass, planet position

0.98
1.01

4

3

ap

1.00
Fig.
4. Planet trapping in the frame co-rotating with the vortex.
5(t=510)The left upper3panel
demonstrates
a planet initially set inside the(t=440)1(t=350)0.99
1.02(position 0), comes out and moves to the position 1 andvortexthen return2 (t=390)
back to2. Afterward, it oscillates around position 5
4 (t=470)
0.981.01
until
it is locked to10thevortex at this location. We
4illustrate the310t(orbit)planet’s behaviour in the right upper panel in a more exaggerated1.00
way to show the radial migration of the planet as well. The
Fig.
4. Planetin semi-major
theframe co-rotating
theby
vortex.
lower
panel trappingshows theaxis of thewith
planettime.
5(t=510)3 (t=440)
The
left
upper
panel
demonstratesaplanetinitially
set
insideThe0.99
numbers
denote the times of the marked positions in thetheleft
1 (t=350)
vortex(position0),marked
comes theoutpositiveand movestothe position
1 fromandupperpanel.
Weandnegative
torques
渦から遠い所でトラップ
then
back
Afterward,
it oscillates
aroundfrom
position5
thereturnvortexby theto+2.and− signs. The
indirect torquethe star
0.98
4
3
10
10
until
is opposite
locked tosign
theofvortexat tthis
location.
We illustrate the(orbit)has itthe
the vortex
torque(Sec.4).
SGが入るとタイムスケール↑
planet’s behaviour in the right upper panel in a more exaggerway
show
the radialmigrationof the planet
as well.
The
4. to
Planettrapping
inembryothe frame co-rotatingwiththe vortex.
on: How a vortex ated
canFig.
shepherd
a planetary
lower
panel
shows
thedemonstrates
semi-major axisof theplanettime.
Theleft upperpanel
a planet
initially
set byinsidethe1.06
O0 1left(position
0),
comes
out
and
moves
to
the
position
andThevortex
numbers
denote
the
times
of
the
markedpositions
in
the
O
1.04
thenpanel.
returnback
to 2. Afterward,
oscillates
aroundposition
O90
uppermarkedthe positiveit andnegative
torques
from5
OWe
1.02
it is
to thevortexTheat thislocation.
WefromillustratetheI 0 startheuntilvortexbylocked
the + and− signs.
indirecttorquethebehaviour
in
the
right
upper
panel
in
a
more
exaggerI
hasplanet’s
theopposite
sign
of
the
vortextorque(Sec.4).
180
1.00
ated way to show the radial migration of the planet as well. The
I
0.98
lower
panel Ishows the semi-major
axis of the planet by time.
0.96numbers denote the times of the marked positions in the left
The
1.06
upper
panel.
We marked the positive and negative torques
O0 from
0.94
O
1.04
300
400
500
600
700
800
900
1000
1100
the vortex
by the + and − signs.t (Thetorque from
the
star
O90
or bi t indirect)O
1.02
has the opposite sign of the vortex torque (Sec.4).
I0
I 180 disFig.
1.00 6. Semi-major axis evolution of four selected models
played in Fig. 5.

Fig. 7. The trapping of the planet in discs with diﬀerent background viscosity. The upper panels demonstrate the locking position of the planet toward the center of the vortex. The lower
panel shows the change of planet semi-major against time.

• 渦が大きい場合には惑星トラップ
apは妨げられる傾向•

Fig. 7. The trapping of the planet in discs with diﬀerent background viscosity. The upper panels demonstrate the locking position of the planet toward the center of the vortex. The lower
panel shows the change of planet semi-major against time.
Fig. 8. The evolution of planet semi-major axis for diﬀerent
background surface densities. The higher amplitude and longer
oscillation of the planet semi-major axis before trapping shows
the stronger interaction between the planet and the vortex in thecase with Σ = 10−3 .

ap

ap

•

Fig. 7. The trapping of the planet in discs with diﬀerent background viscosity. The upper panels demonstrate the locking position of the planet toward the center of the vortex. The lower
panel shows the change of planet semi-major against time.

0.981.06

I

I

O0
0.961.04
O
Fig. 15. Semi-majoraxisevolution for the standard models
but
O90
O
with double
heightfor thebump
models).
yellow
−4
0.94
1.02
face
density= 10, the(Height2
planetmigrates
moreTheslowlybut is300
400Σ BG 500
600
700
800
900
1000 I 01100
t
(
or
bi
t
)
and blue lines
represent
the
model
with
and
without
self-gravity,
trapped
I 180
1.00 eventually.

