Star Forma)on Newsle0er#274 16-20本目
麻生 有佑

Dense gas in the Galac)c central molecular zone is warm and heated by turbulence
Adam Ginsburg et al. (2015)
Ø Central molecular zoneは極限的な物理状態での星形成を調べられる最も近い場所である。
Ø あらゆる物理を測れるガス温度がダスト温度と異なるという問題が永らくある。密度が足りないとdecoupleして合わない可能性があるが、銀河円盤では合って、中心では合わない。
Ø 加熱方法も光子、X線、乱流、宇宙線など議論が分かれている。

Ø –0.4°<l<1.6°, APEX, p-H2CO (J=3-2)@~214 GHzなど, Δv=0.8 km/s, binary dila)on(マスク)。
Ø 温度マップを作るために32,1-22,0/30,3-20,2の比をとった(比が高いと高温)。
A&A proofs: manuscr
Ø dendrogram clamp ﬁndでS/Nが良い領域を抜き出す。

Ø RADEXでモデルグリッド; n=102.5-7 cm–3で20個、dV/
dR=5 km/s/pc, N(H2CO)=1011-15.1 cm–2で30個、T=10-350 Kで50個で計算したものをスプライン補間して100x101x250。
Ø 観測側はH2CO 30,3-20,2の強度、速度、線幅、
32,1-22,0/30,3-20,2比、32,2-22,1/32,1-22,0比、CH3OH 42,2-31,2の強度で記述されるGaussiansと、Clump ﬁndの結果の2種類を考えた。
Fig. 6. The derived temperature4vs the measured
ratio
5 cm
–3 p-H CO S (3Ø Clumpごとに得られた(ベストフィットの)ライン比と温度の関係を得て、n~10
symbols represent leaves (compact と10clumps) andthe small symbols repre
The temperature includes constraints from the assumed constant H CO
color coded by signal-to-noise in the ratio R , with gray S/N < 5, blu
で温度マップを作った。
The black lines show the modeled temperature as a function of R for
2

⌫

2,1

2

1

1

abundance XH2 CO = 1.2 ⇥ 10 9 . While the measured values appear con
permit any density in the range 104 cm 3 < n < 105 cm 3 ; this plot mDense gas in the Galac)c central molecular zone is warm and heated by turbulence
Adam Ginsburg et al. (2015)の続き
Ø AV>10に興味があるのでUV加熱は無視し、星や超新星も局所的であり、X線は密度が3桁足りない(Ao+2013)。
Ø 宇宙線と乱流の加熱がダストと輝線の冷却と平衡とする。
Ø 宇宙線は深くまで浸透するので加熱は一様とするとζCR<10–14/s

A&A proofs: manuscri(a)Fig. 10. Plots of the p-H2 CO-measured temperature against the pred
5.1. While no trend is obvious in the data, most of the high signal-to-noline. There are few points with TG,measured > TG,turb , so turbulence gen
to heat the gas to the observed temperatures. The regions with predic
explanation: either their heating rate has been overestimated or their co
described in Figure 12; the blue symbols are compact ‘clump’ sources aindicate regions with lower limits on the measured temperature, i.e. any
TG > 150 K. The gray dashed line shows the TG,turb = TG,measured rela1level are outlined. (b) The same as (a) for the dendrogram-extract
fg > 3 are excluded. Regions with TG,turb > 180 K are indicated withthe ratio R1 , with gray S/N < 5, blue 5 < S/N < 25, green 25 < S/N <in blue and gray are ignored, the agreement between the predicted tem

which point the dust and gas strongly decouple (Laibe, Gonzalez &Maddison 2012). The second parameter involved is the dust-to-gas
ratio ϵ. For ϵ of order unity, dust affects the gas as much as gas affects
dust due to aerodynamic drag. Although the dust backreaction due todust inertia is negligible in most astrophysical systems, it becomes
critical in regions of discs where dust concentrates, such as particles
traps.
In this Letter, we present a mechanism where grain growth andfragmentation result in dust distributions that produce apparent gaps
when the dust backreaction is taken into consideration. We provide
the details on our numerical simulations and synthetic observables
in Section 2, discuss the origin of the apparent gap in Section 3, andconclude in Section 4.

