ICRR共同利用小研究会「高エネルギー現象で探る宇宙の多様性I」

宇宙の大規模構造形成・銀河形成の観点から

Anatomy of Cosmic Gas— gal. cluster, SN / AGN / CR feedback
Ken NAGAMINE(Osaka / K-IPMU / UNLV)
長峯健太郎

Keywords:「多波長・マルチメッセンジャー天文学」「幅広いスケールの分野」「多角的な視点」

Outline• Cosmological Structure Formation — halo formation, nonlinear P(k) 10min• Distribution of Cosmic Gas — shock heating, cold ow, gal. cluster 10min• SN & AGN feedback — SF quenching, metal enrichment, topheavy IMF at high-z, SMBH—gal co-evolution

15min
15min

• Cosmic Ray Feedback

15min• おまけ（宣伝）

2min

fl

~65minHistory of Structure Formation in the Big Bang Universe
Cosmic microwave
background radiation:

Dark Age

Accelerated expansion
by dark energy~380,000 years after the
Big Bang

Big BangAge of the universe ~13.8 billion years
3© NASAbecause of their extremely
he SDSS, APM and angued as measuring the largeth constant bias, since thefor the red-tilting effect
The Percival et al. (2001)
cannot be directly plotted
licated window functions.
ure 35, but restricted to ar spectra, those for SDSS,e power spectra are decoreye” when examining thismps and wiggles between``Concordance ΛCDM model ”

WMAP, Planck
SN Ia(⌦M , ⌦⇤ , ⌦b , h,

8 , ns ) ⇡ (0.3, 0.7, 0.04, 0.7, 0.8, 0.96)(the so-called “737 cosmology”)⌦DM ⇡ 0.26

• Successful on large-scale (>1Mpc)• Can we understand galaxy formation in this context?
Fig. 36.— Same as Figure 35, but restricted to a comparisonof decorrelated power spectra, those for SDSS, 2dFGRS and PSCz.
The similarity in the bumps and wiggles between the three power
spectra is intriguing.

Matter power spectrum

FFT
simulatelaxy power spectrum measurerne in mind when interpretingm measurements are those from
effect of luminosity dependent
the APM survey (Efstathiou
oints are from Tegmark et al.
ange 21 < r ∗ < 22 — see also

Tegmark+ (2004)

sFig. 37.— Comparison of our results with other P (k) constraints.
The location of CMB, cluster, lensing and Lyα forest points in this“Back-bone of structure”

nonlinear P(k)linear P(k)Continued support for ΛCDM

Chabanier+’19

Halo bias

nonlinear growth∝ k −2

k

1.5

halos

P(k)

dark matter

linear P(k)∝ k −3

MultiDark sim. Klypin+’16

1st BAO peak@ k ∼ 0.07

Halo bias

2nd peak@ k ∼ 0.13

halos

dark matter

Klypin+’16

Ishiyama+21

The Uchuu (宇宙) simulation
Ishiyama+21

50 M lyr

100 million light yr

10 M lyr

1 billion light yr

First Galaxy Formation in Atomic Cooling Halos`First Star’ DM Halo

Cooling Rate

~106 M⦿
atomic cooling(SF via H2 cooling)

H2 cooling`First Galaxy’ Halo~108 M⦿
Atomic Cooling Halo

Tvir ~ 104 K

Bryan & Norman ’98

814

cf. Tvir ∼ 10 K for M = 10 M⊙(galaxy clusters)

k
m
s
/
Mass-scale of Galaxies, Clusters
1000
100
10

Cooling FunctiontlD o ., ( :l k r /P r n ;t /" )equilibrium curve

tcool = tdyn

t, .
H'

!n )

Brems

b0

gal clusters[erg cm3 s-1]

: cooling rate per unit vol. [erg cm-3 s-1]

104

t

Lo6

l0 7

r /K
Blumenthal+’84;

Peacock textbook, p.572

Gas Temperature

Cosmic Gas：phase diagramWarm-Hot Ionized Medium(WHIM)