Fig. 8. The evolution of planet semi-major axis for diﬀerent
background surface densities. The higher amplitude and longer
oscillation of the planet semi-major axis before trapping shows
the stronger interaction between the planet and the vortex in thecase with Σ = 10−3 .
Fig. 8. The evolution of planet semi-major axis for diﬀerent
background surface densities. The higher amplitude and longer
oscillation of the planet semi-major axis before trapping shows
the stronger interaction between the planet and the vortex in thecase with Σ = 10−3 .

The fate of scattered planets
B.C.Bromley et al.

• 巨大惑星によって弾き出された惑星とガス円盤の相互作用を計算•

遠い惑星が低離心率になる時の円盤状態を調べる• power-lawの低質量円盤モデル••

時間依存した散逸のどちらかを入れる•Σ*exp(-t/τ)の散逸•

時間とともにinner edgeが広がる(transition disk)

巨大ガス惑星は5AUに置く• 計算結果(次ページ)•

重い惑星ほど円軌道になりやすい•

expで散逸する円盤のほうが小さい半径で円軌道に•

inner holeが広がる場合、軌道の多様性が進む。スーパーアースでも100AUを超えて円軌道になる• HR8799の惑星(super-Jupiter at 70AU)を説明できるかもexp(-t/τ)散逸のケース

inner holeが広がるケース

面密度

カラーは惑星の初期位置: 100(赤), 200(緑), 300(青) AU
Fig. 4.— Simulation outcomes depicted as elliptical orbital paths for a suite of models. The x-y(細)1,3, plane;
10, 30(太)areMearthcoordinates線の太さは惑星の質量:
give planetary positions in theirorbitalmodelsoﬀset from one another sothat they are sorted in rows and columns according to model parameters. The first and second

Kinematics in Partially Ionized Molecular Clouds: Implications for the Transition to
Coherence
N.D.Balley et al.

• 分子雲における電離度が、形成されたコアの速度構造や分子輝線プ
ロファイルに与える影響を調査•

Balley&Basu(2014)の発展 (Paper1)• 電離度を調べるには中性物質とイオンの振る舞いを調べる事が重要•

電離度分布を変えた際の中性分子とイオン分子の振る舞いを計算
4• モデル計算Bailey, Basu & CaselliSL•

partially ionized, isothermal,
magnetic interstellar molecular cloud•

電離度の分布として、step-like(SL)とcosmic lay(CR)を採用

CR•

non-ideal MHD IDL code (Basu &Ciolek 2004)を使用•

パラメータはPaper1のもの+2つFigure 1. Column density enhancement maps of clump/core regions. Each panel is a zoom-in of the full 64πL0 x 64πL0 region that
focuses on the region containing the most evolved clump/core in each model. Contours show the visual extinction value in 2 magnitude steps
starting at AV = 2. Panels are organized to depict increasing frequency of perturbations (initial perturbation only (left), perturbations
every 10t0 (middle), and perturbations every 5t0 (right)) with top row depicting models with the step-like ionisation profile and the bottom
row depicting models with the cosmic-ray only ionisation profile. Panels show each model at the last time of the respective simulations.
Top row (from left to right): Model I (t/t0 = 143.6), Model II (t/t0 = 80.7), Model III (t/t0 = 70.0). Bottom row (from left to right):

end pointの面密度分布

6Bailey, Basu & Caselli• velocity ﬁeld(右上図)•

SLのほうが速度分散大きい• 合成スペクトル(右下図)•

コアと低密度領域(LDG)の中性
分子とイオン分子

91Model IIILDG Neutral
LDG Ions

Core Ions with the most structure apthe AV 0.8= 2 mag contour
pearing within the AV = 4 mag contour where the aver3−3
age volume
0.6 density is above 4.4 × 10 cm , thus indicating that the majority of our region of interest would
indeed be detected by these two tracers. For the lowdensity 0.4
gas, we assume the majority of the neutral particles are carbon monoxide (CO) and the majority of ions13+
are H CO0.2 . No chemistry is included in the models, so
constant abundances of the various species are assumed,
as well as optically thin conditions. The eﬀects of the
inclusion 0of a simplified chemistry and radiative transfer
will be discussed in a future paper.

