Formation of supermassive
black holes at high-z
Kohei Inayoshi (稲吉恒平)
Kavli Institute for Astronomy and Astrophysics
Peking University
初代星・初代銀河研究会 @ 名大

初代星研究会と私2010年3月（長崎大学）

M1

2011年1月（愛媛大学）

M2

DCBH (UV/X-ray/CR)

2011年12月（九州大学）

D1

DCBH (cold acc shock)2013年1月（北海道大学）

D2

DCBH (pulsation instability)

2014年1月（鹿児島大学）

D3

DCBH (3D simulation)

2015年1月（東北大学）

PD

BH growth / DCBH

2016年10月（金沢大学）

PD

BH growth / binary BH / GW

2018年2月（呉高専）

PD

BH growth

2019年11月（名古屋大学） 助教

BH growth / DCBH
※ review talkThanks to all of your supports
The Assembly of the First
Massive Black Holes
Kohei Inayoshi,1 Eli Visbal,2 and Zoltán
Haiman3
1

Published in 2020
(will be on arXiv soon…)
Annu. Rev. Astron. Astrophys. 2019.
58:1–74
https://doi.org/10.1146/(please add
article doi)
Copyright c 2019 by Annual Reviews.
All rights reserved

Kavli Institute for Astronomy and Astrophysics, Peking University, Beijing
100871, China; email: inayoshi@pku.edu.cn
2
Department of Physics & Astronomy, University of Toledo, Toledo, Ohio 43606;
email: Elijah.Visbal@utoledo.edu
3
Department of Astronomy, Columbia University, New York, NY 10027; email:zoltan@astro.columbia.edu

Keywords
black holes, cosmology, first galaxies, active galactic nuclei, quasars
Abstract
The existence of ⇠109 M supermassive black holes (SMBHs) within
the first billion year of the universe has stimulated numerous ideas forthe prompt formation and rapid growth of BHs in the early universe.
Here we review ways in which the seeds of massive BHs may havefirst assembled, how they may have subsequently grown as massive as⇠109 M , and how multi-messenger observations could distinguish between di↵erent SMBH assembly scenarios. We conclude the following:Thanks to all of your supports
Special thanks to 大向さん & 細川さん !!

OK, let’s start!

Cosmological QSO population
Development of
Galaxies, Stars, Planets.

Dark ages / EoR

Inflation

today
(z=0)

First SMBHs

First Stars@z>20
13.7 billion years

Quantum
fluctuation

High-redshift monster BHs
11
1011

1.0

0.9

0.8

cosmic age (Gyr)
0.7

11

1077

196SHELLQs/Subaru
QSO samples (z>6)
DELS/DES

Pan-STARS1

10

Pan-STARS1

CFHQ

CFHQ

log (M /Msun)log (M /Msun)

BH mass (Msun)

1088

age (Gyr)SHELLQs/Subaru
DELS/DES

10
1010

1099

0.6

SDSS

SDSS

othersothers

9

8

Ed
di
ng
to
nlim

7

66
*λEdd=1 is assumed for BHs
if not mass measurements

6

7
7

7

redshift z8
8redshiftredshiftz

ite
8

d

Mseed = 10-100Msun
zseed = 35
ε = 0.1
Planck ΛCDM
9

9
9

Inayoshi et al. (2020)

central density and temperature of the cloud to ρ c = 1.67 × 10−20 g
cm−3 and T = 8000 K, giving a mass and radius of 1.17 × 105 M⊙
and 10.8 pc, respectively. Although we here do not impose an external FUV radiation, H2 is collisionally dissociated for ρ " 10−20 g
cm−3 and T " 6000 K. Note that we neglect the dark-matter gravity since the cloud is already bound by the self-gravity of its gas.
Our simulation box size is (50 pc)3 and refinement is controlled byinsisting that one Jeans length is resolved by at least 64 grid cells(e.g. Turk et al. 2012). Under this condition, the simulation uses 23
out of the allowed 25 refinement levels, ensuring we are resolvedby the above criteria at all times and giving a limiting resolution of!0.1 au.
The development of turbulence in the central region of forming
first galaxies has been suggested by numerical simulations (e.g.
Wise & Abel 2007; Greif et al. 2008). In the initial phase of collapse with ∼10−20 g cm−3 , the turbulence is still subsonic in thecloud. To consider the density and velocity perturbations due to theturbulence, we initially impose a subsonic velocity field (the root
mean square of the velocity is set to 0.1cs ) with power spectrum
P(k) ∝ k−4 , which corresponds to the so-called Larson’s law for
the contemporary star-forming regions (Larson 1981). To ensurethat the turbulence is adequately resolved, we select the maximum
k-mode value of 1/10 of the number of cells across the cloud.
We consider the non-equilibrium primordial chemistry of 9−+++
) and 13 hyspecies (H, H2 , e− , H+ , H+
2 , H , He, He , and He
drogen reactions selected to reproduce the correct thermal/chemical
evolution of the warm atomic-cooling cloud (reactions 3, 4, 7−10,
12, 15−18, 28, and 32 in table 2 of Omukai 2001). We adopt
the reaction rate coefficients updated by the following studies: 7–
10 (Coppola et al. 2011), 15 (Martin, Schwarz & Mandy 1996),
17 (Stibbe & Tennyson 1999), and 28 (Ferland et al. 1992). The
four helium reactions originally included in ENZO are also present,
although they are not relevant in our calculation. We initially assume a uniform distribution of ionization degree with 10−4 and H2
molecular fraction with 10−7 , respectively (e.g. Shang et al. 2010).
At high density, the chemical reactions proceed faster than the cloud
collapse and chemical equilibrium is achieved. To smoothly connect the non-equilibrium chemistry to that of equilibrium, we solve

spatial scales; from the top-left clockwise, labution (∼1 pc), the collapsing core (∼0.1 pc),
region, and the protostar formed at the centre (∼portion of the cloud undergoes the runaway clence forms filamentary structures that channcentral region (ρ ∼ 10−8 g cm−3 ), feeding thebottom panel presents the density distribution a
At the end of this simulation, the protostellar mand its radius ≃2 au. These values are consisof the stellar-structure calculation by Hosokaw
assumed a steady and spherical accretion.
Fig. 2 shows the evolution of mass-weight(a) density, (b) temperature, and (c) H2 fractio

Rapid SMBH assembly

BH mass

109-10bright QSOs

MBH,0

gas accretion

BH mergers
(GWs)

6-7

MNRASL 445, L109–L113 (2014)(tH ~ 1Gyr)redshiftseed BH
formation

Figure 1. Density distribution in the plane throug
four spatial scales: from top-left, clockwise: the lation (∼1 pc), a collapsing core by the H− free–bou(∼0.1 pc), the central region around the protostar (∼protostar (∼10 au).

10-20

Rapid SMBH assembly
heavy seeds
・direct-collapse

109-10
BH mass・dense clusters

rapid growth(super-Eddington)light seeds

MBH,0
Eddington limited

6-7(tH ~ 1Gyr)

10-20

redshift

Theory of early BH formation
3

Titans of the Early Universe
2 HOW COULD SUPERMASSIVE STARSFORM?