Baryon overdensity

Shull+’12

Cosmic Gas
Warm-Hot Ionized Medium(WHIM)

in gal. groups & clusters

CGM
Shock-heated gas
in galaxy clusters& groups

Hydrogen neutral fraction

Highly ionized gas

SN feedback
Choi & KN ‘09(in DM halo)

Shull+’12

Missing Baryon Problem

Shull+’12

Movies: zoom-in sim
Gas Density

log(Temperature)

AGORA L12 GADGET3-Osaka sim.
cf. Roca-Fabrega+21

log(Metallicity)

Shimizu, KN+19

https://sites.google.com/site/santacruzcomparisonproject/1st-order Galaxy formation
Phase Diagram
DM halo

DM halo forms.
Virial shockheatingTvir

gas accretion&
shock-heating

IGM

COLD FLOW

Cooling,dissipation, &disk galaxy forms

Rees & Ostriker ʼ77, White & Rees ʼ78, Fall & Efstathiou ʼ80, White & Frenk ʼ91, Mo, Mao & White ʻ98

Metal absorption lines

Tumlinson+’17, ARAA

Necessity of Feedback
SNe & AGN

e
e
F
SN

?
k
c
dba

DM halo mass function n(M ) / M

2<latexit sha1_base64="(null)">(null)</latexit>

α ∼ − 1.2
@z=0Observed
Galaxy Stellar MF (LF)

(cf. Press-Schechter ’74; Sheth-Tormen ’01)
Tinker+, ….

AGN Feedback?
Schechter functiondnMMΦ(M) == Φ*
exp −( M* )( M* )dM
αlog (Mass)𝛥

log ( dn/dlog M )Number density of gal. per logM

Shaping Galaxy Mass Function via FeedbackdnLLΦ(L) == Φ*
exp −( L* )( L* )dLαStellar-to-Halo Mass Ratio (SHMR)~1012M⦿SN Feedback

AGN Feedback?

Abundance Matching

Behroozi+’10, ’13, ‘18(cf. Ilbert+’10; George+’11; Leauthaud+’12)

SN feedback ef ciency(IMF) &CC SNe:~30% of this couples to ISM as kinetic E. of galactic wind (GW)

energy—massdeposition rate
2ϵw
vw =
( η )

1/2

GW ef ciency:−6

ϵw ∼ 10fi)

fi

mass-loading factor:

η∼2

−1

c ∼ 300 km se.g. Springel & Hernquist (2003
Cen, KN, Ostriker (2005)

SN vs. AGN (SMBH) feedback(Etot,II ~ 100 foe ~ 1053 erg)

EK per SN II ~ 1 foe ~ 1051 erg

Type II SN occurs ~0.01 per M⦿ of SF for an IMF → 1048 —1049 erg / (1M⦿ of SF)
48

49

10 − 10ϵK,SN ∼
M⊙c 2

48

49

10 − 10∼1033+21

−6

∼ 10

−5

− 10

(per 1M⦿ of SF)