potential of the star and planet (see Fouchet et al. 2007). The discis initially set up as in Fouchet, Gonzalez & Maddison (2010) witha constant surface density $(r) = 19.67 kg m−2 with radial extension from 4 to 120 au, a vertically isothermal temperature profile
T(r) = 15 (r/rp )−1 K, and is free to evolve. Grains have an initially
uniform size s0 = 10 µm and are allowed to grow and fragment.
In Gonzalez et al. (2015), we presented a suite of simulations thatvary the fragmentation threshold: Vfrag = 10, 15, 20, 25 m s−1 and+∞ (i.e. growth only), evolved for 100 000 yr or ∼400 planetary
orbits, and discussed our choice of parameters and their influenceon the results. In particular, our disc temperature is colder thanwhat observations suggest (e.g. Williams & Best 2014). Changing its value would affect gap formation (only moderately√ for ahigh-mass planet), as well as fragmentation (Vrel ∝ cs ∝ T ). In
the mid-plane, St is independent of T, leaving dust dynamics unchanged. In this Letter, we focus on the Vfrag = 15 m s−1 case, forwhich we see a second apparent gap.
Fig. 1 shows the gas and dust volume densities ρ g and ρ d , aswell as the grain size s in the disc mid-plane at the end of thesimulation. The planet at 40 au has opened a gap in the gas phaseand much of the disc interior to the gap has been accreted on toss low-density area.
relIn the dusts phase,ss thetheSPHstar, leaving a central
gap is more prominent, with sharper edges. Some material remains
in the inner disc and grains have accumulated at the dust trap located at the gap outer edge. These properties are commonly seen
in studiesJ of planet gaps in gas+dust discs (see e.g. Paardekooper& Mellema 2004; Fouchet et al. 2010). An 5
unexpected feature isfrag ring in the dust disc at r ∼ 90 au visible
the second high-density
in the density map (centre panel of Fig. 1) and in the radial profile of the mid-plane volume density (Fig. 2). Large grains (mmto cm in size) are abundant in the high-density regions (right-handpanel of Fig. 1 and bottom panel of Fig. 2), where relative velocities between grains are lowest and grain growth is more efficient.
The dust population is best characterized in Fig. 3, showing the
distribution of grains in the meridian plane and their radial size

ALMA images of discs: are all gaps carved by planets?
J.-F. Gonzalez et al. (2015)

We study the evolution of gas and dust in protoplanetary discs using
our 3D, two-phase, smoothed particle hydrodynamics (SPH) code,
described in Barrière-Fouchet et al. (2005). We self-consistently
treat the aerodynamical drag between gas and dust, SPH
includingthe backreaction of dust on gas. We follow the standard SPH
frag
implementation of artificial viscosity (Monaghan 1989) and use
αSPH = 0.1 and βSPH = 0.5, corresponding to a uniform Shakura &o parameter*∼ 10−2 (see Fouchet et al. 2007,
Sunyaev (1973)
αSS o–1for a discussion). Arena & Gonzalez (2013) showed that the expected properties of turbulence are naturally reproduced by the
SPH formalism. Grain growth is treated as described in Laibe et al.
(2008), via the prescription developed by Stepinski & Valageas(1997) based√ on the determination of the relative velocity of grains,
Vrel ∝ cs αSS St/(1 + St), as a function of local properties of thegas phase (e.g. its sound speed cs ) and grain size. We implement
grain fragmentation by defining a velocity threshold, Vfrag , to which
we compare Vrel . Collisions occurring at speeds below the thresh-

