銀河形成シミュレーションにおけるガス・星形成・フィードバックの物理矢島秀伸 	
  東北大学 学際科学フロンティア研究所	
第４回銀河進化研究会＠大阪大学 2017/6/7-­‐6/9	

アウトライン	• イントロ	  • 銀河形成における重要なガスの物理	  	
  １：星形成	  	
  ２：フィードバック	
  	
  ３：星間ガス	
  • まとめと今後	イントロ	
  （より大きいダイナミックレンジへ）	
これから	
これまで	
大規模構造	  銀河	   星間ガス	
  ・宇宙再電離	  ・クラスタリング	
  ・環境効果	

・星形成率	
  ・輻射特性	  ・銀河風	

>10 Mpc	~kpc	

・多相構造	
  
f	
・重元素	
  
・電離構造	

星形成	
BH	
  

1-100 pc	

<1pc	

SKA	
 すばるHSC	
 ALMA	  JWST,	
  TMT	

空間スケール	

銀河形成におけるkey	
  processes	
電離光子→	
  宇宙再電離	

アウトフロー	
フィードバック	  （輻射圧、輻射加熱、SNe）	
角運動量輸送	MBH	

POPIII星?	
ダスト生成	

星形成率	

共進化?	

検出可能性	  
	  (JWST:2018~)	

+環境効果（Major	
  Merger,	
  UV	
  background,	
  Odal	  ﬁeld）	

銀河形成の重要なバリオン物理	

銀河	
星形成	

フィードバック	

星間ガス	

銀河形成の重要なバリオン物理	

銀河	
星形成	

＊星形成効率	  ＊初期質量関数	
  
＊星種族	
  
＊金属量、乱流	LEdd 4⇡GmpMBH
ṀEdd ⇠ 2 =c⌘c T⌘
星形成効率	
Ṁ = min ṀBondi, ṀEdd✓◆
MBH
星形成効率	
t = 0.45 Gyr ⇥ ⌘ ⇥ ln
0
Mstar
M
BH
✏⌘
Mgas ⇢0
⇢DM =
2
r(1
+
r/R
)
シミュレーションにおける星形成効率(のようなもの)	
s⇢gas⇢˙ star ⇠ ✏
tdyn
⇢
˙
f
=
A
SN
star
高赤方偏移銀河内はどのような値が適切だろうか？	

異なる環境での星形成効率 	

The Astrophysical Journal, 768:74 (22pp), 2013 May 1

tdepl (Gyr)PHIBSS z=2-3
PHIBSS z=1-1.5
z=0 SFGs
EGS13011166 z=1.53 resolved

0.1
0.3

Tacconi	  et	
  al.	  (2013)	
  	

2

4 z=1-1.5 (N=50)
PHIBSS
mass-matched
ms COLDGASS

LIRGs

0.1

ie
s
cl
ou
10

1

2

3

10

typical uncertainty

5

4
-2

log (Σmol gas (Msunpc ) )

la
ct
ic

-6

-3

la
x
ga

-4
ds

High-­‐z	
  v.s.	
  low-­‐z	

-2

-2

ga

-1

0

0%

1.5 Gyrlog(SFR) [M yr-1]

N=1
0.7 Gyr

SFR	

0

0.2

normal spirals

1% %1fraction-1

ULIRGs2-2

log ( Σstar form (Msunyr kpc ) )α=4.36

Lada+(2012)	

The
Astrophysical1
0.2
0.5 Journal,
1 745:1902(6pp), 2012 February
5

Bz
Ks

3

Tacconi et al.

0-1.0-8 -0.5004
0.5
6log (tdepl (Gyr))log(M) [M ]
2

8

10

12

Mgas	

Figure 2. SFR–molecular-mass diagram for local molecular clouds and galaxies

from the Gao& Solomon(“Kennicutt–Schmidt”(2004a) sample. The solid
symbols forcorrespond
Figure 7. Molecular-gas–star-formation-scaling relation. Left: molecular-gas–star-formation surface-density
relationshiprelation),z ∼ 0 tomeasurements of dense cloud masses either from extinction observations of“near main-sequence” SFGs (open gray circles: Kennicutt 1998a; Gracia-Carpio 2008; Genzel et al. 2010; Armus et al. 2009), z = 1–1.5 SFGs (black filled circles:
the galactic clouds or HCN observations of the galaxies. The open symbols
PHIBSS, including Daddi et al. 2010a; Magnelli et al. 2012b), z = 1.53 EGS13011166 (filled cyansquares:toGenzel
et al. 2013),= 2–2.5
detected
SFGsclouds(openandcorrespondmeasurementsof totalzcloudmassesof the samecrossed red squares, PHIBSS). The dotted gray and red lines mark the best-fit linear (N = 1) fits to thegalaxies,
low-z andhigh-zRight:
distributionforof depletion
eitherfrom data.
extinction
measurements
the galactictimescales
clouds or CO
nferred from the 50 z = 1–1.5 SFGs in the left panel (black shaded histogram). The distribution hasobservations
a median of= 0.7 Gyrandgalaxies,
a dispersionof 0.24dex.theForthefortdepl
the galaxies.
For thepentagons
represent
locationsnormalwhileofthe
of starburstasgalaxies
are(Saintonge
represented by
SDSS selected, z ∼ 0 COLDGASS survey for main-sequence galaxies matched to the same mass ofrange
andspirals,
coverage
thepositions
main sequence
PHIBSSsquareshistogram).
(LIRGs) andAinverted
Triangles
high-zet al. 2011a), the corresponding depletion timescale and scatter are 1.24 ± 0.06 Gyr and ±0.23 dex (bluetypical triangles
Poisson(ULIRGs).
error for both
low-represent
and high-z
BzK galaxies. The star formation rates for the Gao and Solomon galaxies have
distributions is indicated.

