数値シミュレーションで探る
初代銀河形成過程 	
矢島 秀伸 (東北大学)
長峯 健太郎 (大阪大学)	

第３回銀河進化研究会 6/1-6/3@東北大学	

初代銀河形成研究の位置づけ	
現在	

すばる	

OUCHI ET AL.

JWST:2018~	

TMT:2024~	

21 We estimate the 2σ limits of total magnitudes in the vicinity of this object

to be J = 24.3, H = 24.0, and K = 23.8.

22 http://ssc.spitzer.caltech.edu/mips/apercorr/

LAE,	  LBG,	  SMG	Vol. 696

58◦ 46′ 25′′
Declination1166

and z′ data with the assumption of z = 6.6 LAE spectrum shape(Shimasaku et al. 2006). Figure 3 shows that the profiles ofthis object is clearly more extended than that of point-spreadfunction (PSF) of our image (FWHM = 0.′′ 8). The outskirts ofthe profile reach a radius up to ≃1.′′ 5 and possibly to ∼2′′ . If wedefine the isophotal area, Aiso , as pixels with values above the 2σsky fluctuation (26.8 mag arcsec−2 in NB921), the isophotal areaof Himiko is Aiso = 5.22 arcsec2 in the NB921 image. Figure 4
presents the isophotal area of our z = 6.6 LAE candidates as afunction of total NB921 magnitude and average NB921 surfacebrightness. The average NB921 surface brightness, ⟨SB⟩, isthe value of an isophotal flux divided by the isophotal area,
where the isophotal flux is the one summed over the isophotal
area. We mark possibly extended (FWHM > 1.′′ 2) sources withfilled squares to distinguish between bright pointlike and faintextended sources with a comparable isophotal area. The rightpanel of Figure 4 plots an isophotal area distribution with ahistogram, which visualizes the extended nature of our object
among all of the LAE candidates. Figure 4 indicates that there
are no LAEs similar to Himiko. We confirm that the brightest
source from the previous 0.2 deg2 Subaru Deep Field (SDF)
survey for z ∼ 6.5 LAEs is only as bright as our second brightest
candidate with no significant spatial extent (Taniguchi et al.
2005; Kashikawa et al. 2006), and that our object is distinguishedfrom all the other z ∼ 6.5 LAEs found in the previous studies.
By the comparisons with z ∼ 6–7 LAEs found in previousstudies, we note that the size of this object is outstanding.
The major axes of the isophotal area in NB921 and z′ bands
are ≃3.′′ 1 and ≃2.′′ 0, respectively. Additionally, the NB921 (z′ )image shows potential diffuse components which continuouslyextend by ∼1′′ (∼0.′′ 3) around the isophotal area with a surface
brightness above 1σ sky fluctuation (Figures 1 and 2). Thus, thesize of our object is probably !3.′′ 1 and !2.′′ 0 in NB921 and z′
bands, respectively. Given the fact that this LAE has the unusual
brightness and size, we refer to this object as the giant LAE.
Interestingly, this object is detected at the 4σ level in themedium deep 3.6 µm image from the Spitzer legacy survey ofthe Ultra Deep Survey field (SpUDS; P.I.: J. Dunlop; Figure 1),
while we find only marginal detections (∼2–3σ )21 in the nearinfrared (NIR) images from the UKIRT Infrared Deep Sky Survey Third Data Release (UKIDSS-DR3; Lawrence et al. 2007).
We align Spitzer/SpUDS and UKIDSS-DR3 images with theSXDS optical images, referring a number of stellar objects in thefield. The relative astrometric errors are estimated to be ≃0.′′ 04,≃0.′′ 11, and ≃0.′′ 35 in rms, for optical-NIR, Spitzer/Infrared
Array Camera (IRAC; 3.6–8.0 µm), and Multiband Imaging
Photometer (MIPS; 24 µm) images, respectively. We summarize total magnitudes/fluxes and 2′′ diameter aperture magnitudes of Himiko in Table 1. We define the total magnitude with
MAG_AUTO of SExtractor (Bertin & Arnouts 1996) in the optical and NIR bands. The total magnitudes of Spitzer/IRAC andMIPS bands are obtained from a 3′′ diameter aperture and anaperture correction given in Yan et al. (2005) and the MIPS
Web page22 , respectively. Note that our object is detected in the3.6 µm band, but not in the 4.5 µm band. This is probably
due to the higher noise level in the 4.5 µm band, as we expect
the object to have a fairly flat spectrum at these wavelengths.
Our measured 3.6 µm magnitude of 24.02 (4σ ) would resultin a <3σ detection at 4.5 µm for a flat spectrum (constant AB

初代銀河	  24′′dust continuum

22[C II]( P3/2 → P1/2 )G1Bz = 6.337+50023′′

a22′′hm s
17 06 48.0 47.s 847.s 6Right Ascension (J2000)

velocity dispersion−175

z = 2.092

HFLS3

Figure 2. Composite pseudocolor image of Himiko. The RGB colors areassigned to 3.6 µm, z′ , and NB921 images, respectively. North is up and east isto the left. The image size is 5′′ ×5′′ . The white bar at the bottom right represents
the length of 1 arcsec. The brightest peak with a bluish white color corresponds
to position 1. The position 2 is located 1.1 arcsec west of the position 1.

