THE STAR FORMATION NEWSLETTER: No.239 10 Nov 2012
担当: 塚越(茨城大)2013/1/11
Interaction between brown-dwarf binaries and Sun-like stars
Murat Kaplan et al.
•観測によるBrown Dwarfの問題点•Sun like starとbinaryを組むbrown dwarfはたいていwide orbit(>100AU):brown dwarf desert•close binaryとなるのはbrown dwarf - brown dwarfのケース(<20AU)•Sun like starのwide orbitにあるbrown dwarfは、他brown dwarfとclose binaryになりやすい•Sun like starとbrown dwarf binaryのgravitational interactionを3体計算•Sun like starが重力によってclose brown dwarf binaryをwide orbitに補足するのは難しい•=>wide orbitを回るclose BDバイナリは、disk fragmentationによってその場所で形成された•他のbrown dwarfもfragmentationだろう
Exploring Magnetic Field Structure in Star-Forming Cores with Polarization of Thermal
Dust Emission
Akimasa Kataoka et al.
•磁場内の回転コアの進化と、偏波観測によってどのように見えるかをシミュレート•磁場強度、磁場方向、コア回転を変えて4 models•コアの進化を計算 -> 偏波を計算•偏波はコアのviewing angleに大きく依存する•protostar phaseになるとモデル間で大きな違いが出てくる•中心部(円盤方向)で低偏光度 -> viewing angleを制限•中心部(回転軸方向)でtoroidal ﬁeldの効果により低偏光度 -> 原始星の有無を制限•回転>磁場の場合S-shapeとなる。磁場>回転の場合はhourglassになる。
•形状からcloudのエネルギー状態を制限する•ALMAの偏波観測へ
6
Kataoka et al. 2012

re in Star-Forming Cores

コアの進化: 回転なしのケース

58回転ありModel:原始星フェーズ

Kataoka et al. 2012

viewing angle毎の偏光度の違い

Fig. 5.— Polarization maps for model 1 viewed from the angle◦ , 0◦ ) at four different epochs, which correspond to the(θ, modelφ) = (90Fig. 6.— Polarization maps for1 viewedfrom θ = 0◦ (pole-on view; a), 30◦ (b), 45◦ (c), 60◦ (d), 80◦ (e), anFig. 3.— Density distribution (colors and contours) and velocityFig. 4.— Density distribution (colors and contours)
respective
panels of Figure 3. The polarization degree (colors), po◦andvelocityf); φ = 0 when Mps = 0.5 M" .
vectors (arrows) on the y = 0 cutting plane at four different epochs
larization vector (lines), and column density (contours) are plotted
vectors (arrows) on the y = 0 cutting plane for models 2 (a), 3 (bfor model 1. Elapsed time t and protostellar mass Mps are listed
in each panel. The number density at the center of the cloud ncand c), and 4 (c). Elapsed time t and protostellar mass Mps arein each panel.
or protostellar mass Mps is shown in each panel.
listed in each panel. Panel (d) is a close-up view of panel (b).

