Star	  Forma)on	  Newsle0er	  No.256	  6本目–10本目	
麻生 有佑	  (東京大学	  D1)	

Sizes	  of	  protoplanetary	  discs	  a2er	  star-­‐disc	  encounters	  Andreas	  Breslau,	  Manuel	  Steinhausen,	  Kirsten	  Vincke,	  and	  Susanne	  Pfalzner	  (イントロ)	  •of protoplanetary discs after star-disc encountersat least
at rinit
riastronar being
pproacherial areor a renamicalcer par(3)he innerhe mass

oplanefinition
only ofseveral
methodsity andwer law
ation ofges, theity falls
l 1998).
c’s surp in thes methntrast in

r an ene partiee casedo notdial diserefore,cle disresenta-bservace denn be obpplyinga and e.
s yields
Whenwith the(3) onegnostics(p = 1)

rongestt of theatisticalve to belied.
rber onhe cor-(a)An	  ALMA	  Con9nuum	  Survey	  of	  Circumstellar	  Disks	  in	  the	  Upper	  Scorpius	  OB	  Associa9on	  John	  M.	  Carpenter,	  Luca	  Ricci	  and	  Andrea	  Isella	  (イントロ)	  • 円盤散逸のタイムスケールはクラスターでdisk	  frac)onを比べて測られてきた。	  
• IRは~1	  AUを見るのに対してsub	  mmは外まで見える。	  	  (内容)	  • ALMA,	  880	  µm	  • Upper	  Scorpius,	  K-­‐M	  13/20個検出,	  Class	  II,	  5–11	  Myr,	  	  M*=0.1–1.6	  Mo,	  d=145	  pc	  • Taurus	  (1–2	  Myr)と比較するためにフラックスからMdustを計算。	  
• log	  Mdust/ME=0.68+1.01log	  M*/Moを得る、Taurusでも傾きは1.1。	  
• Upper	  Sco.にはM0.5–M3に対応するM*の天体がないà分布のブレイク?	  
• IRのdiskはM*が小さいほどたくさんあるàM*を揃えてTaurusと比較。	  
• M*がM0.5より軽い組ではMdustはTaurusより25%小さい。	  
• M*がM3より重い組ではTaurusとの違いは微妙。	  ∆δ [arcsec]

J155624.8-222555

J155706.4-220606

J155729.9-225843

J155829.8-231007

contours = -1, 2σ-2, 2σ-2, 2σ-2, 3, 10, 20σ

J160357.6-203105

J160421.7-213028

J160525.5-203539

J160532.1-193315-2, 3, 6, 9σ-2, 3, 30, 130σ

-2, 3, 6σ-2, 2σJ160600.6-195711

Sco PMS 31

J160622.8-201124

J160643.8-190805-2, 2σ-2, 3, 6σ-2, 2, 4σ-2, 2, 3σ

J160702.1-201938

J160823.2-193001

J160827.5-194904

J160900.0-190836-2, 2σ-2, 3, 30, 60σ-2, 2, 4σ-2, 3, 7, 11σ

2 J160900.7-190852 J160959.4-180009 J161411.0-230536 J161420.2-190648
10-1-2, 3, 30, 60σ-2, 2, 3σ-2, 3, 10, 20σ-2, 3, 30, 120σ-2
2 1 0 -1 -2∆α [arcsec]

20天体中13天体で連続波を検出。	

Fig. 3.— Dust mass as a function of stellar mass for the 20 stars in the Upper Sco sample. Sources with 880 µm continuum detections
are denoted by circles, and 3σ upper limits to the 880 µm detections are indicated by triangles. The dashed lines shows constant ratios ofdust mass to stellar mass of 0.01ζ and 0.001ζ, where ζ = 0.01 is the dust to gas ratio. The shaded region shows the correlation between

Upper	  Sco.のMdustとM*の関係。	

ig. 1.— Contour maps of the ALMA 880 µm continuum emission for 20 K-M type stars in Upper Sco. Each map is centered on thedust where
mass solid
andanddisk
mass derived by Andrews et al. (2013) in Taurus, where the width of the region shows the 0.7 dex spread about thelar position after correction for proper motion. The contour levels are indicated in the lower right for each panel,
dottedbest fit relation.
tours indicate positive and negative flux densities, respectively.