MBH
3
⇠ 10
AGN: 10% of Ṁacc c & feedback ef ciency ~0.1 &M?
2<latexit sha1_base64="grwfTjqcPpCEzGt9HCLKvQHvcTg=">AAAB/nicbVDLSsNAFJ34rPUVFVduBovgqiRV0JUU3LgRKtgHtDFMppN26EwSZm6EEgL+ihsXirj1O9z5N07bLLT1wIXDOffO3HuCRHANjvNtLS2vrK6tlzbKm1vbO7v23n5Lx6mirEljEatOQDQTPGJN4CBYJ1GMyECwdjC6nvjtR6Y0j6N7GCfMk2QQ8ZBTAkby7cNeP4bsNveznpKYUJpj+lDz7YpTdabAi8QtSAUVaPj2l3mHppJFQAXRuus6CXgZUcCpYHm5l2qWEDoiA9Y1NCKSaS+brp/jE6P0cRgrUxHgqfp7IiNS67EMTKckMNTz3kT8z+umEF56GY+SFFhEZx+FqcAQ40kWuM8VoyDGhhCquNkV0yFRhIJJrGxCcOdPXiStWtU9q9buziv1qyKOEjpCx+gUuegC1dENaqAmoihDz+gVvVlP1ov1bn3MWpesYuYA/YH1+QO9upVR</latexit><latexit sha1_base64="kf3ywtCYgpYe2jIt+gtbUuBYSKs=">AAACEHicbZDLSsNAFIYn9VbrLerSzWAR3ViSVtCVFN10I1SwF2himEwn7dCZJMxMhBLyCG58FTcuFHHr0p1v47TNQlt/GPj4zzmcOb8fMyqVZX0bhaXlldW14nppY3Nre8fc3WvLKBGYtHDEItH1kSSMhqSlqGKkGwuCuM9Ixx9dT+qdByIkjcI7NY6Jy9EgpAHFSGnLM4+dQCCc3nipIzi8amSZZkcqJDLoSMqhbd2np7XMM8tWxZoKLoKdQxnkanrml9OPcMJJqDBDUvZsK1ZuioSimJGs5CSSxAiP0ID0NIaIE+mm04MyeKSdPgwioV+o4NT9PZEiLuWY+7qTIzWU87WJ+V+tl6jgwk1pGCeKhHi2KEgYVBGcpAP7VBCs2FgDwoLqv0I8RDohpTMs6RDs+ZMXoV2t2LVK9fasXL/M4yiCA3AIToANzkEdNEATtAAGj+AZvII348l4Md6Nj1lrwchn9sEfGZ8/xI+cYA==</latexit>

ϵK,AGN ⇠ 0.1 ⇥ 0.1 ⇥ 10 3 ⇠ 10 5<latexit sha1_base64="Joj/pn5CjHoXau5mZp7/pKzIQcA=">AAACGnicbVDLSgMxFM3UV62vUZdugkVw4zBpFV1JwY3LCvYBnbFk0kwbmnmQZIQyzHe48VfcuFDEnbjxb0yns6itBwLnnnMvN/d4MWdS2faPUVpZXVvfKG9WtrZ3dvfM/YO2jBJBaItEPBJdD0vKWUhbiilOu7GgOPA47Xjjm6nfeaRCsii8V5OYugEehsxnBCst9U3kSBZA20LQUSygcp4i+yE9q2cQ5j2z8iLrm1XbsnPAZYIKUgUFmn3zyxlEJAloqAjHUvaQHSs3xUIxwmlWcRJJY0zGeEh7moZY73bT/LQMnmhlAP1I6BcqmKvzEykOpJwEnu4MsBrJRW8q/uf1EuVfuSkL40TRkMwW+QmHKoLTnOCACUoUn2iCiWD6r5CMsMBE6TQrOgS0ePIyadcsVLdqd+fVxnURRxkcgWNwChC4BA1wC5qgBQh4Ai/gDbwbz8ar8WF8zlpLRjFzCP7A+P4FT72dSQ==</latexit>

Probably depends on
various other factors:
fi

χ≡

ϵK,SN

( ϵK,AGN )(per 1M⦿ of SF)= χ(z, Mhalo, M*, MBH, SFR, C, ⋯)Three Revolutions in Cosmological Hydro Simulations
1990’: 1st
Revolution

Larger scale, medium resolution
w. subgrid models

2012
3rd Rev.

Zoom-in cosmo. sims. w.
better sub grid models
Resolution ~ 10-100pc

e.g. Cen, Ostriker ’92-’9
Katz+ ’96

e.g. KN+ ’01, 04, 0

IC code : GRAFIC (Bertschinger
MUSIC (Hahn & Abel ’11))

Resolution ~ kpc

3

Resolution ~100 kpc

6

~

~

First cosmological, butcoarse calculation

2000
2nd Rev.

Springel & Hernquist ’03

6> 10 K•••© Oku+’21, in prep.