High-­‐z,	
  merging	
  galaxiesは	  
効率良く星が作られている？	
  been adjusted upward by a factor of 2.7 to match those of galactic clouds when
extrapolated to local cloud masses (see the text).

z~0.4-1 Combes ULIRGs
z=1-3 PHIBBS SFG

these galaxies are those determined by GS04 after the upward
adjustment described above.

Another consequ
is that of a corremolecular gas dep
decrease would am
factor we adopted afor our understandi
depletion times forthat describe the loc
we hesitate in dra
particular issue sinof adjustment optio
selected option (2)gas component ofinteresting to note i
dense star-forming
than those estimate
both galaxies and loof any adjustments
As discussed ear
have adjusted thesame constant offs
estimates we aremasses and the ex
the same cloud ma
MCO . To assess thi
masses, MTG , we c
CO observations oobtained CO data f
CfA 1.2 m Millime
12
CO observations
and the integrated C
Applying the stan
1020 cm−2 (K km
integrated intensiti
the mass of each c
to all agree with th

大質量星による分子雲の破壊	
Geen	
  et	
  al.	  (2017)	
３次元輻射流体計算により星団形成、	
  分子雲破壊の計算が行われつつある	  星形成効率の直接計算	
4

S. Geen

20 pc

23.00

22.50

高密度ガス	
100%	

22.25
22.00

10%	
1%	

21.75

低密度ガス	

21.50

log(NH / cm 2 )星形成効率	

22.75

4.2 Myr

8.4 Myr12.6 Myr

16.0 Myr

21.25
21.00
20.75
20.50

Figure 1. Sequence of hydrogen column density maps in simulation L-RT (tff = 4.2 Myr). Each sink particle is shown as a pie chart, whose area is proportional
to the sink particle mass. The fraction shaded in red captures all of the mass accreted in the last 3 Myr, i.e. the fraction in YSOs.

Schaye	
  et	
  al.	  (2010)	

SFR	

フィードバックによる	
  Self-­‐regulaOonのために星形成効率を大きくしても宇宙の星形成率密度は変わらない	

OWLS	
  project	

Figure 10. As Fig. 5, but comparing the SFHs for models with varying
Kennicutt–Schmidt SF laws. Model SFSLOPE1p75 (red, dashed) assumes
a power-law slope n = 1.75 whereas the other models use our fiducial
value n = 1.4. For models SFAMPLx3 (blue, dot–dashed) and SFAMPLx6(olive, dotted) the amplitude of the SF law has been multiplied by factorsof 3 and 6, respectively. All simulations use a 25 h−1 Mpc box and 2 ×5123 particles. At very high redshift, when the SFR in the simulations is

高精度ズームイン計算では?	
Downloaded fro

odel
odel
taltwo
Mpc
and
hold

星形成効率と星形成率密度	高分解能シミュレーションでは？	
Yajima,	
  Nagamine	
  et	
  al.	  (2017)	
高い星形成効率	

200kpc(comoving)	

低い星形成効率	

ガスの柱密度	銀河円盤形成	
高い星形成効率	
 低い星形成効率	
 フィードバックなし	
 2.5
2logΣgas [M¤ pc-2]	

20kpc	

1.5
1
0.5
0
-0.5
-1
3
2.5
2
1.5

ガス円盤の形成	

1

2kpc	
0.5

星形成史	
Yajima,	
  Nagamine	
  et	
  al.	  (2017)	