SMBH formation (not only high-z)

Independent of whether the seeds of the most massivehigh-z quasars are light or heavy, there is a growing
consensus that sustained extremely high accretion rates
are a vital ingredient in their origin. The central problem
is that the growth of black holes via mergers alone isunlikely to produce ≥ 109 M§ black holes by z ≥ 7 (e.g.,
Sesana et al., 2007; Tanaka & Haiman, 2009; Natarajan,
2011, see §4 for further discussion). Rapid accretion rates(& 0.1 M§ yr≠1 ) are also essential if any direct collapse
black hole seeds are to pass through a supermassive(≥ 105 M§ ) stellar evolutionary phase, as the Eddington
limit imposes a maximum nuclear burning lifetime of≥ 2 Myr for these objects.

M.Rees (1984) ARA&A- rapid gas accretion- supermassive star formation

Over the last 2 decades, numerical simulations of thecollapse of primordial halos have grown increasingly
sophisticated. Early numerical work focussed on investigating the collapse of low angular momentum halos(e.g., Loeb & Rasio, 1994; Begelman et al., 2006; Lodato& Natarajan, 2006). It was soon understood that the
dissociation of molecular hydrogen by Lyman-Wernerradiation from the first stars could produce primordial
halos in which collisionally-excited atomic line transitions dominate cooling, allowing the isothermal collapseof such halos at ≥ 104 K (Haiman et al., 1997). Sinceboth the Jeans mass and the infall rate scale with the
cube of the sound speed, and hence as T 3/2 , where T is
the temperature, such “atomically-cooled halos” provide
a natural means for producing the conditions necessary
for the formation of early quasars from massive seeds- dense clusters

Figure 1. Original diagram from Rees (1978, 1984), outlining thepossible formation pathways for supermassive black holes. In thisemission) and H2 cooling (rovibrational line and collision-induced
emission). We also include the suppression of the cooling rate in the
optically thick case by using the optical depth estimated as ρκLc(e.g. Omukai 2001; Shang et al. 2010), where κ includes the H2 -line
opacity and the Rosseland mean opacity considering the H Rayleighscattering, the H2 collision-induced absorption, and the H− boundfree and free–free absorption, and Lc the size of the central core,
which is approximately given by the Jeans length for the sphericallysymmetric cloud in the runaway collapse. Finally, note that we do
not include the heating/cooling associated with the chemical reactions because their effect is negligible during the thermal evolutionof the atomic-cooling clouds.

2 METHODOLOGY
We performed a three-dimensional hydrodynamical simulation of
the gravitational collapse of a primordial-gas cloud using the adaptive mesh refinement code, ENZO (Bryan et al. 2014). Our main
purpose is to investigate the gas dynamics over a wide rangeof the densities (10−21 ! ρ ! 10−7 g cm−3 ). The cloud initiallyhas a spherically symmetric density profile enhanced by a factor f (=1.6) above the critical Bonnor–Ebert (BE) distribution, an
isothermal sphere embedded in a pressurized medium and supported
in marginal hydrostatic equilibrium against gravitational collapse.
According to cosmological simulations (e.g. Wise et al. 2008), at
the centre of a first galaxy with virial temperature "104 K, forming
in an environment where the H2 formation is suppressed, a warm
(T ∼ 8000 K) cloud with ∼105 M⊙ becomes gravitationally unstable at ρ ∼ 10−20 g cm−3 and collapses. Based on this, we set thecentral density and temperature of the cloud to ρ c = 1.67 × 10−20 g
cm−3 and T = 8000 K, giving a mass and radius of 1.17 × 105 M⊙
and 10.8 pc, respectively. Although we here do not impose an external FUV radiation, H2 is collisionally dissociated for ρ " 10−20 g
cm−3 and T " 6000 K. Note that we neglect the dark-matter gravity since the cloud is already bound by the self-gravity of its gas.
Our simulation box size is (50 pc)3 and refinement is controlled byinsisting that one Jeans length is resolved by at least 64 grid cells(e.g. Turk et al. 2012). Under this condition, the simulation uses 23
out of the allowed 25 refinement levels, ensuring we are resolvedby the above criteria at all times and giving a limiting resolution of!0.1 au.
The development of turbulence in the central region of forming
first galaxies has been suggested by numerical simulations (e.g.
Wise & Abel 2007; Greif et al. 2008). In the initial phase of collapse with ∼10−20 g cm−3 , the turbulence is still subsonic in thecloud. To consider the density and velocity perturbations due to theturbulence, we initially impose a subsonic velocity field (the root
mean square of the velocity is set to 0.1cs ) with power spectrum
P(k) ∝ k−4 , which corresponds to the so-called Larson’s law for
the contemporary star-forming regions (Larson 1981). To ensurethat the turbulence is adequately resolved, we select the maximum
k-mode value of 1/10 of the number of cells across the cloud.
We consider the non-equilibrium primordial chemistry of 9−+++
) and 13 hyspecies (H, H2 , e− , H+ , H+
2 , H , He, He , and He
drogen reactions selected to reproduce the correct thermal/chemical
evolution of the warm atomic-cooling cloud (reactions 3, 4, 7−10,
12, 15−18, 28, and 32 in table 2 of Omukai 2001). We adopt
the reaction rate coefficients updated by the following studies: 7–
10 (Coppola et al. 2011), 15 (Martin, Schwarz & Mandy 1996),
17 (Stibbe & Tennyson 1999), and 28 (Ferland et al. 1992). The
four helium reactions originally included in ENZO are also present,
although they are not relevant in our calculation. We initially assume a uniform distribution of ionization degree with 10−4 and H2
molecular fraction with 10−7 , respectively (e.g. Shang et al. 2010).
At high density, the chemical reactions proceed faster than the cloud
collapse and chemical equilibrium is achieved. To smoothly connect the non-equilibrium chemistry to that of equilibrium, we solve

Pop III BHs

Prior starformation(1)

radiation feedback

mini-halo
4
Tvir <~ 10 K

atomic-cooling halo
> 104 K
Tvir ~

Fig. 1 shows the density distribution at the end of the simulation,
where the central density reaches ∼10−7 g cm−3 , for four different
spatial scales; from the top-left clockwise, large-scale gas distribution (∼1 pc), the collapsing core (∼0.1 pc), the central ∼100 au
region, and the protostar formed at the centre (∼10 au). The central
portion of the cloud undergoes the runaway collapse. The turbulence forms filamentary structures that channel material into thecentral region (ρ ∼ 10−8 g cm−3 ), feeding the protostar. The leftbottom panel presents the density distribution around the protostar.
At the end of this simulation, the protostellar mass reaches ≃1 M⊙
and its radius ≃2 au. These values are consistent with the resultof the stellar-structure calculation by Hosokawa et al. (2012), who
assumed a steady and spherical accretion.
Fig. 2 shows the evolution of mass-weighted radial profiles of(a) density, (b) temperature, and (c) H2 fraction. During collapse,

GW recoils

20<M/Msun<140

Collapsing
protogalaxy

3 R E S U LT S

hyper-Eddington accretion
・・
M>>MEdd(4)