Temperature

Gas Density

t=0 to 1 Gyr

Metallicity

Isolated
Galaxy Tests

Shimizu, KN+19

Shimizu, KN+19

Keller & Wadsley ‘14

Issues with GW Feedback• Physical state of out owing gas (multi-phase, ρ, T, vel., M, p, E)• Acceleration sites (disk? near SNe? above mid-plane? CGM?)• Ultimate fate of out ows (unbound? recycled?)• Morphologies (biconical, spherical, lamentary, clumpy)

fi

fl

fl

• Acceleration mechanism (mechanical, radiation, CR)

Multi-phase out ow from galactic plane: cut-out sim.
fast-moving,
low-density out ow

as well as warm fountain(fall-back)Athena MHD code

4kpc

1 x 1 x 9 kpc domain

 

fl

fl

 

Kim & Ostriker ‘18

Resolved SN-driven winds
face-onedge-onlarger scale

GADGET-3 sim.
dwarf gal.
Mvir = 1010 M⊙
Rvir = 44 kpc
c = 17
mgas = 1 M⊙
ϵg, hs = 0.3 pc∼ 106 K

Superbubbles of ~few 100pc
break out from the disk
& expands.

Predicted winds are weaker
than in cosmo. simulations.

Hu ’19cf. Hill+12; Girichidis+16; Martizzi+16; Gatto+17; Hu+16,17; Li+17; Kim & Ostriker ’18

Convergence of resolved sim
Expected KE with 0.28 ESN

underestmate by x3

Injected terminal momentum

Hu ’19

SPH mass resolution: 100, 900, 1000 M⦿
spatial resolution: ~ pc

kpc-scale superbubble

GIZMO SPH

FIRE-2 simulation (Ma+20)

RAMSES-RT
SPHINX
simulation
Katz+19, 21• importance of low-Z CCSN
• top-heavy IMF w. reduced[C/O]• [O I] could be useful too.
• PDR & clumpy ISM are stillunresolved.

 cf. Vallini+15; Pallottini+14,15,17,19; Moriwaki+18; Arata+20 …

AGORA Paper III: Cosmo-Run• 4 calibration steps• only in the 4th step, we turn on our favorite
SN feedback model.

• the only constraint:
9 −1M⋆ ∼ (1 − 5) × 10 h M⊙ targeting theabundance matching result at z=4.

Roca-Fabrega+21

AMR codes

SPH codes

z=8

z=7

z=6

z=5

z=4

Roca-Fabrega+21
4th AGORA paper will be CGM metallicity (e.g. O IV emission diagnostic; Strawn+21, in prep.)

AGN feedback(Active Galactic Nuclei)