Halo-­‐11	
  低星形成効率	  フィードバックなし	

安定した星形成?	
間欠的星形成	

低星形成効率モデルの方が結果的にSFRが大きい	
  フィードバック効率が小さくなるため	

星分布のコンパクトネスにも影響	“Cloud内”の星形成効率は銀河の星形成史、	
  形態には大きな影響	  高赤方偏移銀河内の冷たいガス雲の状態(e.g.密度)が鍵？	
  Figure 6. Left: CR7 with the NB921 ﬁlter/Suprime-cam imaging on Subaru, showing the extent of the Lyα but note that NB921 detects Lyα at only 50%
transmission. Middle: HST imaging in YJ, revealing that CR7 clearly splits into three different components which we name A, B, and C. Right: HST imaging in H,
again revealing the three different components in CR7. We ﬁnd that component A fully dominates the rest-frame UV and is coincident with the peak of Lyα emission
and the location at which we detect strong He II 1640 Å emission. Clumps B and C are much redder, and fully consistent with signiﬁcantly contributing to the IRAC
photometry. Note that because of the colors of the B and C clumps, they completely dominate the mass of the system, and thus the actual mass center of the system
would be located between C and B, and signiﬁcantly away from A. This is fully consistent with a scenario in with PopIII star formation is propagated in a wave fromthe central position toward the outskirts.

星種族：初代星銀河は存在するか？	
Sobral	  et	
  al.	  (2015)	
  
1HeII輝線を放射する明るいLAEの発見	
  
(a)
Sky
0
POPIII星を大量に含む銀河？	
−1

Publications of the Astronomical Society of Japan, (2017), Vol. 00, No. 0

10

CR7
Sobral+15 (X−SHOOTER)
T.S. (Pipeline)
T.S. (ESO 2D)
S.M. (ESO 1D)

0
−1
1

d ["]

1

1

fλ, arbit

0
Figure 7. False0color composite of CR7 by using NB921/Suprime-cam
imaging (Lyα) and two HST/WFC3 ﬁlters: F110W (YJ) and F160W (H). Thisshows that while−1and−1component A is the one that dominates the Lyα emission
the rest-frame UV light, the (likely) scattered Lyα emission seems to extend all
the way to B and part of C, likely indicating a signiﬁcant amount of gas in thesystem. Note that the reddest (in rest-frame UV) clump is C, with B having amore intermediate color and with A being very blue in the rest-frame UV.

1

Lyα

5

Pipeline(b)
No He II

Smooth
しかし、再解析・追観測はどれも否定的	(c)sky resNo He II

Shibuya	  et	
  al.	  (2017)	
  	
  1650 (see	
  also	  Bowler	
  et	
  al.	  )	< 2σ
NoFigureHeII
8. HST imaging in YJ and H allows us to physically separate CR7 into

two very different stellar populations and shows remarkable agreement withour best-ﬁt composite SED derived in Section 5.3. While clump A (see, e.g.,
Figure 7) is very blue and dominates the rest-frame UV ﬂux, B+C are red andlikely dominate the rest-frame optical and the mass. Note that the we simply
show the HST data together withsky
our best ﬁt composite model derived in
Section 5.3 which was solely based on the full photometry and did not makeuse of any resolved HST data.

0
have direct evidence of the intense Lyman–Werner radiation(rest-frame 912–1000
Å) from the brightest UV clump. It is0
therefore possible that the other clumps have emitted as much−5
or even more of such radiation a few ∼100 Myr before,
preventing what1210
is now the site 1215
of young massive
stars (A)
to
1620
1630
1640
1220
form before and potentially allowing for that pocket of metal
free gas to remain metal free.
There
λrest [Å]as we detect some faint continuum (spatially
λrest[Å]are of course, other912 Å rest-frame,
potential interpretations of our observations. In Section 6 wevery compact) at rest-frame ∼916–1017 Å (Lyman–WernerFig.
6. Re-analyzed
VLT/X-SHOOTER
of CR7.
left and right panels denote the
VIS and NIR
arms of the X-SHOOTER
spectra.
The blue
discussthe different
potentialspectra
scenariosin Thedetail.
radiation).
X-SHOOTER
data also
providesa lines
NIR spectrum,indicate the X-SHOOTER spectra in Sobral et al. (2015). The red, magenta, and orange lines depict spectra obtained from (1) reducing the raw data with theallowing to investigate the signiﬁcant excess seen in the J-band
X-SHOOTER reduction pipeline ESO REFLEX (Pipeline), (2) stacking of each 2D 1-exposure spectrum reduced by ESO (ESO 2D), and (3) stacking of each 1D
6. DISCUSSION
photometry from UltraVISTA (McCracken et al. 2012; Bowler
1-exposure spectrum reduced by ESO (ESO 1D), respectively. These lines have been smoothed with a kernel of ≃ 0.4 Å width which is similar to that of Sobral
et al. 2014), indicative of emission line(s). No continuum is
et al. (2015). The gray lines
present
un-smoothed
spectrum obtained from our data reduction with ESO REFLEX. The thin blue and red lines indicate sky
6.1.
The theNatureof CR7detected in the NIR spectrum. However, and despite theOH lines in Sobral et al. (2015) and our re-analyzed data, respectively.
top-left panels show the 2D spectrum of the X-SHOOTER VIS arm. The top-right43.93 0.05 Thes−1
CR7,
with
a OH
luminosityof L Lyaand=(c)10
II feature
1640 Å line was
integrationtime,strongHepanelsindicate(a) skyline, (b) un-smoothedsmoothed 2D erg
spectra, allisof which relatively
are obtained low
from ourdata reduction
with aESO
REFLEX.
The
found(∼6σ),
capableofexplaining
theexcess∼3×
moreluminous
than
any
known
Lyα
emitterwithintheat λrest = 1643 Å appears to be made by sky subtraction residuals. The blue ticks indicate the position of He II whose detection is claimed by Sobral etin
al. the J band
epochSeeof Section
reionization(e.g.,
Ouchi et al. 2013).
(see Figure 5). He II can only be produced if the intrinsic(2015).
4.5 for moredetails.
Our optical spectrum shows that the source is very blue
extreme UV spectrum is very hard, i.e., emits a large number of