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Ø HL Tauで見つかったようなギャップは全て惑星が作ったのか？dust pile-upを代案として
シミュレート。
Ø ダストの成長はストークス数とg/d比で決まる。HL Tauでは数十AUでSt~1となり、そこではガスとdecoupleした粒子が最も速くmigrateする。g/d~1となると、空気力学
2 M E T H O D S A N D R E S U LT S
(aerodynamics)的にダストの慣性がガスに影響を与える。
2.1 Hydrodynamic simulations
Ø 3D, two-phase (gas & dust), SPH, α =0.1, β =0.5, α =0.01, V ∝c √(α St)/(1+St)
Ø VrelとV を比べて、小さいとgrain growth、大きいとsha0ering (サイズが小さくなる)。
Ø Mdisc=0.01 M , M =1M , g/d=0.01, 40 AUに5 M の惑星, 面密度は19.67 kg m–2で一定、温度は15(r/4 AU) K、ダスト粒子は初期に10 μm, V =15 m/s, 10 yrまで。
Ø ギャップの外縁60 AUにトラップされるのは普通。加えて、90 AUにも溜まっている。
Ø 60 AUのダストは成長しきっていて、90 AUのダストは成長中で内向きに移動している。

ALMA images
discs:
are all
gaps carved by planets?
L38 ofJ.-F.
Gonzalez
et al.
J.-F. Gonzalez et al. (2015)の続き