H2 cooling(2)pristine gas

JLW > Jcrit
high vbsm
rapid merger

star formation(H2 cooling)
MNRASL 445, L109–L113 (2014)

no101-2 Msun

105-6 Msun
SMS

Figure 1. Density distribution in the plane through the density peak forfour spatial scales: from top-left, clockwise: the large-scale gas distribution (∼1 pc), a collapsing core by the H− free–bound continuum cooling(∼0.1 pc), the central region around the protostar (∼100 au), and the final
protostar (∼10 au).

if N＊>Ncrit runaway

collisions
(3)

yes

Downloaded from http://mnrasl.oxfordjournals.org/ at Columbia University on December 12, 2014

Theory of early BH formation

First galaxies

103-4 Msun

Inayoshi, Visbal & Haiman (2020) ARA&A

Heavy BH Seedsolumbia University on December 12, 20140

the five clouds have different structures of density, velocity, andtemperature. The resulting stellar masses are also different as
indicated in the figure.
After the formation of a protostellar core at the center ofthe collapsing cloud, we switch to the 2D RHD calculations
for each individual dark matter halo and follow the evolution
during the later accretion stages. Figure 4 shows snapshots fromthree of our examined cases, which exemplify the three different
evolutionary paths (P1, P2, and P3). We see that in each casea bipolar Hii region forms (Figure 4(a)), which subsequently
grows at varying rates as the stellar mass increases. The massaccretion onto the protostar is finally shut off by the strong UVradiative feedback caused by the dynamical expansion of the Hii
region (Figure 4(c); see also Hosokawa et al. 2011). Figure 5
shows the distribution of the final stellar masses obtained inour simulations (a summary is given in Table 2). We see alarge scatter of resulting stellar masses, ranging from 9.9 M⊙

spectrumof the stellar-structure
velocity is set tocalculation
0.1cs ) withbypower
Hosokawa
et al. (2012), who
whichcorresponds
to theso-called
Larson’s law forassumeda steady andspherical
accretion.
orary
star-forming
regions(Larson1981). To ensure
Fig.
2 shows theevolutionof mass-weighted
radial profiles ofulenceis adequately
resolved, we
maximum
During collapse,(a) density,(b) temperature,andselect(c) Hthe2 fraction.
e of 1/10 of the number of cells across the cloud.
cosmological scales
der the non-equilibrium primordial chemistry of 9
Figure 3. Number (upper panel) and masses−+++
as a function of time since first sink formati
,
H
,
He,
He
,
and
He
) and 13 hyH2 , e− , H+ , H+
2
HEAT (blue lines). In the lower panel, solid aand total sink masses, respectively.
ons selected to reproduce the correct thermal/chemical
if no H2
the warm atomic-cooling cloud (reactions 3, 4, 7−10,
mass accreted by a sink is immediately
28, and 32 in table 2 of Omukai 2001). We adopt
rated into the protostar. We calculated t
rate coefficients updated by the following studies: 7–
photosphere R⇤ with an analytic fit fromindependent of metallicity for M⇤ .
et al. 2011), 15 (Martin, Schwarz & Mandy 1996),
Omukai 2009).
& Tennyson 1999), and 28 (Ferland et al.monolithical1992). The
present,
reactions originally included in ENZO are also
single protostar
collapse
3 RESULTS
y are not relevant in our calculation. We initially asWe coolingrun a simulation HEAT that includeH
rm distribution of ionization degree with 10−4 and H2
Figure 1. Density
distribution
in termthe inplane
larradiation
(via
the
last
Eq.
normally…
4
action with 10−7 , respectively (e.g. Shang et al. 2010).
~10four spatial scales:
fromKwithout
top-left,clockwisetion
NOHEAT
heating.
The therm
dust is shownin Figure
1. As
ty, the chemical reactions proceed faster than the cloudtion (∼1 pc), aandcollapsing
core
by theH−thefnofine-structure
fragment
line cooling by [C II] and(∼0.1 pc), the centralprochemicalisothermalregion
at TCMBaround⇡ 40 K theat densitieH2equilibrium
cooling is achieved. To smoothly conprotostar (∼10 density,au). the lines become optically thicequilibrium chemistry
to
that
of
equilibrium,
we
solve
2-3
40

35

30

25

z form20

15

10

er of dark matter halos that host star-forming gas clouds.
ows the distribution of redshifts when the central gas density

High-z star formation−3 . The histograms are colored according to the virial masses

ale displayed at the top.
of this figure is available in the online journal.)

gical simulations to follow the formation of thes clouds that gravitationally collapse in the centerr halos. The histogram in Figure 2 shows that our
0 dark matter halos has a wide range of masses–106 M⊙ distributed over redshifts z = 35–11,
h are at z = 20–15. Figure 3 shows an example ofas density concentrations arising in five such dark
ogether with insets of their zoomed-in structure,

ed gas density distribution at z = 25 in one of our cosmological simulations. We show five primordial star-forming clouds in a cube of 15 kpc on ashow the zoom-in to the central 1 pc region of the clouds at the respective formation epoch. The masses of the first stars formed in these clouds are
and 343 M⊙ , respectively.
of this figure is available in the online journal.)

5

~10 K
fragment

Pop III stars

45, L109–L113 (2014)M⇤ ⇠ 100 M

Safranek-Shrader
et line-of-sightal. (2016)
Figure 2. Mass-weighted
projections of gas density (left column) and temperature (right column) at the end of NOHEAT (top three rows)

but cools back to TCMB after reaching dgas and dust collisionally couple. Isothues until reaching densities ⇠ 5 ⇥ 101
marking the opacity limit for fragmenta
depth due to dust exceeds unity ( esc .
loses its efficacy, and the gas begins teffect of protostellar dust heating is molumbia University on December 12, 20140

the five clouds have different structures of density, velocity, andtemperature. The resulting stellar masses are also different as
indicated in the figure.
After the formation of a protostellar core at the center ofthe collapsing cloud, we switch to the 2D RHD calculations
for each individual dark matter halo and follow the evolution
during the later accretion stages. Figure 4 shows snapshots fromthree of our examined cases, which exemplify the three different
evolutionary paths (P1, P2, and P3). We see that in each casea bipolar Hii region forms (Figure 4(a)), which subsequently
grows at varying rates as the stellar mass increases. The massaccretion onto the protostar is finally shut off by the strong UVradiative feedback caused by the dynamical expansion of the Hii
region (Figure 4(c); see also Hosokawa et al. 2011). Figure 5
shows the distribution of the final stellar masses obtained inour simulations (a summary is given in Table 2). We see alarge scatter of resulting stellar masses, ranging from 9.9 M⊙