8Fig. 4.— Left: Cumulative distribution of the stellar masses for the Taurus (grey) and Upper Sco (black) samples for stellar masses

10

nd magnetic field perpendicular to eachthat the tangential component of the magnetic field
discussion of oblique shocks. Here weshock planeis reversed
inmbipolar	  the post-shock
region. Thes
Forma9on	  of	  M
agne9zed	  P
restellar	  Cores	  w
ith	  ADiﬀusion	  ck equations and write them in a moresolutions are commonly specified as intermediate shocks
tailedjumpconditions.
and	  Turbulence	  Wu 1987; Karimabadi 1995; Inoue & Inutsuka 2007). Aplane-parallelshock
these two field-reversal solutions, we notice that rf,expChe-­‐Yu	  Chen	  and	  withEve	  Cuniform.	  Ostriker	  preρ0 and ionization-recombination equiapproaches the hydrodynamic jump condition (rf,hydro =he shock front is in the x–y plane, thewhen θ → 0, and rB,exp (θ), 2 is smaller in magnitude than(イントロ)	  the z-direction (v0 = v0 ẑ), and the upsolutions when θ < θcrit . Thus, we classify this set of sols in• themass/ﬂux比で重力収縮が決まる。ISMでは磁場が強く、星では弱いが間の変化はx–z plane, at an angle θ to therf,exp (θ), 2 and rB,exp (θ), 2 as the quasi-hydrodynamic s
未だ謎。	  
+ B0 cosθ ẑ) such that Bx = B0 sin θ isThis quasi-hydrodynamic solution can create gas compr• 準静的なambipolar	  iﬀusionでは10t_ﬀかかり(観測2–5t_ﬀ)、GMC内はﬂowがある。	  
nt perpendicular
to the flow. dTheparammuch stronger than the regularly applied fast shock con	  m valueof the Mach number and plasma
and may be the reason that when θ < θcrit , even ideal
CO12
therefore become(内容)	  simulations can generate shocked layers with relatively• 両側のﬂowに対して傾いた磁場を想定し、ショックを解く。	  
mass-to-flux ratio (see Sections 4 and 5 for more details)
B02
1
1 10
n~103	  cmof–3ではM>10	  oが必要で観.
(1), • m/fの臨界値はM/ΦB=1/2π√Gだが、B~10µm,	  ≡=Thedefinition
θcrit can beMderived
from Equatio22
β0測されるコアの1Moで収縮するにはambipolar	  8πρ0 cs
βx sin θ2diﬀusionが必要。	  
which turns
negative when 1 − (2 cos θ/β0 M2 ) >1 − (2rf (θ) cos2 θ/β0 M2 ) < 0:
• ideal	  HD,	  等方収縮では、あるθより大きい(ﬂowと磁場が垂直側)とpost	  shockで磁MHD
shocksMare
describedby compres場が大きくなり、ISMà星の磁場の弱まりに合わない。	  
nd magnetic
field:
β0 M222cos θ > cos θcrit =.
Bstream⊥,downstream
2r(θ)
f crit• , シミュレーションではθと電離度χがパラメーター。	  
r B⊥ ≡.
(2)
B⊥, upstream
ream• 磁場が強いとコアはゆっくりでき、磁場に垂直なフィラメントができる。	  
Using Equation (5) and considering only the terms ∼M• 磁場が弱いとambiがあって中心部が良く圧縮されコアがすぐできる。	  
becomes0) and(A14) in CO12, we have√2
• 中間の磁場だと磁場に沿ったサブフィラメントネットワークができる。	  
cos θcrit
β0 M
"2 !
,
≈2sin
θ2
θ
sin
2 cos2 θcrit
2=M+
1+or√
β0 M2β0β0 M
2"!
"2sin θcrit +sin θcrit − 1 = 0.