sumoto 2012, Bate et al. 2014) show thas performing simulations in which the sink particles are
tween gas accretion onto the protosta
significantly more massive than individual stars. However,Possible smoking-gun
evidenceforini)almass
segrega)onin
re-virialized
post-gasaccording to the high resolution hydrodynamic calculations
driven outflow
results in star formatio≈ 30% in the neighborhood of proto-st
Machida & Maatsumoto (2012) and Bate et al. (2014),
expulsion globularby
clusters
sub-pc scale. This would imply the overal
and direct observation (Lada & Lada 2003) only small fracHosein Haghi et al.tion(2015)of gas convert into stars.
clump (sub-pc to pc scale) would be les
ing its innermost regions. Such proto-sta
Based on the hydrodynamical simulations without feedare the only available ones that can track
back, Kruijssen et al. (2012) also claimed that the SFE isØ 銀画面のほとんどの星が埋もれた星団でできたかもしれない。
and outflow (driven by radiation and mag
higher in the regions where stars have formed, and henceØ 高密度領域ではIMFは観測的にも不変↔︎M
Fのスロープαと中心集中度c=log(r)に massiv
t/rcmoredividual proto-stars. Other
the long-term evolution of star clusters is only weakly afthat intend to form star clusters, either
相関 (De Marchirela)on)。
fectedby gas removal. They argued that this is because
or include only partial feedback. The SFof the low gas fractions within the pre-cluster cloud core,lations can never be reliable. The Dalecaused by the accretion of gas The
on to
the ofstars.
In this case 5eﬀect
gas expulsion
Ø ガスの動きは計算せず球対称重力ポテンシャルとビリアル平衡として、時間変化させ
lations that go up to globular cluster m
it however needs to be explained how the outer parts of theTable 1. The initial and the final parametersof theN-bodyruns after
t=100 Myr
evolution.
Column1 gives out theparticle representing a sub-cluster itself,
gas cloudareexpelled
without
heatingandpushing
る。ガス質量はτ
=0.6Myrからτ
で減少(gasexpulsion)。
the model name, in which the first twoDdigit numbers after ‘rh’ denote
initial half-mass radius in parsec,
GEthe
ticles themselves are of pc scale. Of courinner
gas ofcontentsof theeﬃciency
pre-cluster
core.
It also is
the last two digit numbers denote thepercentage
the star formation(sfe), andcloudthe degreeof mass5
Ø星団メンバーのN体はNBODY6を使う。銀河中心から8.5
kpcを220
segregation is written after the letter’S’. The initial
mass of allmodels is M = 10HIIM⊙ regions. Column 2ifgives
all overkm/sの円軌道。
the clump (the SFE can, in pra
diﬃcult
to explain
ultracompact
gasthecan siminitial 3D half-mass radius. Columns 3, 4 and 5 contain the adopted star formation eﬃciency, gas expulsion
the /router regions), gas expulsion should
ply beaccreted
ontostars
despitetheirluminosity.
In pcでrØ 潮汐外場は天の川銀河のもの。Half
mass半径は1,5,7,10
timescaleand initial degree of segregation,
respectively.
Column6 is thethe initial
fillingfactor, wherert0 =54.3 pc
h tが>0.1, <0.1を試す。
for all models. Note that the gas massisnotincludedinthecalculation
of
the
tidal
radius.
The
final
slope
of
the
impact on the cluster. This is the init
the short-lived
UCHIIm phase, the gas is heated significantly
5 Mmass
0.85 inside the tidal
radiusis presented
massfunctionafter 100 Myr ofM=10
evolution in the
0.50 ! M ⊙ ! Kroupa
Ø
N=170000,
canonical
IMF(α=1.3と2.3),
T=0-100
4, rangeothat weMyrplausibly adopt in our computa(upbeing
to 10
K following
or more)in theimmediatesurroundingin column 7, the initial canonical value
2.3. The
columns
give thefinal mass, the3D half-mass of theradiusthe simulated star clusters
aftersegrega)onは、S=1でfull100
Myr of evolution,
the finalconcentrationc = segrega)on
logto), as well of
We assume that the gas expulsion
t /rcoutflow
stars(within0.05 pc),
whichthen leadsthe
10 (ranØ ofprimordialmass(重いほど中心近くにある)as the two-body relaxation time of the star clusters after 100 Myr of evolution.
bedded phase at a certain time, τD , w
residual gas on larger scales.
f
0.6 Myr and corresponds to the UCHI
in α2
theMr<rt rsimulations
by
Model
rh0
ε
τGE However,S(rh /rt )0
c
Trh
hf
[pc]
[Myr]
[0.50 − 0.85]M⊙
[M⊙ ]
[pc]
[Gyr]
time, it is assumed that the mass of gDale, Ercolano, & Bonnell (2012) including the influ(1)
(2)
(3)
(4)
(5) photoionizing(6)(7)(10) stars(11)
nentially on the gas-expulsion characteri
enceofradiation (8)from(9)O-type
it isover which the natal gas is removed fromrh10-S0-sfe25
10
0.25 shown
1
0.0
36741
22.3
1.38 and13.14 agrees
that 0.18
the global2.16
SFE is between
20-30%
rh1-S0.9-sfe33
1
0.33
0.1
0.9
0.02
2.27
67392
4.4
2.40
1.81
time-varying total mass in gas is given b
is0.02
observed instar-forming
and embedded
rh1-S0.9-sfe25
10.25 with
0.1 what
0.9
2.21
65016 regions
5.2
2.10
1.50
rh3-S0.9-sfe25
30.25 clusters
0.1
0.9 which0.05 suggests
1.85that the
34930
12.9
1.55
4.40
star formation eﬃciency
Mgas (t) = Mgas (0) e−(t−τD )/τGE .
rh5-S0.9-sfe25
5
0.25
1
0.9
0.09
1.96
54274
18.6
1.61
9.36
can
at
most
be
ε
≈
30%
(see,
e.g.,
Lada
&
Lada
2003),
rh7-S0.9-sfe25
7
0.25
1
0.9
0.13
1.45
41716
21.6
1.34
11.02
rh10-S0.9-sfe33
10
0.33 although1
0.9 on0.18
22.6
1.195
the scale 1.08
of whole31486
molecular
clouds11.81
the star
The impact of the early gas expulsi
rh10-S0.9-sfe25
10
0.25
1
0.9
0.18
0.81
20060
21.5
1.178
10.08
formationeﬃciency
is2.28at most66038
a few(Murray
namical evolution of star clusters has b
rh10-S0.9-NoGE
10
1.0
–
0.9
0.18
16.2 percent
1.70
8.84
rh1-S0-NoGE
1
1.0
–
0.0
0.02
2.29
73827
1.9
2.65
10.10
2011). High-resolution mapping in star-forming regions has
Baumgardt & Kroupa (2007) using a lar
rh3-S0-NoGE
3
1.0
–
0.0
0.05
2.27
74379
4.3
2.44
12.40
uncoveredthat
outflows,
even
from
low-mass
proto-stars,
tegration. They showed that the survival
rh5-S0-NoGE
5
1.0
–
0.0
0.09
2.30
72668
7.0
2.42
8.01
rh10-S0-NoGE
10
1.0– very0.0 powerful
0.18
2.22
71039 clearing
12.7
1.84
are
in destroying
andcloud7.50
material
erties of star clusters are strongly influe
rh20-S0-NoGE
20
1.0
–
0.0
0.36
2.26
55029
22.8
1.67
6.30
away
from
the
regionsub-psof the expulsion of the residual-gas that
rh30-S0-NoGE
30
1.0
–
0.0
0.54 star-forming
2.34
10864 on35.0
0.86 scales
3.40 (e.g.,
Bally et al. 2011, 2012, McGroarty et al. 2004, Frank et al.
depending on the parameter values foPossible smoking-gun evidence for ini)al mass segrega)on in re-virialized post-gas
expulsion globular clusters
Hosein Haghi et al. (2015)の続き
Ø 初めにmass segrega)onがないとMFはほぼ時間変化しない。
Ø あると、)dally underﬁlling (rh/rt<0.1)だとIMRがα=1.3の範囲(m<0.5 Mo)のみ時間変化。
Ø mass seg.とtidally ﬁllingだと0.5<m0.85 Moでもベキが浅くなる。
Ø mas seg.とgas expulsionによって、cが大きいほどαが急という観測の傾向を再現。
Ø αの振動は外側の速度分散と関連。外の軽い星は中に、中の重い星は外に(revirialize)。
Ø 本当は銀河系も時間変化するので形成時期で潮汐場は異なる。
10