初代星銀河は存在しないのか？	

o detection of metal lines (see howor metal line detections in CR7). In,
system of star forming galaxies at a
5 kpc. Based on SED modelling the
7 could host a POPIII star cluster

just after virialization is,⇡R4⇡R ↵ n初代星銀河は形成可能か？	
f Ṅ
>+ ⇡R D↵ n
4⇡D
3
ion
esc

ion

2
vir
2

3
vir

B

2
H

2
vir

B

2
H,IGM ,

(1)

Ṅion is the ionizing photon emissivity of the star-forminge.g.,	
  Visbal	
  et	
  al.	  (2016,	
  2galaxy,017),	  	 galaxy to
DYisajima&Khochfar	
  the distance from the(2017)	
  star-forming
the target halo, Rvir is a virial radius of the target halo, ↵Bch a POPIII star cluster, a massiveis the case B recombination coefficient, nH and nH,IGM arecollapse. Visbal et al. (2016) sugTHII	
  >	
  Tvir	
光電離によりミニハロー時代
the meanhydrogen number density in the halo and the interCR7 could be suppressed by photogalacticの星形成を停める。	
  medium, respectively. nH and nH,IGM just depend
hboring galaxies, and that only once3
nough the gas can collapse. Pallot-電離	
on redshift, nH = c ⇥ nz=0c ⇠ 200 is
H,IGM (1 + z) , where7
その後、初代銀河クラスに
the meanover density and nz=0
cm 3 . The
that indeed cosmological hydrodyH,IGM ⇠ 2.0 ⇥ 10
second termon the right-hand side of the above equation isuce conditions as required for the
なった後コラプスさせてスター
the absorption of ionizing photons by the intervening IGM,
sters. However, the mass they find
バーストを起こせばよい	and muchsmaller than the first term. Using above equaby the observations of Sobral et al.
tion we estimate the critical distance for target haloes to be
le in forming a massive POPIII starby a neighbouring galaxy as
e life-time of massive stars, a few
THII	
  <	
  Tionizedvir	✓ ion ◆ ✓◆collapse and the formation of thefesc
SFR
Dcrit < 73 kpc
s to take place. If the gas has suffi0.1
10 M yr 1
mentum it will settle in a large scale✓◆ 1✓◆ 5 (2)3
e place on a viscous time scale (e.g.,Mh1+z⇥n efficient mechanism to fuelコラプス、	  a cen109 M
8time scales and to destroy such discion
fescis the escape fraction of ionizing photons from
POPIII星形成	
a et al. 2006). However, timescales(鄭さんトーク)	
theneighbouringgalaxy (Yajima et al. 2009, 2011, 2014;
pping merger rate at high redshifts
Paardekooper et al. 2015; Wise et al. 2014), Mh is thes process (e.g., Khochfar & Burkert

初代星銀河の形成環境	
4

Yajima et al.

Yajima&Khochfar	
  (2017) 	

Smidt	  et	
  al.	  (2016)	

POPIII	
  gal.?	
  (Dual	
  emiqer)	

Bright	  LBG/LAE/QSO	

HI	
  
21	
  cm	
  signal	

atur
is ve
phot
inde
et al
by u
(200

Hen
betw

LLy↵

The
of F
gala
The
limit
CR7
recen
[OII
enric

HII	
Bright	  LBGs付近の”そこそこ”	  明るいhigh-­‐EW	  LAEsを探す	  ただし、HeIIは暗い、、	Figure 1. Top panel: Halo mass range where angular momentum
transport efficiently occurs due to Ly↵ radiation drag. Dash lines
represent virial temperatures of 1 ⇥ 104 , 3 ⇥ 104 K and 8 ⇥ 104 K,
respectively. Middle panel: Ly↵ luminosity expected in the halo
mass and redshift. Dash line shows the Ly↵ luminosity for halowith the virial temperatures of 3 ⇥ 104 K.