spectrumof the stellar-structure
velocity is set tocalculation
0.1cs ) withbypower
Hosokawa
et al. (2012), who
whichcorresponds
to theso-called
Larson’s law forassumeda steady andspherical
accretion.
orary
star-forming
regions(Larson1981). To ensure
Fig.
2 shows theevolutionof mass-weighted
radial profiles ofulenceis adequately
resolved, we
maximum
During collapse,(a) density,(b) temperature,andselect(c) Hthe2 fraction.
e of 1/10 of the number of cells across the cloud.
cosmological scales
der the non-equilibrium primordial chemistry of 9
−
+
++
,
H
,
He,
He
,
and
He
) and 13 hyH2 , e− , H+ , H+2
ons selected to reproduce the correct thermal/chemical
the warm atomic-cooling cloud (reactions 3, 4, 7−10,
28, and 32 in table 2 of Omukai 2001). We adopt
rate coefficients updated by the following studies: 7–
et al. 2011), 15 (Martin, Schwarz & Mandy 1996),& Tennyson 1999), and 28 (Ferland et al.monolithical1992). The
present,
reactions originally included in ENZO are also
single protostar
collapsey are not relevant in our calculation. We initially as3 distribution in the plane
rm distribution of ionization degree with 10−4 and H2
Figure 1.3 Density
c
s
・high
mass
/ no
fragments:⇠ ⇢spatial/action with 10−7
, respectively
(e.g.
Shang
et al. 2010). MJfour
p from top-left, clockwise
J scales:⇢ty, the chemical reactions proceed faster than the cloudtion (∼1 pc), a collapsing core3/2
by the H− f/T3(∼0.1 pc),
chemical equilibrium is achieved. To smoothly conMJthe central
cs region around the pro
protostar・verytohigh⇠ (∼10⇠au).
equilibrium chemistrythat ofaccretion
equilibrium, rate:
we solve Ṁacc
40

35

30

25

z form20

15

10

er of dark matter halos that host star-forming gas clouds.
ows the distribution of redshifts when the central gas density

High-z star formation−3 . The histograms are colored according to the virial masses

ale displayed at the top.
of this figure is available in the online journal.)

gical simulations to follow the formation of thes clouds that gravitationally collapse in the centerr halos. The histogram in Figure 2 shows that our
0 dark matter halos has a wide range of masses–106 M⊙ distributed over redshifts z = 35–11,
h are at z = 20–15. Figure 3 shows an example ofas density concentrations arising in five such dark
ogether with insets of their zoomed-in structure,

ed gas density distribution at z = 25 in one of our cosmological simulations. We show five primordial star-forming clouds in a cube of 15 kpc on ashow the zoom-in to the central 1 pc region of the clouds at the respective formation epoch. The masses of the first stars formed in these clouds are
and 343 M⊙ , respectively.
of this figure is available in the online journal.)

545, L109–L113 (2014)t↵Ghigher temperature is requiredWays to suppress H2 formation
SMBH

IMBH

H-cooling halo

SMBH

IMBH

H-cooling halo

Streaming Vel.

Origin of High Accretion :Origin(A) Thermalof High Accretion

Lyman-Werner irradiation

pristine gas

H2 +LW ! 2 H

baryonic streaming motion

pristine gas
10222ve↵ ⌘ cs + vbsm

Bromm & Loeb 2003; Shang +2010; Latif +2013; Johnson +(2013); Regan +2014; Inayoshi +2014; Sugimura+ 2014; Visbal +2015; Latif +2016; Chon+2016; Hirano+2018; Inayoshi+2018; Wise +2019; Luo+2019 etc…

Ways to suppress H2 formation
SMBH

IMBH

H-cooling halo

SMBH

IMBH

H-cooling halo

Streaming Vel.

Origin of High Accretion :Origin(A) Thermalof High Accretion

Lyman-Werner irradiation

pristine gas

H2 +LW ! 2 H

baryonic streaming motion

pristine gas
10
2 of the2two effects
2
Combinationve↵ ⌘ cs + vbsm
leads to H2 suppression further…

see Wenxiu’s talk

Bromm & Loeb 2003; Shang +2010; Latif +2013; Johnson +(2013); Regan +2014; Inayoshi +2014; Sugimura+ 2014; Visbal +2015; Latif +2016; Chon+2016; Hirano+2018; Inayoshi+2018; Wise +2019; Luo+2019 etc…

Metal-free gas collapse phase10 5

if no H2…
temperature (K)

10 4

10

H cooling3

10 2

H2 cooling(Z=0; PopIII)

1

105

Omukai (2001)

1010

number density (cm -3 )

1015

10 20

Metal-free gas collapse phase10 5

if no H2…

H cooling

temperature (K)

10 4

10 3

10 2

H2 cooling(Z=0; PopIII)
Inayoshi, Omukai & Tasker (2014)

1

105

1010

number density (cm -3 )

1015

10 20

No fragments / high Mdot!
Becerra et al. (2015)

0.5 pc500AU

gravitational collapse

Figure 10

Regan et al.500AU(2014)

500AU

accretion disk

・no/weak fragmentation
no/weak fragmentation・high accretion rate

Distribution of gas density in an ACH in which H2 cooling is suppressed by strong LW radiation
with JLW >⇠ Jcrit . The results
2 on the left (Becerra et al. 2015) and right (Regan et al. 2014a)
panels are from two di↵erent simulations, which overall find a similar behavior. The central
collapsing region does not undergo a major episode of fragmentation despite the complex flow
structure caused by turbulence (left panels, down the few 1000 AU scales, and panel labeledpristine & H free gas

No fragments / high Mdot!
Becerra et al. (2015)

0.5 pc500AU

Regan et al.500AU(2014)

500AU

Clump migration is quick!
unlikely to affect the central star
(Inayoshi & Haiman 2014; Sakurai et al. 2015)

gravitational collapse

Figure 10

see Matsukoba-kun’s talkaccretion disk

・no/weak fragmentation
no/weak fragmentation・high accretion rate

Distribution of gas density in an ACH in which H2 cooling is suppressed by strong LW radiation
with JLW >⇠ Jcrit . The results
2 on the left (Becerra et al. 2015) and right (Regan et al. 2014a)
panels are from two di↵erent simulations, which overall find a similar behavior. The central
collapsing region does not undergo a major episode of fragmentation despite the complex flow
structure caused by turbulence (left panels, down the few 1000 AU scales, and panel labeledpristine & H free gas

log r (cm)No fragments / high Mdot!

Figure 5. Profiles of the radial velocity (solid) and tangential velocity (long-dashed) at theparison, one-half of the Keplerian velocity (short-dashed) and sound speed (dotted) are show
Inayoshi, Omukai & Tasker (2014)
Regan et al.500AU(2014)

100
0.5 pc・accretion
sun/yr)M * rate(Msun(M/ yr)

Becerra et al. (2015)

10
10

c3s
Ṁacc ⇠
/ T 3/2
G

11~ 1Msun/yr

0.1
0.1

500AU

500AU

0.01
-2
10

-4

-3

-2

-1

0

1

2

3

4

5

6

log M
M r (Msun)logr (Msun)

5
2
M
⇡
Ṁ
⇥
t
⇠
O(10
)M
Figure 6. Profile of the mass infall rate (Ṁ?∗ = −4πρr
enclosed m
accvrad ) as
lifea function of the
Figure 10

accretion
Thegravitational
horizontal collapse
line shows the critical value 10−2 M⊙
yr−1 , disk
above which the protostar ev
phase (Hosokawa et al. 2012).