Star formation, structure, and formation mechanism of cometary globules: NIR
observations of CG 1 and CG 2
disk tends to form in the direction perpendicular to theMinja
M. Makela
etFor
al.thispolarization vectors (the direction of the B-vector of themagnetic field lines
for model
4 (Figure 4[c]).
cates that the magnetic field lines are gradually inclinedmodel, because the initial magnetic energy (γ0 = 0.57)
toward the line-of-sight direction.
electromagnetic waves), respectively. Hereafter, we call
•ガム星雲にあるcometaryglobule(CG1,CG2)の形成メカニズムを調査
dominates the initialrotational energy (β0 = 0.01), theFrom the polarization vectors in the ξη plane, we canthistype of figure a polarization map. Viewed from◦◦cloud rotation is expected to have very little effect on thedirectly confirm that the magnetic field lines are inclined(θ, φ) = (90 , 0 ), the ξ and η directions in the obser•2つのcometaryglobuleの形成メカニズム
dynamical cloud evolution (fordetails, see Machida et al.
from their initial direction. In other words, the magnetic
vation grid correspond to the y and z directions in the2006). In addition, because of the strong magnetic field,
field lines point to the protostar around the center ofsimulation grid, respectively (see Figure 1). Thus, eachmagnetic braking is more effective in •supernova
model 4 than in
explosionshock
the cloud. Thus, we can imagine an hourglass structure
panel of Figure 5 corresponds to the same panel in Figmodels 2 and 3.
of the magnetic field lines on the polarization maps. Inure 3. (Note that the distributions on the y = 0 and
For all the models, over half of the initialcloud mass
•radiation(RDI):
radiation
O staraddition,the column density contours show a disk-likexdriven= 0 planesimplosion
are the same because
theUVsystem
is axisym- fromfalls onto either the protostar or the circumstellar disk
metric. In addition, because this model maintains theby the end of the calculation. This means that we calcuaxisymmetry, the physical quantities do not depend on
lated the cloud (or protostar) evolution until the Class the0 azimuthal viewing angle φ. Thus, we show only theor I protostar stage. We will show the evolution of thepolarization map with φ = 0◦ for this model.)magnetic field configuration in the next subsection.
Figure 5(a) is a polarization map of the initial state formodel 1, in which all the polarization vectors are aligned
4.2. Expected Distributions of Polarization
with the η-axis. Figure 5(b) shows the polarization just
After MHD simulations were performed, we calculated
before protostar formation. At this epoch, the polarizathe polarization of the thermal dust emission accordingtion vectors are distorted near the center of the cloud,Fig. 8.— Polarization maps for model 1 at dconfiguration, which indicates the existence
pseudo◦ , 0◦ ) when Mwith (θ, φ) of= a(90ps = 0.5 M"
disk, as described in §4.1.
the central region in the left panel enlarged f
To investigate the effects of viewing angle θ on thepolarization, we selected a snapshot when the protostar
has a mass of Mps = 0.5 M" . Figure
6 shows
protostarat ◦polarizaviewing◦ angles of θ = 3
tion maps with different viewing anglesof θ6).
= 0This(a),feature30(Figurewas explained◦◦◦◦(b), 45 (c), 60 (d), 80 (e), and 90(f); Onthetheazimuthal§4.2.1.
otherhand, in the rotati◦◦
angle is fixed at φ = 0 . The viewing
angle θ =0 corpolarization
degreeappears in two reg
responds to the pole-on view and andθ =lower-left
90◦ corresponds
regions of the center, as•観測と解析•NIR photometry: SOFI/NTT telescope, SIRIUS/IRSF•WISE, Spitzer, IRAS, AKARIのアーカイブデータ•3つのYSOを新検出: CG1に2つ、CG2に1つ(初YSO検出)•headとtailの質量比から形成メカニズムを議論 -> どちらもRDIの可能性がある

B. Nisini et al.: Mapping water in pro

Mapping water in protostellar outﬂows with Herschel: PACS and HIFI observations ofL1448-C
B. Nisini et al.
•shock tracerおよびchemical processのkey moleculeとして水分子の分布を調べる•観測と解析
•Hershel- HIFI & PACS: o-H2O (2̲{1,2}-1̲{0,1}), (1̲{1,0}-1̲{0,1})•JCMT: CO(3-2), SiO(8-7)•LVG radiative transfer calculation•shock spotにclumpyなH2O分布•平均温度300-500K、密度5E6 1/cm^3•shock heated spotをトレース。COのようなentrained gasとは異なる•H2Oのアバンダンスはoutﬂow内でほぼ一定(0.5-1.0 E-6)
•H2O/CO比に速度依存性 -> physical conditionの違いを反映•高速度域で低H2O/SiO比。shockモデルでは再現出来ず、protostellar windから出てきたdust-gree gas
Fig. 7. Contours of the H2 O 110 -101 emission integrated in three diffからのmolecular formation
sion in the same velocity bins (gray scale). The velocity ranges with rB. Nisini et al.: Mapping water in protostellar outflows with Herschelin the upper side of each panel. Green crosses mark the positions ofグレーがH2O

12Fig. 6. Overlay between H2 O 110 -101 (black), SiO(8-7) (blue)
and CO(3-2) (red) towards the L1448-C position.