angular momentum is not transported efficiently, i.e., tangtdyn , a rotationally supported disk forms with Rdisk ⇠ Rvir(e.g., Mo et al. 1998), whereis the spin parameter. In
the latter case star formation will proceed quietly follow-

with
high
gala
the o
the
POP
dete
will
grat
The
hour
with
2.4

As P

銀河形成の重要なバリオン物理	

銀河	
フィードバック	＊超新星爆発	  ＊AGN	
  ＊ダストへの輻射圧	
  ＊光電離加熱	  (＊ライマンアルファ輻射圧)	

フィードバック	(藤田さんトーク)	

熱圧	

運動量輸送	

超新星爆発	
  星によるUV電離加熱	
  AGNによるX線電離加熱 	

超新星爆発	
  ダストへの輻射圧	
  ライマンアルファ輻射圧	
  BH	
  disk	  wind	  AGN	
  jet	

Supernova	
  feedback	
熱エネルギーを注入	

運動量を注入	(coolingを一定時間止める)	
  SOnson	  et	
  al.	  (2006)	
  	
  (余計な事はしない)	
  EAGLE	
  project	  (Schaye+)	
  Dalla	
  Vecchia	
  et	
  al.	  (2012)	
  Yajima	  et	
  al.	  (2017)	

(現象論的)	
  Springel	
  et	
  al.	  (2005)	
  	
  illustris	
  project	  (Vogelsberger+)	
  (点源爆発の結果を導入)	
  FIRE	
  project	
  (Hopkins+)	
  
Kimm	
  et	
  al.	  (2014)	
  	

es the expec-

Fo
the gas willgas, tc . If this condition is satisfied, 1then
cm−3
107.5start
K to expandby th
adiabatically, doing work on its surroundings and converting ther×gf ηe (ηH II + ηHe II + ηHe III )asbedecreased
eithermal energy into kinetic energy. The ratio ts /tc can"$ n%2 & T '1/2
(7)
which can−3usuallybeaccomplishedbyincreasing
by increasing tc ,−23H
≃ 1.12 × 10 erg cm s−1
−3
h≃
which
means
the
the temperature, or by decreasing t1s ,cm
107.5increasingK
ave the same
temperature
spatialresolution.
(1 +and/or
XH )(1the+ 3X
熱エネルギー注入モデル(e.g.,	
  Dalla	
  Vecchia	
  &	
  Schaye	
  (12)
2012)	
H)×,nversely proThe sound crossing
time across a resolution element of length h
wher
8XH2
terms of the １：Sound	
  
is
the a
c
rossing	
  
O
me	
  
where ηi =&ni /nH 'is1/2the number density of species i relative toneighbours is
timeh
h
µm
H
	
   hydrogenthe plasma to be fully ionized in the=and we assumed
ts =
1/2
tc
cs We γassumed
kB
T
≃
last
step.
for
simplicity
that
the
Gaunt
factorgf≃	
  $ µ %1/2 & T '−1/2 & h 'ts
1.4[consistentwiththeapproximatedequation
given
byTheuns5
2. ×Column
an(9),= 1.1510 yr density of7.5gas in Halo-11, Halo-11-lowSF,
	
   Figureet al. (1998)
our fiducial
temperature
increase]a primordial
0.6 The10
100andpcthe
variationsof forHalo-11.
top K
panelsshow
filamentary net(8) circles
composition.
Hence,
we radius
obtaina⇡radiativeof topanels
show
thevirial18 time-scale
pc.cooling
Theistime
middleshow
where
cs is O
the
local
sound
speed.
This
relatedthe
２：Cooling	
  
me	
  
(free-­‐free)	
'1/2!t =the& time-step,
physical
7 definition
kpc. The
represent
stellarsurfacede%panels$=
Ct,
classicalof lowerthe
Courant
Ch/c
−1
s
s
n
T
H
wher
7
have
quite
di↵erent
morphologies
dependingonthetreatment
of
≃
3.26
×
10
yr
t
c
where
C (<1) is the Courant
factor. 7.5
h at least one−3
priate1 cm
10 K
We definecoolingtime as'&particle and%1/2radiative$ µ theWef (XH )
ρϵuan be driven(13),
×loss:
value
cooling=0.6 ,
(10)
tc =
0.13
SSP at once,
$
$
only
✓
◆✓
◆ 1/2 to 10
ases that are
where
where
u and $ are the internal energy and radiative cooling rate
臨界密度	mg3 and ρT
inefficient.
−1
−1
per
unit
volume,respectively,
is
the
gas
density.
The internal
well
100
.
(3
=⇠
XH (1
+ Xcm
)
(14)
f (XHn) H
H ) (1 + 3XH7.5
4
10 K
10 M

:

熱圧駆動	
  vs.	
  輻射冷却	

⃝C 20120.752) MNRAS≃ 0.13. 426, 140–158
andf (XTheH =Authors,
C 2012 RAS
Monthly
the Royal
Astronomical
In
thisNotices
wetreat
all Society
star
andtime-scales
gas particles aThus,
wework,canofwrite
for theratio between
the⃝
two

輻射圧（ダスト、ライマンアルファ）	
(一色さんトーク)	
  

hv/c	

hv/c	

hv/c	

Mstar
多重散乱、再吸収により光子の運動量を効率良く	
  ✏⌘ガスに輸送する（柱密度が大きい方が良い）	
Mgas
Lyα RHD of galac?	