• 密度が107	  cm–3を超えたときにコアの物理量を測ると、密度、サイズ、質量、磁場、CFE,	  CMF,	  質量とサイズの間の冪、は大体観測と合っている。	  
• コアの質量とサイズはθとχに依存しない。

The AstrophysicalChenJournal,
785:69 (20pp), 2014 April 10& Ostriker

The Astrophysical Journal, 785:69 (20pp), 2014 April 10

Astrophysical Journal, 785:69 (20pp), 2014 April 10C

Chen & Ostriker

parameter that controls ambipolar diffusion effects in oursimulations, following CO12. We use typical values of themean neutral and ion molecular weight µn and µi of 2.3mH
and 30mH , respectively, which give the collision coefficient(see Equation (20c)) between neutrals and ions α = 3.7 ×
1013 cm3 s−1 g−1 .
The physical parameters defining each model are ρ0 , v0 =|v0 |, B0 = |B0 |, θ , and χi0 . We set the upstream neutral number
density to be n0 = ρ0 /µn = 1000 cm−3 in all simulations,consistent with typical mean molecular densities within GMCs6(e.g., Larson 1981; Williams et al. 2000; Bot et al. 2007; Bolattoet al. 2008). We choose the upstream B0 = 10 µG as typicalof GMC values (Goodman et al. 1989; Crutcher et al. 1993;
Heiles & Crutcher 2005; Heiles & Troland 2005) for all ourFigure 2. Schematic configuration for our simulations.
simulations. To keep the total number of simulations practical,(A color version of this figure is available in the online journal.)
we set the large-scale inflow Mach number to M = 10 for
all models. Exploration of the dependence on Mach numberof ambipolar diffusion and of core formation has been studied√in previous simulations (CO12 and GO11, respectively). For
three-dimensional velocity dispersion δv = 3 · δv1D at theour parameter survey, we choose θ = 5, 20, and 45 degrees
e of the simulation box would beto represent small (θ < θcrit ), intermediate (θ > θcrit ), and"!
large (θ ≫ θcrit ) angles between the inflow velocity and cloud√ πGαvir 1/4 1/2 1/2 1/4
magneticChen
field.
For each θ , we conduct simulations with χi0 = 3,v(L(20pp),
3 10
M cs ρ0 Lbox 1/2 . (22)box ) =85:69
2014 April& Ostriker
1510, and ideal MHD to cover situations with strong, weak, and no
ambipolar diffusion. We also run corresponding hydrodynamicemphasize the influence of the cloud magnetization instead ofsimulations with the same ρ0 and v0 for comparison.
perturbation field, our simulations are conducted with 10%
A full list of models is contained in Table 1. Table 1 also−1
he value δv(Lbox ), or δv = 0.14 km s with αvir = 2. With
lists the steady-state post-shock properties, as described iner δv(Lbox ), simulations can still form cores, but because
Section 2.1. Solutions for all three types of shocks are listed-self-gravitating clumps can easily be destroyed by strong
for the θ = 5◦ (A5) case. For the θ = 20◦ and θ = 45◦
ocity perturbations and no core can form before the turbulent
cases, there is only one shock solution. Also included inrgy dissipates, it takes much longer, with corresponding
Table 1 are the nominal values of critical mass and radius forher computational expense.
Figure 4. Similar to Figure 3, butfor model A5X3 withsymmetric
upstream magnetic field
nearly parallelto
to the(θ = 5 ) and low ionization underfraction (χFigure= 3).5. Similar to Figure 3, but for model A45ID, with 45 angle between upstream v and B, and ideal MHD.
spherically
volumes
beinflow
self-gravitatingthese(A color version of this figure is available in the online journal.)(A color version of this figure is available in the online journal.)
steady-state post-shock conditions, as discussed in Section 2.2
3.2. Model Parameters(small θ and/or small χi0 ) (seeand strong(large θ and/or(11),large (12),filaments
perpendicular
the magnetic
field.“thermal”
The timescales
Equations(14),and to(15)).
Bothhydrodynamic solution. This is a direct evidence omeasured
the average density
within the z = 0.5 pc ± ∆x layer
χi0 ) magnetic effects in the shocked gas. The thickness of theat which compressed layers become
gravitationallyunstableand“magnetic”
criticalmasses
arelisted.
In
mostmodels,of a transient stage of ambipolar diffusion (CO12att=0.3
Myrformodel
A5X3
and found that this overdense
A schematic showing our model setup is shown inFigure2.
post-shock layer is very different for these two extreme cases.
and start to form cores also differ. Note that in thecases
with5−3
Table 2 lists the physical properties of the poslayern ≈ 1.4
× 10 cmlayer, which exceeds the steady-MHD
Especially at early time (0.3MMyr), structureis alsodifferentand ambipolardiffusion(Figures3,and4),hasa highly
compressed>
MM≫M
indicating
thatthesimulation box is 1 pc on each side and represents
aremag,sphth,sphmag,sph⊙
measured at t = 0.2 Myr, as well as the corresp
shockregion.
jump condition
predicted
in these two cases, with stronger magnetic effects producing
forms in the center of the post-shock
Quantitatively,
we in Table 1 even for the quasipost-shock regions are dominated by magnetic support, andn within a GMC where a large-scale supersonic converging
109
either ambipolar diffusion
or anisotropic condensation would
w with velocity v0 and −v0 (i.e., in the center-of-momentum
be needed to form low-mass supercritical cores, as discussedme) collides. The z-direction is the large-scale inflow direcin Section 2.2. On the other hand, the quasi-hydrodynamic, and we adopt periodic boundary conditions in the x- andshock solution for models with θ < θcrit (i.e., A5 cases) has
rections. We initialize the background magnetic field in theMmag,sph < Mth,sph < M⊙ downstream. If this shock solution
ud, B0 , in the x–z plane, with an angle θ with respect tocould be sustained, then in principle low-mass supercriticalconvergent flow. For simplicity, we treat the gas as isothercores could form by spherical condensation of post-shock gas.
at temperature T = 10 K, such that the sound speed is
In order to collect sufficient statistical information on the= 0.2 km s−1 . The neutral density within the cloud, ρ0 , isv