4. Overlay of the JCMT CO(3-2) (left panel) and HIFI H2 O 557 GHz (right panel) emission on the H2 O 179 µm map. The
and red contours represent
integrated
in the velocity
ranges(−100,+4)(+6,+100)km s−1 , respectively.iceContours
The emission
thermalreactivityofHCN
andandNH
interstellar
analogues
3 inrawn from 25 to 200 K km s−1 , with steps of 20 K km s−1 , for the CO(3-2) map, and from
2 to 12 K km s−1 , with steps of 1.2−1""""m s , for the H2 O 557
GHzmap.
TheHPBWof14(JCMT)
and38(HIFI)
isindicated.
the two molecules trace the primary jet cannot be ruled out. InJ.A. Noble et al.

this respect, initial results presented in Panoglou et al. (2011) for•HCNとNH3(->NH4+CN-)の反応に関する実験
broad H2 O survival in disk-winds seem promising, predicte EHV gas in L1448, comparing it with the gas respon- is associated with the strong intermediate velocity
the molecularpeaking
around the ambient velocity. The association
for the wing emission, and foundsignificant
differences emission•fouriertransform
infrared
spectroscopy
ing
that
a significant fraction of water is synthesized in jets fromeen these two components. They found, in particular, that of SiO emission with the EHV collimated jet is a well-knownclasssourceshaving a mass accretion rate of 5 10−6 M" yr−1 ,(e.g. 0Hiranoet
EHV gas is relatively
rich in O-bearing ice内で低温度の熱反応のみで形成される
species and poor in feature characteristic of several class 0 sources•interstellar
that SiOinaring molecules compared to the wing regime. Thus the ob- al. 2006; Codella et al. 2007); it has been suggested
implyinga massflux rate of the order of that measured in L1448.
the jet is either directly synthesized in the dust-free jet acceld drop in the H2 O 1•反応レートは温度に依存
10 -101 /CO(3-2) ratio in the EHV regime
Finally,wenotethat the timescales to increase the waterore easily understood if the water-emitting EHV gas has a eration region (Glassgold et al. 1991; Panoglou et al. 2011) orabundanceto values X(H2 O)>10−5 in a gas with T > 400 K arematerial where silicon
is releasedr temperature and/or•NH4+CN-はCN-イオンに対応する吸収が4.78
density that the gas responsible for the originates in shocked ambientµmあたりに出てくる(e.g.order
Gusdorfof 100 yr (Bergin et al. 1998), which match well theemission. Kristensen et al. (2011) compared the excitation into the gas phase by the disruption of grain coresof the•他、反応パラメータなど
itions for water in the EHV gas with those responsible for et al. 2008).
dynamicaltimescale for the L1448 jet propagation of the orderwing emission towards the central position, but were unable
An origin in the primary jet has been recently ∼supported
150 yrby(Hirano et al. 2010).
entify significant differences in the two regimes. Santangelo both interferometric observations in the HH212ofobject (Cabrit et(2012), on the other hand, found that at the L1448-R4 po- al. 2007) and by the molecular survey conducted onWith
regards to the broad H2 O emission at intermediate vethe L1448n the high velocity H2 O emission is associated with gas at R2 EHV bullet by Tafalla et al. (2010), who showed that the bullocities,Kristensen
et al. (2011) suggested an origin in shocks
nsity about an order of magnitude lower than that of the gas lets possess a peculiar chemistry with respect to the standard outcaused
bytheinteraction
between the outflow and the envelope.
onsible for the low velocity emission.
flow wing emission, suggesting an origin different from shocks.
Suchshocks
would
beexpected
to produce significant SiO emisThe fact that SiO(8-7) is more prominent than H2 O 557 GHz inthe EHV bullets may either be an excitation effectthe resultsion,or sincethe disruption of grain cores occurs at shock speeds
SiO and H2 Oof an enhanced SiO/H2 O abundance ratio, or both. The excita->∼Tdsfterw
dthhea
CCtm

4F
liw
ov
amr

1we(aaEarly-stage young stellar objects in the Small Magellanic Cloud
J.M. Oliveira et al.
•小マゼラン雲の34 YSOに対して、opt-IR-radioを用いた同定作業•4天体でPAH emission、14天体でice absorption•30/33天体がmassive star formationにある段階•残り3つはdusty diskとenvelopeを持つintermediate mass YSO•1つはD-type共生星•SMCでは初となる