2003; Scannapieco & Oh 2004; Murray et al. 2005; Springel
on larger,
et al. 2005; Di Matteo et al. 2005; Booth & Schaye 2009).
chev et al.
AGNs are powered by mass accretion onto a super massive
3; Murray
2826
KUROSAWA,PROGA,&NAGAMINE
Some
early
work
with
4⇡GM
⇢
BH
(SMBH).
To
illustrate
howthegrowth
of
SMBH
can
be
Begelman Bondi-Hoyle-Lyttleton
BHṁa = ↵ 2
↵
⇠
100
2
3/2
2002),
and
∼0.1
or
∼0.2
(Yu
&
Tremaine
2002).
On
the
other
in
the
equatorial direction due to the ge(cs + v )can be characterized, letself-regulated
and how AGN feedback; Thacker massaccretion rate:
hand, Cao & Li (2008) find that !r is relatively low (∼0.08) for
effect.
826 theKUROSAWA,PROGA,& NAGAMINE89
us first express
radiation
luminosityduetoaccretionasal. 2007;
MBH < 10 M" and relatively high (!0.18) for MBH ! 10 M" .
In all simulations performed by KP
Since
the exactof howtheTremaine
accretion2002).
luminosityof otherfrom an
inflowwith sub-Keplerian2009, andRadiative output from2002),and mechanism∼0.1 or ∼0.2
(Yu 2&On thein theequatorial
direction due tr
accretion:a BHhand,couples
to&thesurrounding
gas
is
not
well
known,
SDH05
with
high
temperatures
at
large
radii
L
=
#
c
,
(1)
Ṁ
Cao
Li
(2008)
find
that
!
is
relatively
low(∼0.08)foreffect.
arrapport thisand M
BS09<simply
a
wide
range
of
the
polar
angles,wh
8 assume that La couples only thermally (and 9
10 M" and relatively high (!0.18) for MBH ! 10 M" .
In all simulations performed
BH
isotropically)tothesurrounding.
Using
Equation
(1),
the
rate
of
in
a
relatively
narrow
polar
angle
ran
emission
Since thethe
exactsimplest
mechanismassumptionof how the accretion
luminosityof lumifrom an inflow with sub-Kep
where
we
invoke
such
that
the
deposition
tothesurrounding(theAGN
feedback
rate)
Figure2).
However,forthesuper-Ed✏
radiative ef energy
ciency1rM-σrelationa BH couples to the✏surrounding
gas
is
not
well
known,
SDH05
with
high
temperatures
at
larg
BH
Shakura&Sunyaev’73
N spectra(rest-mass
⇡
10
in
SDH05
is
writtenastemperature
at large radii, the inflow prnosityis
proportional
to
the
mass
accretion
rate
(
Ṁ
)
and
a
a
energy conversion)
and BS09 simply assume that La couples only thermally (anda wide range of the polar ang
the equatorial plane, resembling a thinmmediatetothesurrounding.
Using
Equation
(1),
the
rate
of
in
a
relatively
narrowpolaran
2 efficiency
radiative (orisotropically)
the rest-massconversion)
(#
).
Both
#
r
r
Ėf = !f La = !f !r ṀBH c ,(12)outflow arises in a wide range of radii
energy deposition to the surrounding (the AGN feedback rate)
Figure2).
However,
for
the
s
−1
rview). In
right
panel
in Figure 2). The geometry
and
Ṁ
are
uncertain.
For
example,#ranges
from
∼10
in
a
2
a
r
in SDH05
iswrittenastemperatureatlarge
radii,
thei✏
Feedback ef
ciency
Ė
=
✏
L
=
✏
✏
Ṁ
c
is
the
efficiency
of
the
AGN
energydeposition
to
where
!
f
f
outflowsolutionisverysimilar
to
a
ra
f
f
a
f
r
a
−11
)
of
the
BH
standard,theradiatively
disk
to ∼10for
sphericallytheequatorial
plane,
resemblinsurroundingefficientgas, and isthina freeparameter
which
is
to
a luminous Keplerian accretion disk (2
Ėf = !f LaBS09= !f !rfind
ṀBHthec , models with (12)Proga outflow
arises
in
a
wide
range
o
e velocity
beconstrainedbyobservations.
et
al.
1998;
Proga
&
Kallman2
symmetric accretion from a low-density medium (e.g., Shakurarightsolutions,panel in Figure
2).
The
g
!
=
0.15
matches
observations(e.g.,
the
Magorrian
relation
f
is
only
very
EddingtonṀoutso-calledisrelations)the efficiency
ofMeszaros
theenergydeposition
where& Sunyaev
Shapiro1973;
1975;
Soltantoweaker
1982;
outflowsolution
is verysimilarand 1973;
the MBH!f–σvery well,andAGNsimilarly
SDH05
find
correlation
is mainly
causedbythesurrounding
gas,
and
is
a
free
parameter
which
is
to
ardt et al.
a luminous
Keplerian
Springel,Di
Matteo+
’05
that
!
=
0.05
matches
observations
well
(see
also
Di
Matteofdirectionof escaping
photonsaccretio
and theYu &from
Tremaine
2002)
while
the
mass
accretion
rate
dependson
Requirement
be2005).
constrainedby observations.
BS09
find
the
models
withwords,Progaet al. 1998;
Proga &
Kal
et
al.
In
the
study
of
BS09,
they
find
that
the
global
BH
the radiation
is emitted
mostl
feedback MBH-σ
rel. constrained
poorlyconditions
andgeometry
atwhereaslarge& Schaye
Dubois+’12 Magorrian!f =density
0.15 physicalmatches
observations(e.g.,
therelation
numberat Booththe current
era’09;(zero
redshift)andthe BHis oEddington
solutions,Ṁoutmainl
the inflowing
gas occurs
distances.
and
the
M
–σ
relations)
very
well,
and
similarly
SDH05
find
scaling
relationsareverysensitive
toachoice
of
!
,
and
the
BH
weakercorrelationismainly
ca
f
3
As
it
has
been
discussedandshodistances fromthe
BH.
✏f ⇠!10= 0.05Bellovary+’10
thatmatches
observations
well(seeofalso
former is nearly
inversely proportional to thevalue!f . Di Matteothat self-consistentlyfdirection of escaping
photons
determined
prehDi
Matteo+’05
A commonto study
estimateṀtheytothatadopt
the BHanalytic
a isfindet al.method
2005). In theof BS09,
the global
the canradiationis emitted
quasarwords,
radiation
significantly
reletters to nature