Bv

xyz

◦

i0◦

Figure 8. Statistical distribution of core mass (left) and size (right) from all models combined.
(A color version of this figure is available in the online journal.)

質量とサイズはパラメー
ターに依らない。	

Figure 9. Median (diamond) and ±25% values (vertical bars) of core mass, size, average magnetic field, and normalized mass-to-flux ratio for different model

The Astrophysical785:69(20pp),
2014
Aprilhigher
10
Chenangle& Ostriker
parameters. InJournal,
each figure
fromleft toright,
X corresponds to increasing ionization, and larger A corresponds to larger
θ , or increasing pre-shock(upstream) magnetic field Bx = B0 sin θ parallel to the shock front. The dashed line in the core size plot (bottom left) indicates the lower limit (0.015 pc) of resolved
core size; for our simulations, 0.015 pc ≈ 3∆x.

磁場に沿った方向に流れてくる非等方収縮	

cores found in the simulations did not form isotropically. Wediscuss this further in Section 6.
To further investigate the relationship between core masses
and sizes, we binned the data set by log Lcore and calculate
the average core mass and mean density for different model

We note in particular that for the θ = 20◦ and θ = 45◦ ideal
MHD cases, the masses in Figures 8 and 9 are more than anorder of magnitude lower than the limits for a spherical region
at post-shock conditions to be magnetically supercritical, aslisted in Tables 1 and 2. This implies that the low-mass bound14