?	

velocityb 0.′′ 35×0.′′ 23c ∆v = 100 km s−1

初代星	220800d ∆v = 100 km s−1b, High-resolution (FWHM 0.35”x0.23”) maps of the 158-µm continuum (a) and [CII] line emission (b)	  za,obtained~6-­‐11	at 1.16 mm with the PdBI in A-configuration, overlaid on a Keck/NIRC2 2.2-µm adaptive optics

Figure 3: Gas dynamics, dust obscuration, and distribution of gas and star formation in HFLS3.	  z~10-­‐15	image (rest-frame UV/optical light). The r.m.s. uncertainty in the continuum (a) and line (b) maps is 180and 400 µJy beam-1, and contours are shown in steps of 3 and 1σ,!starting at 5 and 3σ,!respectively. Az=2.092 galaxy (labeled G1B) identified through Keck/LRIS spectroscopy is detected ~0.65” north of
HFLS3, but is not massive enough to cause significant gravitational lensing at the position of HFLS3.
Faint infrared emission is detected toward a region with lower dust obscuration in the north-eastern partmagnitude), consistent with our tentative 1–2σ detection andthe large error with 24.62 ± 0.73.
of HFLS3 (not detected at <1 µm). The Gaussian diameters of the resolved [CII] and continuum
2.2. Spectroscopic Confirmation
We carried out spectroscopic
follow-up observationsare
withemission
3.4 kpc x 2.9 kpc and 2.6 kpc x 2.4 kpc, suggesting gas and SFR surface densities of!!!
Keck/DEIMOS and Magellan/IMACS. The DEIMOS observations were conducted with the 830G grating and the GG495
Σgas
1.45. Thex 104 Msun pc-2 and! ΣSFR = 600 Msun yr-1 kpc-2 (~0.6 x 1013 Lsun kpc-2). The high ΣSFR isfilter in the nonphotometric nighton 2007=November
consistent with a maximum starburst at near-Eddington-limited intensity. Given the moderate opticaldepth of! τd<~1 at 158 µ m, this estimate is somewhat conservative. Peak velocity (c) and F.W.H.M.
velocity dispersion (d) maps of the [CII] emission are obtained by Gaussian fitting to the line emission ineach spatial point of the map. Velocity contours are shown in steps of 100 kms-1. High-resolution
CO J=7-6 and 10-9 and H2O 321-312 observations show consistent velocity profiles and velocity
structure (Figures S5-S7). The large velocity dispersion suggests that the gas dynamics in this system
are dispersion-dominated. See Supplementary Information Sections 3 and 5 for more details.
Table 1: Observed and derived quantities for HFLS3, Arp 220 and the Galaxy
HFLS3
Arp 220*
Milky Way*!
redshift6.33690.0181
Mgas (Msun)a(1.04+/-0.09) x 1011
5.2 x 109
2.5 x 109
Mdust (Msun)b
1.31+0.32-0.30 x 109~1 x 108~6 x 107c1010
M* (Msun)~3.7 x 10~3-5 x 10~6.4 x 1010
Mdyn (Msun)d2.7 x 1011
3.45 x 1010
2 x 1011 (<20 kpc)
ef40%
15%
1.2%

〜２０２０年代	

Figure 3. Surface brightness (SB) profiles of our object. The open squares showa SB profile in NB921 image, and the filled circles denote that of a continuumsubtracted (Lyα) image. The dotted line is the Gaussian profile representing thePSF of our image with an FWHM = 0.′′ 8.	  z~20	

2323 See Section 3.1 for a possible inclusion of emission lines in the IRAC bands.

初代銀河はいつ頃どのように形成され、星形成を行うのか？	  
初代銀河内の物理状態は？	  
どのように最遠方ライマンアルファ輝線銀河、ライマンブレイク銀河、サブミリ銀河へと進化していくのか？	

– 10 –14

初代銀河内の星形成	

versus luminosity and is usually parameterized as φ∗ e−L/L (L/L∗ )α , where φ∗ is the normalization,L∗ is the characteristic luminosity, and α is the faint end slope (see Fig. 3). The luminosity function
was assumed to have an MU∗ V of −19.5 and −18.8 at z ≈ 8 and z ≈ 10, respectively (based onpredictions from our z ≈ 4-7 fitting relation9,22 ), while α was taken to be −1.74.
∗光度関数	
Bouwens+11	

宇宙の星形成史	?	