・no/weak fragmentation
pristine & H free gas
no/weakfragmentation
exceeds ∼ 10 M⊙ , the SMS collapses
to
a
BH
by
the
generalrela・high accretion rate

Distribution of gas density in an ACH in which H2 cooling is suppressed by strong LW radiation
with JLW >the left (Becerra et al. 2015) and right (Regan et al. 2014a)⇠ Jcrit . The results
2 on5
panels are from two di↵erent simulations,which overall find a similar behavior. The central
collapsing region does not undergo a major episode of fragmentation despite the complex flow
structure caused by turbulence (left panels, down the few 1000 AU scales, and panel labeled

Chandrasekhar 1964).

gure 11sity on December 12, 2014

numerical simulations (e.g.
8). In the initial phase of colulence is still subsonic in theocity perturbations due to thebsonic velocity field (the root
0.1cs ) with power spectrum
e so-called Larson’s law forns (Larson 1981). To ensure
lved, we select the maximumof cells across the cloud.
1055
primordial chemistry of 9
He+ , and He++ ) and 13 hyting s
e
r
c
ar
c
e the correct thermal/chemicaltas
hly ssiveg cloud (reactions 3, 4, 7−10,
g
i
H
a
m
of Omukai 2001). We adoptr・ crit )
e
4
p
4
・ >M
su
by the following studies: 7–10
(Mn, Schwarz & Mandy 1996),
28 (Ferland et al. 1992). The
sunuded in ENZO are also present,
calculation. We initially astion degree with 10−4 and H2
3 Figure 1. Density distribution in the plane through the density peak forively (e.g. Shang et al. 2010).103 four spatial scales: from top-left, clockwise: the large-scale gas distribuMsun/yrs proceed faster than the cloudtion (∼1 pc),1.0
a collapsingcore by the H− free–bound continuum cooling(∼0.1 pc), the central region around the protostar (∼100 au), and the final
achieved. To smoothly conprotostar (∼10 au).Msun/yr
that of equilibrium, we solve
0.1

Rapidly accreting protostar

10

~ 1M

stellar radius (Rsun)

K

K
10
4

/yr

growth to~105Msun

10

5 K

0.03 Msun/yr

10

Pop III stars

Ac

1

11

10
10

10 22
10

ilit
y

AM

gZ
n
i
t
cre

ta
b

10
10

tars
s
S

in
s

10 -3 Msun/yr

R

1022
10

G

stellar radii (Rsun)

10

Teff

00
= 50

10 33
10

stellar mass (Msun)
stellar mass (Msun)

10 44
10

10 55
10

Hosokawa+ (2012,2013)
Inayoshi et al. (2020)

Rapid BH accretion

Super-Eddington accretion
・GR/Radiation/MHD simulations (Ohsuga+ 2009; Jiang+ 2014; Sadowski+ 2015)- photon trapping in a diskṀ
No. 1, 2005

- radiative flux to the poles
1000
3
10

42%

RMHD (Jiang et al. 2014, 2017)

radiative luminosity L / L Edd

Luminosity

at small scales

10%

GRRMHD (Sadowski et al. 2015)

100
100

ṀEdd
is
possible!
SUPERCRITICAL ACCRETION FLOWHow?

photon trap

RHD (Ohsuga et al. 2005)
L/・ vinfc=Mradiation

10
10

11

0.1
0.1
0.1
0.1

0.

99

7
0.
a

Hyper-Eddington

=0

1%

1

1

Rion < RB

0.1%

10

10

・ ・

100

100
・
・

BH accretion rate M / M Edd

M/MEdd

10 33

10

1044

10

accretion

he
the
mu
tiv
bo
ga
iso
ati
tha

ma
ho
are
esc
thr
dis
be
sen
(C
fal
the

ge
be
de
ex

Bondi accretion limit

・・

M>>MEdd OK!
Eddington limit(gravity vs. radiation)

Bondi accretion limit
RB (>>RSch)Gas accretion within RB(gravity vs. gas pressure)

・

GM•
M•
RB = 2 /c1
T1

MBmass inflow from RB・・

M>>MEdd OK!
Eddington limit(gravity vs. radiation)

2
ṀB ' 4⇡⇢1 RB c12
3/2/ ⇢1 M • T 1

Bondi accretion limit
RB (>>RSch)Gas accretion within RB(gravity vs. gas pressure)

photoionization/
heating (UV&X-rays)
・

GM•
M•
RB = 2 /c1
T1

MBmass inflow from RB・・

M>>MEdd OK!
Eddington limit(gravity vs. radiation)

2
ṀB ' 4⇡⇢1 RB c12
3/2/ ⇢1 M • T 1

BH accretion in multi-scales

July 20

Ohsuga & Mineshige-1larger scales: < O(kpc)-6-6log ρ/ρ0log ρ/ρ0-1Model B ~O(100)*RSch ~ 0.1-1
Model Cmpc
smaller scales:

Ohsuga et al. (2011)

25RS

atterns near the black hole for models A, B, and C, from left to right, respectively. Also plotted are normalized
re time-averaged over 6–7 s for models A and C and over 9–10 s for model B.
ne journal.)(e.g., Illustris, EAGLE, FIRE)

standard disks

Slim disk

Shakura & Sunyaev (1973)

Abramowicz et al. (1988)

also expands in theistribution at t = 1 sHere, 512 × 512 gridsve confirmed that thege by a lot even if we= 0.1 RS (1024 × 512
tion and setting thehole region, we startcosmological simulations

vz (z/R)] is the R-component of the velocity. We employ(rc2 , zc ) = (60 RS , 60 RS ) for models A and C, and (rc2 , zc ) =(30 RS , 30 RS ) for model B. That is, in the present study Ṁout
and Lkin indicate the mass and kinetic energy transported upward
only by the high-velocity outflow (vR > vesc ) per unit time. In
addition, we plot the luminosity of the trapped radiation, whichis evaluated by! 2RS

BH accretion in multi-scales

July 20

Ohsuga & Mineshige

larger scales: < O(kpc)

-1
-6

-6

log ρ/ρ0log ρ/ρ0-1Model B ~O(100)*RSch ~ 0.1-1
Model Cmpc
smaller scales:

Ohsuga et al. (2011)

25RS

atterns near the black hole for models A, B, and C, from left to right, respectively. Also plotted are normalized
re time-averaged over 6–7 s for models A and C and over 9–10 s for model B.
ne journal.)(e.g., Illustris, EAGLE, FIRE)

standard disks

Slim disk

Shakura & Sunyaev (1973)

Abramowicz et al. (1988)

also expands in theistribution at t = 1 sHere, 512 × 512 gridsve confirmed that thege by a lot even if we= 0.1 RS (1024 × 512
tion and setting thehole region, we startṀṀ

cosmological simulations

vz (z/R)] is the R-component of the velocity. We employ(rc2 , zc ) = (60 RS , 60 RS ) for models A and C, and (rc2 , zc ) =ispossible! B. That is, in the present study ṀEdd(30 RSis, 30possible!
RS ) for model
Eddoutand Lkin indicate the mass and kinetic energy transported upward
only by the high-velocity outflow (vR > vesc ) per unit time. In
addition, we plot the luminosity of the trapped radiation, whichis evaluated byṀ ' min(Ṁat small scales, ṀBondi )without resolving RB! 2RS