V/Vesc	

L
Frad = ⌘c輻射力	Lya輻射流体計算	

Smith	  et	
  al.	  (2017)	
Figure 5. The ratio of the shell velocity v sh to the local escape velocity
v esc ≡ (2 GM<r /r)1/2 as a function of time for different halo masses Mvir

D
輻
現
(一

フィードバック	

熱圧	

輻射圧	

低密度	

高密度	
  
光学的厚さ 大	星形成領域、星間ガスの密度（or	
  柱密度）を知ることが重要	
  (高分解能シミュレーション	
  or	
  高分解能面分光)	

銀河形成の重要なバリオン物理	

銀河	

＊電離度	
  
＊多相構造	  ＊乱流場	  ＊冷たいガス流入	

星間ガス	

星間ガス多相構造	
Hot	

Myers	  (1978)	

Log	
  温度	

HII	

Warm	

星形成	
Cold	

Log	
  ガス密度	分子雲	熱的・重力不安定性	
Mstar
Mgas
L
Frad = ⌘c✏⌘

銀河のダイナミクスと共にこれら不安定がどのように	
  Mstar
起きるか？	
 ✏ ⌘

Mgas

冷却時間
v.s. 音波通過時間 v.s.	
   自由落下時間	
Mstar

M / T 9/2⇢ 2⇤ 3

✏⌘

L
= ⌘
c

M / T 3/2⇢ 1/2✓ ◆gasrad@
冷却時間＜音波通過時間	
<0
@T p

MF

銀河+多相構造計算	

L
Frad = ⌘
9/2
2
3
M
/
T
⇢
⇤
c
熱の圧力	
3/2

1/2

M 2/ T3 ⇢◆M / T ⇢ ✓⇤
@
<0
3/2
1/2
M / T ⇢ @T p
9/2

✓

@
@T

◆

<0

重力	

音波通過時間 ＞ 自由落下時間	

p

荒田さんトーク!	

the halo virial radius, Rv. The entropy, log K 5 log(T/r ), in units of thevirial quantities, highlights (in red) the high-entropy medium filling the haloout to the virial shock outside Rv. It exhibits (in blue) three radial, lowentropy streams that penetrate the inner disk, seen edge-on. The radial flux
_ 5 r2rvr, in solar masses per year per square radian, where
per solid angle is mr is the gas density and vr the radial velocity. It demonstrates that more than
90% of the inflow is channelled through the streams (blue), at a rate thatper cent of the sphere. When viewed from a given direction, the column
density of cold gas below 105 K is above 1020 cm22 for 25% of the area within
the virial radius. Although the pictures show the inner disk, the disk width is
not resolved, so associated phenomena such as shocks, star formation andfeedback are treated in an approximate way only (see density maps andadditional cases in Supplementary Figs 3–5). Kvir, virial entropy.

冷たい(~104K)ガス降着	

uniformly in r, using the fact that the velocity along the streams isroughly constant (Supplementary Information, sections 5 and 6).
This is convolved with the halo mass function23, n(Mv), to give