AGN feedback ef ciency

more abrupt in the more massive systems. Here, the total gas supply
for accretion is larger, and the gravitational potential well is deeper,and so the black hole has to grow much more before its released
energy is sufficient to expel the gas in a quasar driven wind, whichthen terminates further nuclear accretion and star formation. For
the same reasons, the initial growth of the black holes, which isregulated by the properties of nearby gas, depends on the total mass.
It is faster in more massive systems, which can therefore reach theexponential, Eddington-limited growth phase more easily. The
lifetime of the active black hole phase, however, increases for smaller
black hole masses, implying that low-luminosity quasars should be
more numerous than bright ones. This is consistent with them
residing in a greater number of smaller galaxies and with what has
been found in recent surveys26,27.
The dependence of black hole growth on galaxy mass yields a
relation between the stellar spheroid of the remnant galaxy and its
central black hole. Figure 3 shows the black hole mass versus thestellar velocity dispersion of the merger remnants from our simulations, compared with observations. We show simulations with six
different galaxy masses, each of which has been run with threedifferent initial gas mass fractions of the galaxies’ disks. Remarkably,

ou
we
wh
(w
th
of
M
sc
fra
ga
m
ho
in

m
en
ho
bu
ge
fo
co
ta
lo
M
sim
bl
lif
ro
fe
co
ex

ho
sta
tim
co
qu
ap
hi
U
qu
th
ea

Re

fi

 

fi

fi

1.

Figure 3 The relation between the final black hole mass, M BH, and the velocity dispersionof stars, j, of our galaxy merger simulations compared with observational measurements.
Filled circles show the masses of the black holes and the bulge velocity dispersions

2.
3.

AGN feedback energy
Radiative luminosity:
Feedback Energy:··2
Ef = Ef dt = ϵf ϵrc Maccdt∫∫

E.g., growth to 108 M⦿ SMBH:(cumulative feedback energy)

Bulge potential energy:So, ϵf > 0.01 can give Ef > Ebulge

Mbulge
MBH

−3

∼ 10

Ferrarese & Meritt ’00

AGN jet
電波観測~1.5 kpc

NASA

9

MBH ∼ 6.5 × 10 M⊙

Chandra + VLA(blue) (red)EHThttps://www.sciencesource.com/archive/Supergiant-Elliptical-Galaxy--M87--NGC-4486-SS2712619.html

IllustrisTNG

AREPO code — voronoi tessellation

Two-mode AGN feedback model
Eddington-limited accretion:

Eddington ratio threshold

high
low

thermal (quasar) mode

kinetic (jet) mode

(maintenance mode)

weaker coupling in
low-ρ environment

Weinberger+’18

MBH — σ relation
Illustris (Mseed ~ 105 M⦿/h)

TNG (Mseed ~ 106 M⦿ )Too small σ in simulations?
Or, too massive seed + too ef cient growth?
fi(at low-mass end)

Li+’19

z=4

z=2

z=0

Habouzit+21

Illustris

TNG100TNG300

Horizon-AGN

EAGLE

SIMBA

Habouzit+21

FOREVER22 projectz=3 10 cMpc PCR regions

Yajima+21

High-res. jet models

Martizzi+19
AthenaMukherjee, Wagner+18
Jet-ISM interaction

Bourne & Sijacki ’17
AREPO~200 pc

cf. 大須賀さんtalkCR feedback(Cosmic Ray)CR feedback• non-thermal component — comparable E-density to Eth, Emag, Eturb• transport mechanism, interaction with EM waves — coupling with thermal plasma(e.g. Zweibel ’13,17; Amato & Blasi’18)• Sources of scattering? — streaming instability, background turbulence—> CR m.f.p.