BH accretion in multi-scales

July 20

Ohsuga & Mineshige-1larger scales: < O(kpc)-6-6log ρ/ρ0log ρ/ρ0-1Model B ~O(100)*RSch ~ 0.1-1
Model Cmpc
smaller scales:

Ohsuga et al. (2011)

25RS

atterns near the black hole for models A, B, and C, from left to right, respectively. Also plotted are normalized
re time-averaged over 6–7 s for models A and C and over 9–10 s for model B.
ne journal.)(e.g., Illustris, EAGLE, FIRE)

standard disks

Slim disk

Shakura & Sunyaev (1973)

Abramowicz et al. (1988)

also expands in theistribution at t = 1 sHere, 512 × 512 gridsve confirmed that thege by a lot even if we= 0.1 RS (1024 × 512
tion and setting thehole region, we startcosmological simulations

vz (z/R)] is the R-component of the velocity. We employ(rc2 , zc ) = (60 RS , 60 RS ) for models A and C, and (rc2 , zc ) =(30 RS , 30 RS ) for model B. That is, in the present study Ṁout
and Lkin indicate the mass and kinetic energy transported upward
only by the high-velocity outflow (vR > vesc ) per unit time. In
addition, we plot the luminosity of the trapped radiation, whichis evaluated by

How can we connect the two different scales?
Construct the global solution including RB !
! 2RS

Gas supply from large scales
・ low density case ( ṀB,ini ⇠ 10ṀEdd )Ciotti & Ostriker (2001)
Milosavljevic+(2009)MILOSAVLJEVIĆ, COPark & Ricotti (2011, 2012)

RB

・
MEdd

2.5

−1
−6

Rion > RBaccretionrateṀ (10 M yr )

3.0

2.0

1.5

1.0

1.5
1.0

0.5

L (1040 erg s−1 )

L148

0.5
0.0

0

1

2

3

4

5

0.0

time (10 yr)
time (103 yr)
3Figure 1. Central accretion rate and radiative luminosity as a function of time
for radiative efficiency ϵ = 0.1.

ṀB / ⇢1 M•2 T13/2

2.4. Mesh Refinement and Initial Conditions

burst accretion(radiation heating)We simulate a cylindrical domain in the upper hemisphere0 < (R, z) < 1 pc with 18 levels of mesh refinement. Since
each AMR block contains 8 × 8 zones, the maximum spatial
resolution in the simulation is ∆xEdd∼ 3 × 1012 cm. We requirethat the central hole, with radius rhole = 1014 cm, is always
resolved at the highest level of mesh refinement. We achieve
pseudo-logarithmic gridding by capping the resolution at radiushṀ i <<⌧ ṀGas supply from large scales

lowdensity case ( ṀB,iniVol.⇠69610ṀEdd )LJEVIĆ,・
COUCH,& BROMM
L148

Ciotti & Ostriker (2001)
Milosavljevic+(2009)MILOSAVLJEVIĆ, COPark & Ricotti (2011, 2012)

0.0

RB

sphereSince
spatialrequirealwayschieveradiusreventse. Theartially
0.5, at

−1

2.0

1.5

1.0

1.5
1.0

0.5

0.5
0.0

n of time

・
MEdd

2.5

L (1040 erg s−1 )0.5

Rion > RB−6

1.0

L (1040 erg s−1)

1.5accretionrateṀ (10 M yr )

3.0

0

1

2

3

4

5

0.0

time (10 yr)
time (103 yr)
3Figure 1. Central accretion rate and radiative luminosity as a function of time
for radiative efficiency ϵ = 0.1.

ṀB / ⇢1 M•2 T13/2

2.4. Mesh Refinement and Initial Conditions

We simulate a cylindrical domain in the upper hemisphere
Figure 2. Gas number density n (upper panels) and neutral fraction χH (lower
panels) at two separate instances around t ∼ 4000 yr. The left panels 0show< (R, z) < 1 pc with 18 levels of mesh refinement. Since
the flow during a central accretion minimum; infalling gas is clearly visible ateach AMR block contains 8 × 8 zones, the maximum spatialr ∼ 1.5 × 1016 cm. The right panels show the flow during a central accretion
maximum. The structure near the central axis, in the conical region marked
byresolution
in the simulation is ∆x ∼ 3 × 1012 cm. We require
14
dashed lines, is an artifact of our having set the optical depth to zero along thethatthecentralhole,
withradiusr=
10cm, is always
hole◦axis, for R/|z| < tan 5 , to avoid numerical issues arising from the coordinate
resolved at the highest level of mesh refinement. We achieve
singularity.

burst accretion(radiation heating)hṀ i <<
⌧ ṀEdd

pseudo-logarithmic gridding by capping the resolution at radiusRapid BH accretion in first galaxies
・high density case ( ṀB,ini ⇠ 500 ṀEdd )ity

3

Takeo et al. (2018,2019)
density

10-5

2.0

1
0

0.5

10-9

0.1RB

-0.5
-11

10
-1.0

-2

BH

0.01RB

RB

0.0

・ ac

10-6

c

M

io
v
e
pr

(

-1

0.0

10-7

log (ρ/ρ∞)Rion < RB

1.0log(ρ/ρ∞)Lbol / LEdd2

1.5

Rapid BH accretion in first galaxies
・high density case ( ṀB,ini ⇠ 500 ṀEdd )ity

3

Takeo et al. (2018,2019)
density

10-5

2.0

2

1.5

10-7

log(ρ/ρ∞)Lbol / LEdd

0.5

log (ρ/ρ∞)1

1.0 would
mechanical feedback
Rionto<rapid
RB accretion…
lead

0

see Takeo-kun’s talk -9
-1

0.0

10

0.1RB

-11

10
-1.0

-2

BH

0.01RB

RB

-0.5

0.0

10-6

・ ac

c

M

(

io
v
e
pr

Global hyper-Eddington accretion
8
Inayoshi, Haiman & Ostriker (2016); Sakurai, Inayoshi & Haiman (2016)

10 5

(a)
6

34
1010

accretion rate M/MEdd

・ ・23
・ ・ 1010

4

hyper-Eddington

1010 2110

density (cm -3 )accretion rate M/MEdd

5

3(Rion < RB)

2
1

0.11

sub-Eddington (Rion > RB)

10-1
0

0

(b)
1

1

2
2
time (10 5 yr)

3

3

4

4

Figure 5. Accretion rate history for a higher-mass BH (in the(K)time (105yr)

Global hyper-Eddington accretion
8
Inayoshi, Haiman & Ostriker (2016); Sakurai, Inayoshi & Haiman (2016)

10 5

(a)
6

34
1010

accretion rate M/MEdd

・ ・23
・ ・ 1010

4

hyper-Eddington

1010 2110

density (cm -3 )accretion rate M/MEdd

5

3(Rion < RB)

2
1

What is the feeding rate onto a BH?