Bouwens+13	

d nkBTHIIxHII
PHII =/r 2
drxHII / rMadau+98	 ✓◆SFR28
LUV=0.7⇥
10erg s 1 Hz 1⌫
1 M yr 1?	

(25)
(26)

Fig. 10.— The U V luminosity density (right axis) and star formation rate density (left axis) versus redshift. The U V luminosity andSFR density shown at z ∼ 9 (large blue solid circle) are from the present work and inferred based on the relative number of z ∼ 8 and z ∼ 9galaxies found within the CLASH cluster program (see §4.5). These luminosity densities and SFR densities are only consideredFigure
down to3. a Ultraviolet luminosity functions at z ≈ 4, z ≈ 7 and constraints for z ≈ 10. Thelimiting luminosity of −17.7 AB mag – which is the approximate limit of both the HUDF09 probe (Bouwens et al. 2011b) and the present
present
constraints on the stepwise ultraviolet luminosity function at z ≈ 10 (black points andsearch assuming a maximum typical magnification factor of ∼9 and limiting magnitude of ∼27.0 mag. The U V luminosity isconverted
into a star formation rate using the canonical U V -to-SFR conversion factors (Madau et al. 1998; Kennicutt 1998). The upper setof points
upperlimits) are new and are derived from the J125 -dropout candidate galaxies over our ultra-deepat every given redshift and orange contour show the dust-corrected SFR densities, while the lower set of points and blue contours show
HUDFslope
WFC3/IR field. These luminosity functions are a function of the absolute magnitude (thatthe inferred SFR densities before dust correction. Dust corrections at z > 3 are estimated based on the observed U V -continuum
distribution and are taken from Bouwens et al. (2012b). At z ≤ 3, the dust corrections are from Schiminovich et al. (2005) and&is, Reddyluminosity)of galaxies (M1600,AB ) in the rest-frame far-ultraviolet. All error bars are 1σ. TheSteidel (2009). U V luminosity density and SFR density determinations from the literature are from Schiminovich et al. (2005) at z < 2stepwise
luminosity
function at z ≈ 10 is also presented as a 1s upper limit, given the uncertainty(black hexagons), Reddy & Steidel (2009) at z ∼ 2-3 (green crosses, Bouwens et al. (2007) at z ∼ 4-6 (open red and blue circles), Bouwens
et al. (2011b) at z ∼ 7 (open red and blue circles), Oesch et al. (2012b) at z ∼ 8 (open red and blue circles), and Oesch et al.of(2012a)
at 10 candidate. The lowest luminosity point has been corrected for incompleteness. Theour z ≈z ∼ 10 (open blue circle and upper limit). Estimates of the SFR density at z ∼ 9.6 and z ∼ 10.8 as derived in C12 based on the z ∼ 9.6
1,9ultraviolet
Z12 and z ∼ 10.8 C12 candidates are also shown (dark green and magenta solid circles, respectively). Conversion to a Chabrier(2003) luminosity functions at z ≈ 4 (blue) and at z ≈ 7 (magenta) are shown for comparison.
∗IMF would result in a factor of ∼1.8 (0.25 dex) decrease in the SFR density estimates given here. The present z ∼ 9 determination
is in
The luminosity
functions fitted here are of the form φ∗ e−L/L (L/L∗ )α (see Fig. 2 legend). This
good agreement with the trend in the SFR density and U V luminosity, as defined by the Oesch et al. (2012a) and Z12 estimates.

UV	the ratio of the selection volumes at the two redshifts.
The U V LFs we input into the simulations have the∗following parameters: MUV= −22.4, α = −2.0, and∗φ = 5.5 × 10−5 Mpc−3 . These luminosity parameters
were chosen to implicitly include a factor of ∼9 magnification from gravitational lensing – which is the median magnification estimated for sources in our selection– so the eﬀective M ∗ at z ∼ 8 is chosen to be ∼2.4
mag brighter than seen in blank field studies (e.g., Oeschet al. 2012b). The faint-end slope we assume approximately matches what we would expect based on the U V
LF results at z ∼ 7-8 (Bouwens et al. 2011b; Oesch etal. 2012b; Bradley et al. 2012) which point to faint-end
slopes α of −2. No change is required in the faint-end
slope α of the LF, due to the perfect trade-oﬀ between(27)