Kulsrud & Pearce’69

Yan & Lazarian ‘02• Con ne, and isotropize the CR distribution• Can provide signi cant pressure in CGMs

• Earlier work: e ective di usion, grey, steady approximations
Analytic: Ipavich ’75; Boulares & Cox 90; Breitschwerdt+91; Everett+08; Socrates+08; Mao & Ostriker’18; ….

ff

fi

ff

fi

Simulation: Booth+13; Pakmor+16; Ruszkowski+17; Wiener+17; Chan+19; Hopkins+20; Ji+20; Su+20; Hopkins+21; ….

FIRE-2 MHD sim. + CR

— w. di usion, streamingin dwarf, L* galaxies
CR energy ux

advection & streaming
pure di usion (zeroth moment)(—> later expanded to “two moment”)

κ : e ective di usion coe cientcooling

Lagrangian, nite-volume form:

CR loss

adiabatic term

: total pressure (thermal + magnetic + CR)
CR pressure( > 0 ) Streaming loss

conserved total CR energy

ffi

ff

fi

ff

fl

ff

ff

Chan+19

CR energy injection:
default)

Hadronic & Coulomb losses:

Rather than setting

2-moment method: explicitly solve(Volk+96; Ensslin+97; Guo & Oh ’08)

Volumetric gas heating:Composite parallel (B- eld aligned) di usion coe :(Snodin+06; Jiang & Oh ’08; Thomas & Pfrommer ’19)

Update CR ux as

ff

ff

fi

fl

Chan+19

Chan+19
Hopkins+20
Isotropic di usion coe .

is favored against -ray obs.

ff

𝜸

ff

(cf. Quartaert+21)(cf. Lacki+11; … )FIRE-2 MHD cosmo sim. + CR : follow-up to Chan+19

Hopkins+21

Slow out ow produced by CR
No CR(but + rad, conduction, MHD, stellar FB)With CR

Slow out ow accelerated
in-situ by CR

fl

fl

FIRE-2 sim.

Hopkins+21

FIRE-2 sim.

Hopkins+21

Chan+21

Temperature structure above disk plane
Hot infall (T ≳ 105 K, ~50 km/s)
w.out CR

Warm slow (T ≳ 104 K, <200 km/s)
CR-driven out ow- weaker adiabatic losses- do not su er from rad. cooling- PCR slows down the infall

flff

Chan+21- lowers post-shock T

MHD with CRs

Pressure
Energydensity(w.out ϵcr)

Streaming vel.

EOS:
KE-weighted spatial
di usion coe cient(in rela. limit)(Enslin+07; Kulsrud & Pearce’69; Wiener+13,16)
Heaviside-Lorentz system of units.

ffi

ff

AREPO; Pfrommer+17

Spectrally resolved CR
AREPO MHD sim.

Fokker-Planck eq.

Number & E density- accurate CR cooling- E-dep. spatial di usion

CR pressure- grey CR- di usive transport- e ective cooling
Piecewise power-law ptcl distrib:

Girichidis+20, 21

ff

ff

ff

cf. Miniati+01; Yang & Ruszkowski’17; Ogrodnik+21

Spectrally resolvedCR
initial out ow front by
CRs w. 200 - 600 Gev/c

subsequent out ow by~10 Gev/c

flfl

Girichidis+21

Temporally & Spatially varying CR spectrum

Girichidis+21

Summary & Issues• Multi-scale, multi-phase structures in the universe.
• SNe, AGN, CR feedback: all seems relevant for galactic wind
• SN FB model: rapidly improving at ≲pc level — clustered SNe• AGN FB model : still highly uncertain; intermittent / two phase growth of SMBH?
• CR FB model: details of propagation model — di usion coe cient, scattering,

ffi

ff

subgrid turbulence, spectral decomposition

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林航平（一関高専）、東野聡（神戸大学）、身内賢太朗（神戸大）、長尾桂子（岡山理科大）、中竜大（東邦大）

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