_ )~
n (M

ð?
0

_ jMv )n(Mv ) dMv
P(M_ ), obtained by integration
The desired cumulative abundance, n(.M_over the inflow rates from M to infinity, is shown at z 5 2.2 in Fig. 4.
_ , the curve referring to M
_ lies safely
Assuming that the SFR equals M
above the observed values, marked by the symbols, indicating that the_ and thegas input rate is sufficient to explain the SFR. However, M
SFR are allowed to differ only by a factor of ,2, confirming our
suspicion that the SFR must closely follow the gas input rate. The
simulated SFR indeed traces the accretion rate to within a factor of
two, but, given that our disks are poorly resolved, we focus here on
the accretion as the more robustly simulated quantity. Because atz < 2.2 the star-forming galaxies constitute only a fraction of theobserved ,1011M[ galaxies24,25, the requirement for a SFR almost_ , based on Fig. 4, becomes even stronger.
as great as M
By analysing the clumpiness of the gas streams, using the sharp_ (r) profiles, we address the role of mergers
peaks of inflow in the M
versus smooth
flows. We evaluate each clump mass by integratingÐ_ (r)=vr (r)) dr across the peak, and estimate a mass ratio
Mclump ~ (Mfor the expected merger as m 5 Mclump/fbMv, ignoring further mass
12 of theA.galaxy
Dekelloss in the clump on its way in and deviations
baryonand Y. Birnboim
fraction from fb. We use ‘merger’ to describe any major or minor
merger with m $ 0.1, as distinct from ‘smooth’ flows, which include‘mini-minor’ mergers with m , 0.1. We find that about one-third of
the mass is flowing in as mergers and the rest as smoother flows.
However, the central galaxy is fed by a clump with m $ 0.1 less than
10% of the time; that is, the duty cycle for mergers is g = 0.1. A
similar estimate is obtained using EPS merger rates7 and starburst
durations of ,50 Myr at z 5 2.5 from simulations26 (Supplementary
Information, section 5).
From the difference between the two curves of Fig. 4, we learn that_ are to be seen during a
only one-quarter of the galaxies with a given M
merger. The fact that the SFGs lie well above the merger curve even if_ indicates that in most of them the star formation is
the SFR is ,Mdriven by smooth streams. Thus, ‘SFG’ could also stand for ‘streamfed galaxy’. This may explain why these galaxies maintain an(中村さんトーク)	

冷たいガス降着の条件	
  (Birnboim&Dekel	
  2003)	

Dekel	  et	
  al.	  (2009)	
ビリアルショック後面の輻射冷却が	  効くか効かないか	

Figure 2 | Streams in three dimensions. The map shows radial flux for thegalaxy of Fig. 1 in a box of side length 320 kpc. The colours refer to inflow rate
per solid angle of point-like tracers at the centres of cubic-grid cells. The
dotted circle marks the halo virial radius. The appearance of three fairly
radial streams seems to be generic in massive haloes at high redshift, and is afeature of the cosmic web that deserves an explanation. Two of the streams
show gas clumps of mass on the order of one-tenth that of the central galaxy,but most of the stream mass is smoother (Supplementary Fig. 6). The>1010M[ clumps, which involve about one-third of the incoming mass, arealso gas rich—in the current simulation only 30% of their baryons turn intostars before they merge with the central galaxy.
452

LAEsやLBGsのハロー質量	
  （ビリアル温度が~10^6	  K）	

©2009 Macmillan Publishers Limited. All rights reserved

LAEsやLBGsはちょうどモード遷移のあたり？	

ハロー質量	

ハロー質量、赤方偏移に依存する	

Dekel&Birnboim	
  (2006)	

T
ical
from
pres
pen
its
et a
diff
in r
acti
dep
give
of h
high

5 F

16	
 Onc
pro

Figure 7. Cold streams and shock-heated medium as a function of halo mass
and redshift. The nearly horizontal curve is the typical threshold mass for astable shock in the spherical infall from Fig. 2, below which the flows are

the
abo

6

Dekel et al.

The desired cumulative abundance n(> Ṁ ), obtained by integration over the inflow rates
from Ṁ to infinity, is shown at z = 2.2 in Fig. 4. Assuming that the SFR equals Ṁ , the curvereferring to Ṁ lies safely above the observed values, marked by the symbols, indicating thatthe gas input rate is suﬃcient to explain the SFR. However, Ṁ and the SFR are allowed
to diﬀer only by a factor of ∼ 2, confirming our suspicion that the SFR must closely follow
the gas-input∼ 2.2
the star-forming galaxies constitute only a fraction
f the Lyα surface brightness of the MW Galaxy with redshift at z ∼ 14, 10.4, 8.5, and 6.2, respectively.
The box rate.
size isBecause
1 Mpc atin zthecomoving
11
25,26
−1
−2
−2
of
the
observed∼
10M
galaxies,
the requirement for a SFR SFR almost as great as⊙
ates the Lyα surface brightness in log scale in units of erg s cm arcsec .
Ṁ , based on Fig. 4, becomes even stronger.

冷たいガス降着は重要？	

Dekel+09	

Maybe	
  yes,	
  because.. 	

*スターバースト	
  *クランピーディスク銀河	  *高輝線幅LAE	
  *クエーサー形成	Di	
  Maqeo+12	

Yajima+15	

High-redshift clumpy discs in simulations

2155

Table 1. Simulation details. Figure 4. Abundance of galaxies as a function of gas inflow rate, n(> Ṁ ). Shown is the comoving

number density, n, of galaxies with inflow rate higher than Ṁ at z = 2.2, as predicted from our

Comoving box size
28.57 Mpc
analysis
Number of DM particles
7 × 106 of the cosmological simulation. The upper curve refers to total inflow. It shows that
DM mass resolution
5.5 × 105 M⊙
galaxieswith Ṁ > 150 M⊙ yr−1 are expected at a comoving number density n ∼ 3×10−4 Mpc−3
Max. resolution (proper)
35–70 pc
27,28 ). Fluxes as high as Ṁ > 500 M yr−1 are anticipated
4Min. mass of a stellar particle(similar
10 M⊙to estimates in other simulations⊙−3−5

at n ∼ 6×10

Mpc

. The lower curve is similar, but limited to gas input by µ > 0.1 mergers.

quasars (one in each row) across five different redshifts (labeled on each of theand the temperature from ∼10 (blue colors) to 10 K (red colors). The quasar
circles. The white points in the bottom panel show particles that are eventually

Ceverino+10	

観測的な証明は？ 冷却光(Lya)、吸収線系(HI)?	