0.11

see Haojie’s
sub-Eddington(Rion > Rtalk
B)

10-1
0

0

(b)
1

1

2
2
time (10 5 yr)

3

3

4

4

Figure 5. Accretion rate history for a higher-mass BH (in the(K)time (105yr)

Conditions for rapid accretion
10・analytical argument
✓
◆1/3
3Qion
Rion =
4⇡↵rec,B n21

10
10

GM•
RB = 2 / M• T11
c1

10
BH mass (Msun)

1/3
2/3
2/3
n1 / M• n 1

BH mass (Msun)

/L

1/3

Inayoshi, Haiman & Ostriker (2016)77

hyper-Eddington
Hyper-Eddington

610105510 410103310 2

sub-Eddington
sub-Eddington

10

1033

10

10

1044

1055

density (cm -3 )density (cm-3)
9. Summary of the results for diﬀerent values of the BH
Hyper-Eddington conditions (FigureR< andRBnumber)massionMBH
density of the ambient gas n∞ . Each⇣n1
105 cm 3

⌘✓

M•
104 M

◆

&

✓

symbol indicates whether the final result is hyper-Eddington accretion
3/2(circle), or constant or episodic accretion at the rate of
1 ! (1 − 10) ṀEdd (square). The dashed line marks the√boundary
4 between the two accretion modes: MBH,4 n∞,5 = 1/ 2 ≃ 0.71.
Edd the conditions required
The region above the boundary indicatesto realize steady hyper-Eddington accretion.

T
10 K

◆

Ṁ
ṁ =
Ṁ

5000
500

n ∝ r−2 as expected for an isothermal collapse. The profile is notcompletely smooth due to the presence of small dense clumps within
the halo (as shown in Fig. 1). All simulations show similar profiles
with the maximum number density obtained in each simulation
found to be ≈1 × 1011 cm−3 . The result here are in very good
agreement with those of our own previous simulations as well as

Rapid growth of seed BHs

in a first galaxy (Tvir = 104K)

Takeo, Inayoshi +(2018, 2019)density

Regan et al. (2014)

gas density (cm-3)

xire
re
he
vm

Rion < RB

r -2
n1 , T 1

103cm-3

10pc

M•

radius (pc)halo. As expected for a collapse mediated by atomic hydrogen cooling the temperature
e. Right-hand panel: the density profile for each halo. Number densities reach values
Fig.1:
hermal profile).

Rapid gas accretionis possible for BHs

Gas density halo
distribution
in the outervirial temp
region simulation for hyper-Eddington accre5
tion[3]. WhileMBH. 2 intense⇥ 10radiation
Tvir,4collimatedM towards the poles produces a cavity, rapid accretion from large scales through the equatorialn ∝ r−2 as expected for an isothermal collapse. The profile is not
xicompletely smooth due to the presence of small dense clumps withinrethe halo (as shown in Fig. 1). All simulations show similar profiles
rewith the maximum number density obtained in each simulationhe11 assembly
ative effects
during
the
DCBHs
209
cm−3 . The resultofhere
are in very goodfound to be≈1 × 10vagreement with those of our own previous simulations as well asm

Rapid growth of seed BHs

in a first galaxy (Tvir = 104K)

density

Regan et al. (2014)

Rion < RB

r -2
n1 , T 1

103cm-3

10pc

Downloaded from https://academic.oup.c

Rion (pc)
gas density (cm-3)

Smith et al. (2017)

Takeo, Inayoshi +(2018, 2019)M•

radius (pc)・
halo. As expected for a collapse mediatedN
byion
atomic hydrogen cooling the temperature
e. Right-hand panel: the density profile for each halo. Number densities reach values
Fig.1:
Gas density halo
distribution
in the outervirial temp
Figure 5. Estimated ionization front rIF as a function of the strength ofthehermal profile).
regionsimulation for hyper-Eddington accreionizing source Ṅion , which is proportional to the black hole mass M
• at
Rapid
gas
accretionroughly ∼1048 s−1 M−1.
The
calculation
is based on a Strömgren analysis5tion[3]. Whileintenseradiation
collimated
to⊙
M
.
2
⇥
10
T
M
in which we integrate along rays in the simulation – see equation (3) and BH
vir,4
wards
the
poles
producesacavity,rapid accreisdiscussion.
possiblefor
BHs
the subsequent
The threshold forbreakout
is one to two orderstionof magnitude lower in the polar directions. The arrows indicate the effectof from large scales through the equatorialn ∝ r−2 as expected for an isothermal collapse. The profile is not
xicompletely smooth due to the presence of small dense clumps withinrethe halo (as shown in Fig. 1). All simulations show similar profiles
rewith the maximum number density obtained in each simulationhe11 assembly
ative effects
during
the
DCBHs
209
cm−3 . The resultofhere
are in very goodfound to be≈1 × 10vagreement with those of our own previous simulations as well asm

Rapid growth of seed BHs

in a first galaxy (Tvir = 104K)

hyper-Eddington
accretion terminates

r -2
103cm-3

density

Regan et al. (2014)

10pc

Rion < RB

n1 , T 1
Downloaded from https://academic.oup.c

Rion (pc)
gas density (cm-3)

Smith et al. (2017)

Takeo, Inayoshi +(2018, 2019)M•

radius (pc)・
halo. As expected for a collapse mediatedN
byion
atomic hydrogen cooling the temperature
e. Right-hand panel: the density profile for each halo. Number densities reach values
Fig.1:
Gas density halo
distribution
in the outervirial temp
Figure 5. Estimated ionization front rIF as a function of the strength ofthehermal profile).
regionsimulation for hyper-Eddington accreionizing source Ṅion , which is proportional to the black hole mass M
• at
Rapid
gas
accretionroughly ∼1048 s−1 M−1.
The
calculation
is based on a Strömgren analysis5tion[3]. Whileintenseradiation
collimated
to⊙
M
.
2
⇥
10
T
M
in which we integrate along rays in the simulation – see equation (3) and BH
vir,4
wards
the
poles
producesacavity,rapid accreisdiscussion.
possiblefor
BHs
the subsequent
The threshold forbreakout
is one to two orderstionof magnitude lower in the polar directions. The arrows indicate the effectof from large scales through the equatorialImplications to observations
The& AstrophysicalJournal, 830:53 (7pp), 2016 October 10
ReinesVolonteri

9

n

8
7
6

9

10

11

12

log (Mstar / Msun)

lo

ca

5

at
io

7

10

i
a
in
th
a
a
V
fa
s
q
im
th
p

el

9

Wang et al. (2016)11log (MBH / Msun)log (MBH / Msun)

11

Reines & Volonteri (2015)

lr

ophysical Journal, 813:82 (13pp), 2015 November 10

9

10

11

12

13

log (Mdyn / Msun)