JWST,	  TMT	

analytic representation has recently been shown9 to fit well at z ∼ 7 and later times. The presentof z ∼ 8 sources predicted by this model LF searchshoweda also allow us to estimate the value of L∗ at z ≈ 10 – assuming that the luminosity
resultsvery similar magnitude dependence as seen for our z ∼ 8function at z ≈ 10 has the same values of φ∗ and α ≈ −1.7 as have been found to describe ultraviolet
sample.
Using the above simulation procedure and luminosity
aforemen- function results from z ≈ 7 to z ≈ 4 (refs 1, 2, 9, 27-29). Doing so allows us to constrain
tioned LF, we repeatedly added artificial sources to thereal CLASH observations for all 19 CLASH clusters, created catalogs, and repeated our z ∼ 8 and z ∼ 9 selections. In total, we repeated the described simulations
20 times for each cluster field to obtain an accurate estimate of the total number of sources (selection volumeand redshift distribution) we would expect to find in eachsample, given the described luminosity function.
In total, we find 657 sources that satisfy our z ∼8 selection criteria and 383 sources that satisfy ourz ∼ 9 selection criteria, based on the same luminos-

ダスト	IR	

ALMA	

Egrav ⇠ rdust h(2)HII
2 RvirPth⇠ 4⇡RHIInHkBTHII(6)51
ESN ⇠ 0.007 ⇥ Mcl ⇥ 10 erg(3)
LionIon✏ Prad⇠(7)4t=⌧Salln10 c⇠ 0.46 Gyr(4)1 ✏LUV 重力(促進)	  vs.	  フィードバック(抑制)	
dust
Prad⇠ 1 ⌧dust(8)⌧/ c(5)rdust
LHII
Ly↵ ttrap 2 0.68✏Ly↵ Ṅion ttrap
Ly↵
基本的に星形成率に比例	  ⇠ 4⇡R⇠HIInHkBTHII
(6)
Prad
⇠Pth
(9)
ctcross(Lcion,	  tcrossLIon,	  SN	  rate)	  LionUV
PPrad⇠(7)⇠ 	  Ṁ(10)→	  	  P	  =SNe	  fconv*c	  SvFR	!1/3
LUV
dust
3Ṅ
(1SNあたりMout~10Msun,	  v~3000km/s)	 (8)ion⌧dustP⇠d nkBTHIIxHII
rad
RHII =(11)2c2 33
P=/r(25)
HII4⇡↵n	  fconv(SN)~2x10
B H 	  dyn/SFR	
drLLy↵ ttrap
0.68✏Ly↵ Ṅion ttrap
Ly↵
2Prad ⇠⇠
(9)xHII / r(26)GMfgashcFgrav
tcrossctcross✓◆⇠(12)
星形成率がハロー質量に比例するとして	  2
R
SFR
UV28
重力と比較してみる	
PSNe ⇠ Ṁvirv(10)L⌫ = 0.7 ⇥ 10
erg s 1 Hz 1 (27)Mhfgas
1 M yr 1!1/3
SFR = ✏SF(13)✓◆0.663Ṅiontdyn
RHII =(11)非一様効果、星形成効率、フィードバック効率	  232n⇠
10cm(28)4⇡↵ncrit
333/2thB	  (1H121(通常:	  ~0.1,	  バースト:e~1)	
Mhalo/ feSFf+z)(14)
はシミュレーションが必要	
SF convSF
10s
GMh2fgas⇣ ✏ ⌘3 ✓ f ◆3 ✓ 1 + z ◆ 3/2
Fgrav ⇠ 1 2(12)
SFSN
11M
=4.4
⇥
10M
Rvir
crit
0.1
20
11Mhfgas(29)
SFR = fSF
(13)
tdyn

Key: フィードバック	目的	
初代銀河の星形成史を調べる	  Stellar	  feedbackがどのように進化に影響を与えるか?	  

20Mpcの計算ボックスからいくつかのハロー(Mh>1011Msun	  at	  z~6)を抜き出して、ズームインを使って高精度計算	  宇宙論的流体計算	  Code:	  Gizmo	  (Hopkins	  2015)	  Zoom-­‐in	  Iniaal	  Condiaon	  	  (MUSIC:	  Hahn+2011)	  	  分解能	  Mgas	  ~	  104	  Msun	  MDM	  ~105	  Msun	  Sofening	  ~	  330	  pc	  (comoving)	  Or	  	  minimum	  ~4pc	  (physical)	  using	  	  Adapave	  sofening	  nth,SF=10-­‐100/cm3	  20Mpc/h	

5	

ガスの分布	SNe	  feedback	  なし	w.	  SNe	  feedback	  	500	  kpc	  (comoving)	

Halo0	  (Mh=2x1011	  Msun	  at	  z=6)	

SNeフィードバックなし	

Formaaon	  of	  galacac	  disk	SNeフィードバック	
1kpc	

10kpc	200pc	

30kpc	

Just	  cold	  clumpy	  clouds!	