A method to measure CO and N2 depletion proﬁles inside prestellar cores
Laurent Pagani et al.
•prestellar coreでのCOやN2のdepletion調査は重要だが観測が難しい•COとN2でdepletionする密度が一桁違うのはなぜか•N2はdipole momentなし、COはambientの影響など•DCO+を使用してCOのdepleted regionにおけるabundanceプロファイルを調査する•N2H+, N2D+, H2D+のデータも使って重水素化度の調査 -> COのabundanceに焼きなおす•DCO+(1-0,2-1,3-2)観測 / IRAM 30m•L183をモデルケースとして観測•COのradial方向のアバンダンス変化に比べてN2の変化は緩い•コア外側のCOアバンダンスは星間空間と近い値だが、N2は低い•N2はまだproduction phase
Constraints on the Radial Variation of Grain Growth in the AS 209 Circumstellar Disk
Laura M. Perez et al.
•Tタウリ型星AS209の円盤。ダスト進化のradial方向の調査•SMA 0.88mm, CARMA 2.8mm, VLA 8.0mm, VLA 9.8 mm•0.2-0.5”分解能 -> 26-65AU相当•2layer model, exponential tailの面密度を使用してフィッティング•β(R)を導入して、波長で統一させる•β(R)一定とすると、各波長でT(R)やΣ(R)が異なってしまう•半径が増えるとβも大きくなる。円盤通しての変化量はおよそ1•半径方向のダスト成長を示唆•内側でcm以上のダストサイズ、外側ではミクロンサイズ•Brinstiel+2012のダストモデルと比較•radial-drift dominated population
6Pérez et al.
Grain Growth in AS510 209 Diskn(a) ∝ a−q; q = 3.5n(a) ∝ a−q; q = 3.0

Dust Composition:

Dust Composition:

62% Water Ice
30% Organics
8% Silicates

62% Water Ice
30% Organics
8% Silicates02.5

1σ1.5m1 cm−10.1 1010

β C = 1.0m1m

Dust Composition:
mm

mm
0.7

amax

CARMA

VLA

SMA

200.1

30.

m
m

62% Water Ice
30% Organics
8% Silicates−2406080
Disk Radius [AU]100

101200.01

30

Fragmentation−dominatedDrift−dominated
3σ constraints on amax
−qn(a) ∝ a ; q = 3.0n(a) ∝ a−q; q = 3.50.
5[cm]

10.50

Fragmentation−dominatedDrift−dominated
3σ constraints on amax−qn(a) ∝ a ; q = 3.5

ν11κ [cm2 gm−1]

β0.88−8mm

β I SM0.
12m−2

Σ [g cm ]

3σ 2σ

10100
ν [GHz]300

20406080100120the 2040 disk.60 Black80 line:100
120
Fig. 4.— Left: Dust opacity spectral slope, β, vs. radius, inferred from multi-wavelength
observationsofAS 209
Disk Radius [AU]Radius [AU]best-fit β(R), colored areas: confidence interval constrained by our observations. Verticaldashed-lines indicate the spatial Disk
resolutionof ourFig. 5.— Surface
density (top)onandβ(R)maximum
grain size
as a functionof radiuscalibration(bottom), for two
grain size distributions: q = 3.5 (left),observations, errorbar in top-left corner indicates additional systematic
uncertainty
arisingfromamplitudeuncertainty.

The IRAM-30m line survey of the Horsehead PDR: II. First detection of the l-C3H+
hydrocarbon cation
J. Pety et al.
•馬頭星雲のPDRにおけるミリ波ラインサーベイ•IRAM 30m (WHISPER project)•PDR領域でl-C3H+のラインを検出(星間空間で初検出)•帯域内にある6/8ラインを検出•C3H+ cationと分子の反応は、small hydrocarbonの形成に重要•分布はHCO(ラジカル)や、small carbon chain、C2H(PDRでenhanceする)と似ている•実験室での測定、干渉計観測に繋げるべき
Early evolution of the birth cluster of the solar system
Susanne Pfalzner et al.
•どのようなクラスターから太陽系が産まれたのか•starburst cluster•leaky cluster (OB association)•N体計算を使ってSBCとLCの環境におけるencounter probabilityを調べた•SBCだと中心密度が高すぎて太陽系を作れる円盤が生き残れない -> LCのほうが考えられる•太陽系の性質を制限するencounterは最初の2Myrでしか生じず、その後のencounterは無視出来るレベル•太陽系が高いcircularityを保っている説明にもなる