Downloaded from http://mnras.oxfordjourna

hydrogen and helium, metals and molecularhydrogen, represent
photoion- the vicinity of where the observed massive star-forming galaxies can be
The symbols
ization heating by a UV background with proper self-shielding, star
located once their observed SFR is identified with Ṁ . The sBzK/BX/BM galaxies are marked
formation, metal enrichment, stellar mass loss and stellar feedback.
11 The SMGs respectively brighter and fainter than 5 mJy are marked accordingly.11,10 We
SFG.more
In our stochastic star formation model,
than 90 per centK andmore thanof the stars form at temperatures wellsee
below
103 thethat
overall
gas inflow rate is suﬃcient for the observed SFR, but the small margin implies
half of the stars form below 300 K at gas densities higher thanthat theSFR
must ofclosely follow the rate of gas supply. Most of the massive star formers at a
at everytime-step
10 cm−3 . New stellar particles are generated
the zero-level grid, which is about 5given
Myr atz = are
2. We
use a
SFR
expected
to be observed while being fed by smooth flows rather than undergoing‘constant’ feedback model, in which the combined energy frommergers.
By
studying
thecontribution of diﬀerent halo masses to the abundance n(> Ṁ ), we
stellar winds and supernova explosions is released as a constant
the high-SFR
SFGs and SMGs are associated with haloes of mass 1012 − 1013 M⊙ (SI,heating rate over 40 Myr, the typicallearn
age of thatthe lightest
star that
explodes as a type II supernova. The Fig.
heating13).
rate due to feedback
may or may not overcome the cooling rate, depending on the gas
Figure 1. Face-on gas surface density in galaxy C at z = 2.1. The image
conditions in the star-forming regions (Dekel & Silk 1986; Ceverino
demonstrates violent disc fragmentation into transient features and bound&Klypin2009).
We
also
includetheeffectofrunaway
stars
by
4
8
clumps, resembling observed SFGs and theoretical expectations. The size ofassigning a velocity kick of ∼10 km s−1 to 30 per cent of the newly
the image is 15 × 15 kpc. Surface density is in units of log( M⊙ pc−2 ). For
ties of the modeled galaxy from z ~ 14 to z ~ 6 , including, in clockwiseformed
direction,
star
formation
rate,
emergent
Lyαluminosity,
equivalentstellar particles. As a result, these stars can migrate ∼100 pc
comparison, the surface density of molecular clouds in low-redshift galaxies−2 .
away from thedense regions
and explodeas supernovae
in regions
rest frame, and photon escape fraction of Lyα (ﬁlled circles) and UV continuum
Å, opencircles).
Open triangles
squaresis ∼100 M⊙ pcand(1300- 1600)−19times. This
−1 enhances−2 the−2
with lower densities−18
and longer cooling
and
10
erg
s
cm
arcsec,respectively.
d Lyα properties considering the detection thresholds of the surface brightnesswith10
the simulated galaxies. We wish to verify in particular that theyefficiency of feedback in heating the surrounding gas. Finally, inare consistent with the observed global properties of high-redshiftorder to mimic the self-shielding of galactic neutral hydrogen fromstar-forming galaxies, and their main scaling relations.
the cosmological UV background, we assume that in regions where
the gas density is higher than n = 0.1 cm−3 , the ionizing flux shrinks
to a negligible value (i.e. 5.9 × 1026 erg s−1 cm−2 Hz−1 , the value
4.1 Virial properties of the dark-matter haloesまとめと今後	
•星形成	
  高赤方偏移銀河内の星形成は近傍銀河とは大きく異なる可能性がある	
  高分解能シミュレーションor観測で分子雲の物理状態,	
  	
  星形成効率を理解する事が重要	  	
  • フィードバック	  フィードバックは星形成領域のガスの状態（密度、温度、金属量）に依存	
  これらを理解し、どのようなフィードバックがいつごろ支配的か、そして星形
成史、銀河風への影響を解明する	  	
  • 星間ガス	  高赤方偏移銀河内の星間ガス多相構造はまだ不明な点が多い。すばる、
ALMA、高精度シミュレーションを駆使して解明する事が望まれる	
  冷たいガス降着流は銀河進化に大きな影響を与える可能性がある。	
  	
  しかし、定量的な調査はこれから	  End