銀河中心付近	  ガス密度	

星形成率	

星形成史	
No	  SNe	
フィードバックによりガスが吹き飛ばされる	  →星形成が	  間欠的に	  ダイナミカルタイムより短い時間で星形成率が変動する	

星質量とハロー質量	
No	  SNe	
Z=8	SH03	  Wind	FIRE	

Z=8	 Z=7	Z=6	 SNeフィードバックによりZ=7	星形成が１桁程度抑制
Z=6	される	  Behroozi	 Behrooziらのabundance	  matchingによる示唆とコンシステントな値となる	

10	

ガスの質量割合 vs. ハロー質量	
SH03Wind	

No	  Feedback	This	  work	  (FIRE-­‐like)	

フィードバック:	  	  ハローの成長とともにフィードバッ
クを受けたガスをハロー内にトラップ出来るようになっていく	  	  No	  フィードバック:	  ハローの成長とともに効率的にガスが星にconvertするためにgas	  fracaonが小さくなる	  考察	超新星爆発(FIRE:	  P=eSF=0.1	

(FIRE:	  P20xPej)	
=20xPej,	  fsf =0.05)	

(P=Pej)	
Tvir =20000K	

ダストへの輻射圧（光学的厚さ〜１）	

電離光子輻射圧	

より大きな銀河では	星形成が安定化?	
GN-­‐z11(Oesch+2016)	
IOK-­‐1	  Himiko	Mh~1012Msun	  at	  z~6	Mh~2x1011Msun	  at	  z~6	1012Msunのハローはz<9でSFRが安定(ハロー質量が2x1011Msun程度の時)	  ただし、ハロー中心はガスがフィードバックで除かれる	
Massive	  haloでもGN-­‐z11の星形成率は作れない	Phase	  diagram	
Z=8	

電離加熱	
Warm	

ISM多相化	

星形成	
Cold		  log	  ガス密度[cm-­‐3]		  log	  Nsph/dSbin		  log	  温度[K]	

フィードバックにより加熱	

環境効果:	  UV編	紫外線背景放射	

巨大電離バブル	

HII	HI	

自己遮蔽	IGM電離構造	
15	光電離率	  [s-­‐1]	

紫外線背景放射	10-810-910-10x10000	
x1000	

SFR=1Msun/yr,	  D=7kpc	

10-1110-12

SFR=10Msun/yr,	  D=7kpc	

PHII =SFR	 D	
d nkBTHIIxHII/r 2(25)dr
基本モデル	  xHII / r(26)10-13✓◆(Faucher-­‐Giguere’s	  model)	
SFR
UV28
L10=
0.7
⇥
10erg s 1 Hz 1 (27)⌫ -14
1 M yr 1
自己遮蔽臨界密度	✓◆0.66-1510nth ⇠ 10 1 cm 3(28)10 12 s 1✓◆3 ✓◆ 3/2
10-1656 117⇣ ✏SF ⌘38 fSN 910111+z
Mcrit = 4.4 ⇥ 10 M
0.1
2011
赤方偏移	(29)sMej + Menc(Rcool )1213

星形成史	
100UVB加熱によりさらにSFRが下がる	

星形成率[Msun/yr]	

10Jx103	

10.1

0.01

0.001

6789

赤方偏移	

101112星質量	
Log 星質量[1010Msun]	

10.1

0.010.001星質量はほとんど変わらない	
0.0001

6789

赤方偏移	

101112

まとめ	
宇宙論的流体計算+ズームイン初期条件による高精度シミュレーションで初代銀河の形成、星形成を調べた	  	  ・超新星爆発によるフィードバックで星形成は大きく阻害され、星形成史は間欠的になる	  ・ハロー質量が1012Msun以上になると安定的に星形成が続く	  ・紫外線背景放射が強い環境ではstellar	  feedbackによりガスが低密度になるとUVBによりさらに加熱が効いて星形成率が下がる	  	  今後の予定	  環境の違い、ハロー質量の違いが観測されているLAEs,	  LBGs,	  SMGsにどのようにつながっていくのかを多波長輻射輸送計算とサンプル数を増やして調べる