Haghi et al.

2.5

Slope α [ 0.3 MO• < m < 0.8 MO• ]

2
1.5
1
0.5
0
-0.5
-1
-1.5
-2

Figure 5. The time evolution of the MF-slope, α2 , for main sequence stars within the tidal radius in the mass range 0.50M⊙ !
M ! 0.85M⊙ . The initial number of stars in all models is
N=170000 (corresponding to M = 105 M⊙ ) with an input Kroupa
canonical IMF with the slope of α = 2.3. The initial half-mass

0.5

1

1.5
2
Concentration c=log10(rt/rc)

2.5

The Early ALMA View of the FU Ori Outburst System
A.S. Hales et al. (2015)
Ø 一部の若い星は通常の1000倍の質量降着率を示すことがあり、100年で0.01 Moを中心星へ降着させる(FU Orionis or FU Ori objects)。
Ø 考えられる理由は熱的不安定、重力+磁気回転不安定、disk fragmenta)on、惑星伴星。
Ø 1-10 AUとそれより外の円盤やエンベロープの関係を調べる必要がある。
Ø FU Oriは450 pcにあるバイナリー(FU OriとS)で、バイナリーのFU Oriはそれなりの数がある。
Ø FU Ori, ALMA Cycle 0, Δv~0.1 km/s, 12CO (3-2), HCN (4-3), HCO+ (4-3), θ=0.63”x0.51”
Ø 連続波は北の星にピークがあるが、singleよりdouble Gaussianでよくフィットできた。
Ø 12COは北と南を覆い、HCO+は南にピークがあり、HCNは北にピークがある。
Ø 12COで最も速度勾配が大きい東西にPV図を切るとケプラー回転。

⬅︎連続波(左上)、12CO(右上)HCO+(左下)、HCN(右下)The Early ALMA View of the FU Ori Outburst System
A.S. Hales et al. (2015)の続き
Ø 連続波は南北とも円盤を持つことを示し、半径は45 AU以下。円盤はバイナリー間隔
225 AUの0.2-0.5倍にtruncateされるという予想と合う(Artymowicz & Lubow 1996)。
Ø 直径で3000 AUを超えるエンベロープ成分は検出されなかった。
Ø 円盤のダスト質量は北で2x10–4 Mo, 南で8x10–5 Moだが、光学的に厚い可能性あり。
Ø 中心星質量は過去の研究から北が0.3 Mo, 南が1.2 Moなので、同じ温度を仮定すると重い星の方に軽い円盤があることになる。
Ø 高密度トレーサーのHCO+が南のみで検出されたので、北の円盤は南の円盤より低質量であると示唆される。
Ø 北のみで検出されたHCNは円盤の上空でできるので、北の円盤の方が高温と示唆される。àHCO+の検出と話は合う。