locity (Myers et al. 1991). For each pixel, we assume thatthey are small enough such that they are thermalised,i.e., there is no non-thermal component, so that observed
non-thermal motions will reflect the kinematics. Basedon this assumption, the width of the Gaussian is!
"
F W HM = ∆VT2 = 8 ln 2 c2s .
(9)The height of the Gaussian is given by the FWHM scaledby the neutral or ionic column density. The total spectralline is then constructed by summing up all individualcomponents along the line of sight.
For our analysis, we look at spectra
for VI
lines of sight
Model IVModel
1both on and oﬀ source. For the on-source spectra, we cre4.1. Linewidth and Centroid Velocity Analysis
ate spectra for the two perpendicular lines of sight (NS4.1.1. Method
0.8
and EW) discussed in the previous section. For eachWith the above points in mind, we construct synthetic
line of sight, we create four separate spectra: a low denspectra by assuming Gaussian line shapes for each pixel
sity component and core component for both the neutral
0.6the line of sight velocity in each pixel (i.e., V
such thatxand ionic gas, respectively, assuming that the high andor Vy depending on orientation) is the mean velocity forlow density components are traced by the four molecules
the pixel.
0.4 The width of the line is dependent on thediscussed above. For this analysis, the low density gas
thermal and non-thermal velocities, i.e.,(LDG) is defined as regions along the los with visual ex!
tinction 2 < AV < 7 mag and the cores are defined
0.2
2
2
(8)
F W HM = ∆VN T + ∆VT ,as regions along the los with visual extinction AV > 7
mag. The sources for which we create these spectra are
where ∆V0T2 = 8 ln 2 c2s is the thermal velocity component
listed in Table 2. For each model, the MC designation2and ∆VN2 T = ∆Vobs− ∆VT2 is the non thermal velocity
refers to the “Main Core” which is the most evolved core
-4 corresponding-2
0 to the2observed
4 ve- 6 (i.e.,-2 caused
0 the 2simulation
4 to terminate;-6 ∆Vobs-6 the-4one that
6component with

Scaled Column Density

電離度分布の違いによってコア/
クランプの速度構造に影響を与える

Scaled Column Density

Model
I density enhancement contours. Panels and contours depict the same
Figure 3. Velocity maps at final time for each model withcolumnCore Neutralmodels and levels as Figure 1, respectively.

v/cs

v/cs

• velocity ﬁeldの時間変化
速度の「谷」でコア/クランプが生まれている13

Figure 10. Velocity maps for Model II for six times (in order from top left to bottom right: t/t0 = 75.1, 78.1, 79.1, 79.5, 80.1 and 80.6).
The contours show the visual extinction in 1 magnitude steps starting at AV = 2 mag at the same time as the velocity map.

A near-IR spectroscopic survey of massive jets towards EGOs
A.Caratti o Garatti et al.

• 大質量星のprimary jet( molecular outﬂow)のサーベイ観測•

降着現象と密接に関わるので重要• SpitzerのGLIPSEサーベイから天体を選定•

4.5mum超過があるものがベース (extremely greened object; EGO)

• 最終的に18天体を選出•

中/大質量星に付随する、明らかなH2 ﬂow/jetでかつ明るいもの• 観測とデータ••

ESO New Technology Telescope(NTT)のカメラSoﬁ•

H2とKsのイメージング•

イメージングと低分散分光(R~600)の同時取得

アーカイブデータからSEDを作成

A. Caratti o Garatti et al.: NIR spectroscopy of massive jets from intermediate- and high-mass YSOs

• ノットと天体位置からジェットのposition angleを導出• 多くのサンプルで大きな
precession angleを示す• multiple jetの例もある
Fig. 1. H2 and continuum-subtracted H2 images (left and right panels) of the IRAS 13481-6124 jet. This jet is one of the few examples of jets withsmall precession angles and it has asymmetric lobes. Source and knots positions are indicated in the figures. Blue and red lobes are also reported
according to Kraus et al. (2010).
Fig. 1. H2 and continuum-subtracted H2 images (left and right panels) of the IRAS 13481-6124 jet. This jet is one of the few examples of jets withsmall precession angles and it has asymmetric lobes. Source and knots positions are indicated in the figures. Blue and red lobes are also reported
according to Kraus et al. (2010).