Left: black hole mass vs. total host galaxy stellar mass. All stellar masses are estimated using color-dependent mass-to-light ratios presented in
Figure
MBH vs. Mdyn of the [C II]-detected z>5.7 quasars. The Mdyn forl. (2009; see Sections 2.4 and 3). Our sample of 244 broad-line AGNs for which we estimate virial BH masses from Equation (1) are shown
as red4.
points.
oad-line AGNs and composite dwarf galaxies from Reines et al. (2013) are shown as pink points (including NGC 4395; Filippenko & Sargent1989).
Thez>5.7
quasars
are estimated based on [C II] observations, except for one
xy RGG 118 (Reines et al. 2013) hosting a ∼50,000 Me BH (Baldassare et al. 2015) is the dark green point, and Pox 52 (Barth et al. 2004; Thornton
object, SDSS J114816.64+525150.3 at z=6.42, in which the [C II]-emitting8) is the light green point (see Section 3.1). Fifteen reverberation-mapped AGNs with BH masses taken from Bentz & Katz (2015) are shown as purple
gas(S/S0
at >1.5
kpc scale is turbulent and the CO size is adopted (Riechers et al.
Section 3.2). Dynamical BH mass measurements are taken from Kormendy & Ho (2013) and shown as blue (elliptical galaxies), turquoise
galaxies
cal bulges), and orange (S/S0 galaxies with pseudobulges) points. The gray error bar indicates uncertainties in stellar masses for all points,
andsingle2009; Wang et al. 2013; Willott et al. 2013, 2015; Cicone et al. 2015;
ctroscopic BH masses. The gray lines show various MBH vs. Mbulge relations based on ellipticals and spiral bulges with dynamical BH mass measurements.
Venemans et al. 2016). The red star shows J0100+2802 in this work. For
endy & Ho (2013) “scaled” relation has bulge masses scaled down by 0.33 dex to account for differences in our assumed mass-to-light ratios (seeobjects that do not have an inclination angle estimated with the resolved [C II]
3).

Rarity of IMBHs with
MBH < 105Msun

Overmassive BHs
in high-z QSOs

image, we show Mdyn calculated with different inclination angles (dashedlines). The solid line and the gray region show the local relationship with ±0.3
dex intrinsic scatter. The gray circles are the sample of local galaxies(Kormendy & Ho 2013).

e
h
p
e
F
C
ra
q
+
z
d
s
n
b

Implications to observations

Rapid Growth

The& AstrophysicalJournal, 830:53 (7pp), 2016 October 10
ReinesVolonteri

9

n

8
7

lo

ca

5

at
io

7

10

i
a
in
th
a
a
V
fa
s
q
im
th
p

el

9

Shimasaku
Izumi(2019)
Wang& etal. (2016)11log (MBH / Msun)log (MBH / Msun)

11

Reines & Volonteri (2015)

lr

ophysical Journal, 813:82 (13pp), 2015 November 10

e
h
9
10
11
12
9
10
11
12
13
p
e
Left: black hole mass vs. total host galaxy stellar log
mass. All(Mstellarstar
masses /areM
estimatedpresented in
log
(M
/
M
)
sun)using color-dependent mass-to-light ratios
dynsunFigure
MBH vs. Mdyn of the [C II]-detected z>5.7 quasars. The Mdyn forFl. (2009; see Sections 2.4 and 3). Our sample of 244 broad-line AGNs for which we estimate virial BH masses from Equation (1) are shown
as red4.
points.
oad-line AGNs and composite dwarf galaxies from Reines et al. (2013) are shown as pink points (including NGC 4395; Filippenko & Sargent1989).
Thez>5.7
quasars
are estimated based on [C II] observations, except for one
C
xy RGG 118 (Reines et al. 2013) hosting a ∼50,000 M BH (Baldassare et al. 2015) is the dark green point, and Pox 52 (Barth et al. 2004; Thornton
II]-emitting
object,SDSS
J114816.64+525150.3
atz=6.42,
in
which
the
[C
8) is the light green point (see Section 3.1). Fifteen reverberation-mapped AGNs with BH masses taken from Bentz & Katz (2015) are shown as purpleragas(S/S0
at >1.5
kpc scale is turbulent and the CO size is adopted (Riechers et al.
Section 3.2). Dynamical BH mass measurements are taken from Kormendy & Ho (2013) and shown as blue (elliptical galaxies), turquoise
galaxiesqcal bulges), and orange (S/S0 galaxies with pseudobulges) points. The gray error bar indicates uncertainties in stellar masses for all points,
single- et al. 2013; Willott et al. 2013, 2015; Cicone et al. 2015;2009;andWang
relations based on ellipticals and spiral bulges with dynamical BH mass measurements.
ctroscopic BH masses. The gray lines show various M vs. M+
Venemans et al. 2016). The red star shows J0100+2802 in this work. For
endy & Ho (2013) “scaled” relation has bulge masses scaled down by 0.33 dex to account for differences in our assumed mass-to-light ratios (seeII]objects thatnot havean inclination
estimatedwith theplotted
resolved [CFigure3. do
MBH(panel[a]) andangle
MBH/Mh ([b]),againstMz
3).
with different
inclination
angles (dashedimage, lines
we show
dyn calculated(2002). dI
DottedareMthelocal relation
obtainedby Ferrarese5
lines). The solid line and the gray region show the local relationship with ±0.3
BH
sun
magenta,brighter−27;circles
green,
to −25;fainter thansdex intrinsic
scatter. thanThe grayare−27the sampleof cyan,local galaxiesn2013). in halos that are locked up in SMBHs.
the(Kormendy
fraction&ofHobaryonsb6eRarity of IMBHs with
M < 10 M
BH

bulge

Overmassive BHs
in high-z QSOs

Subsequent growth of BHs…
5

ordinary ACHs (~2σ)

see Chon-kun’s talkrare, massive halos (~5σ)

rapid growth…

no growth…

see Yajima-san’s
talkz
Figure 12

Efficient BH growth in rare regions? (>3σ)

Black hole mass and the accretion rate vs. redshift in two di↵erent simulations. The left panel
shows the growth of a 105 M seed BH in a typical ACH halo, between z = 12 and z = 7.5 (Latifet al. 2018). The host galaxy reaches 3 ⇥ 1010 M at z = 7.5. The right panel shows the evolution

Summary
heavy seeds ~105Msun
・LW radiation / streaming motion

BH mass

109-10MBH,0・high accretion -> DCBHs

rapid growth(hyper-Eddington)
・photo-heating vs. recombination
・ ・・MB/MEdd > 500 is required

light seeds

・overmassive BHs relative to galaxies

6-7

10-20

redshift

KIAA / PKU
Welcome!

Kohei

Wenxiu Li Haojie Hu

Yu Qiu+ 2 undergrads

25 faculties
22 postdocs
80 PhD students

KIAA / PKU
・KIAA fellowships (all fields)
https://jobregister.aas.org/ad/de436c12

・BHOLE postdoc researchers(SMBH, galaxies and coevolution, high-z)
https://jobregister.aas.org/ad/930735dc

・PKING postdoc researchers(Gravitational Astrophysics, Interstellar Medium, Sky-Survey Science)https://jobregister.aas.org/ad/0bb38802

・Kavli Astrophysics Postdoctoral Fellowship (KIAA-IMPU)
https://jobregister.aas.org/ad/3cbd55c3

Thank you!