The Evolu)on of Inner Disk Gas in Transi)on Disks
K. Hoadley et al. (2015)
Ø 惑星ができるr<10 AUは光学的に厚いので調べづらいので、遷移円盤の穴に残ったガスで調べる。
Ø H2は他分子の104倍存在し、5000 Kまで生き残り、UVからもselfshieldされる。しかし双極子がないので、一番強いFUV (912-1700Å)を見る。星からのLyαで励起される。
Ø Hubble Space TelescopeのCosmic Origins SpectrographとSpace Telescope Imaging
Spectrograph (TW Hyaのみ)で撮られたFUV H2データを使う。普通の円盤と遷移円盤。
Ø 速度分解能はCOSでΔv=17 km/s, SITSでΔv=12 km/s。
The Evolution of Gas in Transition Disks

Ø モデルの仮定: ケプラー回転、円盤表面の浅く暖かい層、Lyα吸収はLTE。
Table 1
Ø 面密度はexponen)al tail付きの標準的なモデル。r/z=const.なので2D？
Stellar ParametersaTargetSpect.
Type

M⋆
(M⊙ )

d
(pc)

Av

id
(◦ )

Age
(Myr)

vsini
(km s−1 )ref.aAA TauBP TauCS ChaDF Tau A
DM Tau
GM AurHN Tau A
LkCa15
RECX 11RECX 15
SU AurTW HyaUX Tau A
V4046 Sgr

K7
K7
K6
M2
M1.5
K5.5
K5
K3
K4
M2
G1
K6
K2
K5

0.8
0.73
1.05
0.19
0.5
1.20
0.85
0.85
0.80
0.40
2.30
0.60
1.30
1.75

140
140
160
140
140
140
140
140
97
97
140
54
140
83

0.5
0.5
0.8
0.6
0.0
0.1
0.5
0.6
0.0
0.0
0.9
0.0
0.2
0.0

75
30
60
85
35
55
40
49
70
60
62
4
35
34

6.4 ± 0.2
5.9 ± 0.3
6.4 ± 0.1
6.3 ± 0.5
6.6 ± 0.2
6.9 ± 0.2
1.9 ± 0.9
6.4 ± 0.3
4.0 ± 1.5
6.0 ± 1.0
2.5 ± 0.9
10.0 ± 6.0
6.1 ± 0.3
6.9 ± 0.1

11.4
7.8
...
16.1
10.0
12.4
52.8
12.5
...
...
65.0
6.0
25.4
14.2(+13.7)

2,8,9,11,12,15,17,25
4,8,9,12,15,17,18,30
5,6,15,19,22
9,11,12,16,17,18
3,11,12,15,17,25
3,8,9,11,12,15,17,25
7,9,11,18,23
3,8,10,12,15,17,18
14,15,20,21
15,20,21,32
1,4,9,11,18
3,13,16,24,28,29,31
3,8,11,17,19
25,27,28,29(1) Akeson et al. (2002); (2) Andrews & Williams (2007); (3) Andrews et al. (2011);(4) Bouvier (1990); (5) Espaillat et al. (2007a); (6) Espaillat et al. (2011); (7) France et al.
(2012b); (8) Furlan et al. (2011); (9) Gullbring et al. (1998); (10) Hartmann et al. (1987);(11) Hartmann & Stauﬀer (1989); (12) Hartmann et al. (1998); (13) Herczeg & Hillenbrand(2008); (14) Ingleby et al. (2011); (15) Ingleby et al. (2013); (16) Johns-Krull & Valenti (2001);(17) Kenyon et al. (1994); (18) Kraus & Hillenbrand (2009); (19) Lawson et al. (1996); (20)
Lawson et al. (2001); (21) Lawson et al. (2004); (22) Luhman (2004); (23) McJunkin et al.
(2013); (24) Pontoppidan et al. (2008); (25) Quast et al. (2000); (26) Ricci et al. (2010);(27) Rodriguez et al. (2010); (28) Rosenfeld et al. (2012); (29) Rosenfeld et al. (2013); (30)
Simon et al. (2000); (31) Webb et al. (1999); (32) Woitke et al. (2013).

The Evolu)on of Inner Disk Gas in Transi)on Disks
K. Hoadley et al. (2015)の続き
Ø 最も大丈夫そうな[v’,J’]=[1,4]へのベストフィットから90%のフラックスが入るrinとroutを決める。