Fig. 2. Same as in Fig. 1 but for the IRAS14212-6131 flow. This object represents a good example of precessing flow with only one detected lobe.
The thin dashed lines show the measured precession angle. The positions of sources, knots, masers, EGOs, and H ii regions are indicated in the• 分子輝線(H2)と原子輝線が見える
ケースと、分子輝線しか見えないケー
スがある• ro-vibrational iagramから温度と密度を算出

A&A proofs: manuscript no. Massive_Jets_paper_astroph

gion within a 3 contour level above the sky background. Identified knots along with their coordinates, 2.12 µm fluxes, uncertainties, and projected areas in the sky, are reported in the
Appendix (Columns 1–6 of Table A.1). Third, the AV (H2 ) andT (H ) values derived from our spectroscopy (Sect. 4.3.1 and TaFig. 5. 0.95–2.5 µm2 SofI spectrum of knots 4 and 5 in the G35.2N outflows. An asterisk near the line identification marks the de
2 and 3 sigma.
are indicatedby the symbol“ ”.
bleTelluricA.1)linesare
then used
to deredden
the 2.12 µm line flux andto derive the intensities of all H2 lines
by applying a radiative
Fig. 5. 0.95–2.5 µm SofI spectrum of knots 4 and 5 in the G35.2N
outflows. An asterisk near the line identification marks the de
2 and 3 sigma.
Telluriclinesgas
are indicatedby thesymbol“ ”. in Caratti o Garatti et al. 2006,
code
with
in LTE(seedetails
2008).

For those knots associated with a particular flow but not encompassed by any slit, we adopt an average value of the physicalparameters derived from the flow knots. Then, for each considered knot, our LTE code provides: a) the line intensities involving levels with 0  v  14 and 0  J  29 (Ev,J  50 000 K);b) the total H2 intensity; and c) the H2 luminosity by assuming
the distance to the source provided in Table 1. This process is repeated for all the knots associated with each flow, to estimate LH2of the flow. LH2 values are reported in Column 7 of Table 3. The
inferred values range from 0.5 to 45 L . This roughly reflects thelarge Lbol range in our sample. It is worth noting, however, that
theµminferred
LHof2Knotvaluesbe lower
values,Fig. 6. 1.5–2.5
SofI spectrum
C in themight
IRAS16122-5047
flow. limits
Labels areof
as intheFig. real
5.
because our estimate does not take into account part of the emission
from
the
cold
Hin2 thecomponent,
purerotational
agrams(Boltzmann
plots)
to derive
density, visual traced
exumnbydensitythe5. molecules(Nv,J ) in the upp
Fig.
6. 1.5–2.5µm SofI
spectrumofKnotcolumn
C
IRAS16122-5047
flow. Labels
aretheas inofFig.

• 高温(2500-3000K)、高密度(1018-20 cm-2)領域がH2の起源•

70%はno ionic emission -> 天体近くのoutﬂow cavityにおけるC-type shock•

30%はatomic emissionあり -> magnetically collimated jetのJ-type shock• ジェットのmorphology•

多くのprecessionやmultiple jet -> クラスターによる相互作用• 広がったH2放射+ノット -> 定常降着+バースト成分

A. Caratti o Garatti et al.: NIR spectroscopy of massive jets from intermediate- and high-mass YSOs

ever, in two out of seven cases (flows from IRAS 142126131 and GLIMPSE G035.0393-00.4735; see right panels of
Fig. 2 and B.14) a marginal detection (between 2 and 3 )of H2 emission from the second lobe is observed. This suggests that the higher visual extinction towards one of the lobes
can hamper our observations in the NIR regime. Neverthe低質量星でも見られる関係に乗る
less, the presence of monopolar outflows from high-mass (Zapata et al. 2006; Fernández-López et al. 2013) as well as lowCOで得られてるものと同じベキ数
mass YSOs (Codella et al. 2014) has been observed in a few
sources through submm outflow tracers, such as SiO and CO.
The authors propose that such a monopolar geometry is due toanisotropic ambient cloud conditions or to an intrinsic asymmetry in the flow. In addition, the presence of close companion(s)might divert or even inhibit the ejection (see also Reipurth 2000;
Murphy et al. 2008, and discussion therein). Unfortunately, we
do not have any (sub)mm observations of theses flows to verify any of these scenarios. However, the Spitzer/GLIMPSE images show emission in excess at 4.5 µm located at the positionof the undetected lobe in four out of seven flows (namely IRAS
13484-6100, IRAS 14212-6131, IRAS 15450-5431, GLIMPSE
G035.0393-00.4735). Although the nature of such excess is controversial (e.g. De Buizer & Vacca 2010; Takami et al. 2012), it
has been often associated with outflows (Cyganowski et al. 2008;
Tappe et al. 2012).

• 中心星光度とH2強度に相関あり••• 中心星質量とH2質量にも弱い相関=>重い星ほど質量を多く放出(ただし分散大きい)