Spitzer	  View	  of	  Massive	  Star	  Forma9on	  in	  the	  Tidally	  Stripped	  Magellanic	  Bridge	  C.-­‐H.	  Rosie	  Chen,	  Remy	  Indebetouw,	  Erik	  Muller,	  Akiko	  Kawamura,	  Karl	  D.	  Gordon,	  Marta	  Sewi	  lo,	  Barbara	  A.	  Whitney,	  Yasuo	  Fukui,	  Suzanne	  C.	  Madden,	  Marilyn	  R.	  Meade,	  Margaret	  Meixner,	  Joana	  M.	  Oliveira,	  Thomas	  P.	  Robitaille1,	  Jonathan	  P.	  Seale,	  Bernie	  Shiao,	  and	  Jacco	  Th.	  van	  Loon	  (イントロ)	  • 大質量星はUVや星風で環境を変えるので、銀河の環境がMassive	  Star	  Forma)onへ与える影響を調べるにはYSOが良い。	  
• LMCとSMCが200	  Myr前に衝突してできたBridgeがあり、潮汐系であり、また金属量も小さいという特徴を持つ	  (Bridge:	  1/5–1/8	  Zo,	  LMC:	  1/3	  Zo)。	  
• HI,	  H2の柱密度は大きく、CO	  1–0で分子雲も8個同定されている。	  
(内容)	  • Spitzer	  (3.6–24µm)にSuper	  COSMOS	  Sky	  Surveyと2MASSの可視とNIRを使用。	  
• [8]	  vs	  [4.5]–[8.0]のColor	  Magnitude	  Diagramを使って[4.5]–[8.0]>2,	  [8.0]<14–([4.5]–[8.0])という条件でYSOを決定。	  
• さらにSEDがYSOかlate-­‐type銀河かでYSOを絞る。	  
• 26個のYSOをSEDで分類↓	  Type	  I:	  エンベロープからの放射あり。Ksより短波長で見えない。	  
Type	  II:	  円盤あり。可視で暗く、赤外で明るい。	  
Type	  III:	  ほぼ星。	  • 中心星、ﬂared	  disk,	  ﬂa0ened	  回転エンベロープ、bipolar	  cavityを含むモデルでSEDフィット。	  
• パラメーターはM*,	  Ltot,	  Mdot,env,	  Mdisk,	  age,	  Av	  • Mdot,env/M*とMdisk/M*から進化段階を分類↓	  • Stage	  I:	  >10–6	  yr–1,	  Stage	  II:	  <10–6	  yr–1かつ>10–6,	  Stage	  III	  <10–6	  yr–1かつ<10–6	  • 26個のYSOのうちsingleとわかった17個についてTypeとStageを比べると	  • Type	  I,	  I/II	  -­‐-­‐-­‐	  Stage	  <1.6,	  Type	  II,	  II/III	  -­‐-­‐-­‐	  Stage~2.0	  • Type,	  Stage	  IIIは見つからなかった	  (LMCにはいる、金属量が少ないとダスト散逸も速い)。	  
• M*は最大で10	  MoでLMCの17–45	  Moほど重くない。	  
• Bridgeと同じM*の範囲で可視で見える個数を比べると、Bridgeで81%,	  LMCで56%。	  
• foregroundの減光はともにAv~3.4à分子雲のダストは同じくらいある。	  
• 星周ダストの減光のMdot,env,	  Mdustに対する分布もBridgeとLMCで同じ。	  
àBridgeの方がclumpyだから散乱光で短波長が良く見える	  (進化も速い)。	  
••••SFEを調べるためにあるN(HI)の範囲で見つかるYSOの割合をピクセルの個数比で調べた。	  
N(HI)が大きい比εも大きい。εはN(HI)が大きい側でLMCの1/3だが、小さい側では2倍。	  
N(HI)が大きいと低金属量なので星形成の効率が下がるのだろう。	  
N(HI)が小さいと潮汐力によって星形成が誘発されるのだろう。	  Young Massive Stars in the Magellanic Bridge

19

Spiral	  arms	  in	  the	  disk	  of	  HD	  142527	  from	  CO	  emission	  lines	  with	  ALMA	  ValenSn	  ChrisSaens,	  Simon	  Casassus,	  SebasSan	  Perez,	  Gerrit	  van	  der	  Plas	  and	  Francois	  Menard	  (イントロ)	  • transi)on	  diskは惑星形成と関わり、非対称、warp,	  スパイラルはcompanionの証拠。	  
• HD	  142527はNIRでギャップやスパイラルが見つかっている。	  	  (内容)	  • ALMA	  12CO	  2-­‐1	  (1”x0.7”),	  3-­‐2	  (0.6”x0.4”)と13	  CO	  2-­‐1	  • 南側は分子雲の吸収があるので北側に注目。	  
スパイラルを3つ検出	  (S1はNIRでも検出されている)	  4•Christiaens et al.
最も外のS2では輝度温度、励起温度は11–15	  KでCOがdeple)onしていないので、
ing•S1 andS2 in the CO 3-2 Imap (Figure 3c). Uncertainties in the measured positions of the spirals weredustが少ないか赤道面にある。	  
set to max{binsize, quarter of the beam size}. We thenperform least-squares fits to each spiral independently• two温度からスケールハイトとdisk	  aspect	  ra)oが計算できる。	  
using
di↵erent parametric formulae, as follows.
The polar coordinates (✓, r) of an acoustic wave cre• byradial	  	  ateda planet, c
atut上の極大点でスパイラルを決めて、2種類のモデルでフィット	  location (✓ , r ), on a gaseous disk
can be approximated by (Muto et al. 2012, and references
therein):(Muto+2012,	  Kim	  2011)。	  	  ✓(r) = ✓ + sgn(r r )peakccc⇥(✓✓chcrrc

11+◆1+ 1

11+1↵+

1◆)1↵+✓rrc◆ ↵(1)

where hc is the disk aspect ratio at radius rc . It is as-