Neutrino Emission from Accretion Flowsin Active Galactic Nuclei
Tohoku University

Shigeo S. Kimura
References:
Murase, SSK, Meszaros, 2020, PRL, 125, 011101
Kheirandish, Murase, SSK, 2021, ApJ, 922, 45
SSK, Murase, Meszaros, 2021, Nat. Comm., 12, 5615

ICRC satellite workshop@Chiba Univ.
Aug. 7, 2023

Index• Introducon• Neutrino emission models in Seyfert galaxies• Neutrino emission from AGN coronae• Sub-GeV gamma rays from Seyfert galaxiesti• Summary2Index• Introducon• Neutrino emission models in Seyfert galaxies• Neutrino emission from AGN coronae• Sub-GeV gamma rays from Seyfert galaxiesti• Summary3Cosmic Neutrino Background Spectrum
A combined fit of IceCube’s high energy neutrino data

4

IceCube Collaboration ICRC 2023

Figure 1: Result of the combined
fit of tracks and cascades (in red) under the assumption of an astrophysical SPL neutrino flux. Previous results from measurements using single event channels are shown forcomparison. Note that the sensitive
neutrino energy ranges (as indicatedin the upper panel) and neutrino flavor probed are diﬀerent amongst thediﬀerent samples.

Soft cosmic neutrino spectra→ Medium energy excess (High[7]intensity
@~10
TeV)
R. Naab, “TeVPA 2022,• Origins of cosmic neutrinos are a https://indico.cern.ch/event/1
82486/contributions
new big mystery
•

•

γ

•
νμ

νμ

TargetμπTargetνγπγ

Cosmic ray

νμ

γ

νμμ
ν
π

ν

νμ

p+p → p+p+π• π±→ 3ν+e
• π0→2γγγ
γγπCosmic ray

μ

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p+γ → p+π
• π±→ 3ν+e
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Interac on between CRs & photons/nuclei → Neutrino produc on
Gamma-rays inevitably accompanied with neutrinos
h

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High-energy neutrino production

＜
〜
〜

p

5

Gamma-ray Constraint on Neutrino Sources
Fermi Satellite is measuringcosmic gamma-ray backgrounds••

ν ux@10 TeV > γ-ray ux@100 GeV•Consider sources from which
both γ & ν can easily escape→ t theory to neutrino data
→ γ-ray theory >> γ-ray data

Neutrino theory

γ-ray data

γ-ray needs to be absorbed
inside the sources (hidden source)•+

γ+γ→e +e

−

fl

fi

ffi

X-rays e ciently absorbs GeV γ-rays•

flγ-ray theory

Bechtol et al. 2017
Murase et al. 2013, 2016

Neutrino data6

we focus on scenarios which may produ
7
Under the current constraints imposed byanalyses mentioned above, such LL GRBs a
the originators of the diffuse VHE neutrino f• Choked GRBs

Hidden Neutrino Source Candidates• AGN Core
Kimura et al. 2021

Orphan Neutrinos

Senno et al. 2016

Stall RadiusPr

Choked Jet

ExtendedMaterial
CE•

Most luminous steady source in the Universe•

Source of Cosmic X-ray background•+

γ+γ→e +e−

Progenitor
Core

GRBs failed to penetrate stellar envelope
FIG. 1. Left panel: The choked jet model for je• Stellar envelope absorbs γ-rays
jet is hidden, and such objects may be observ+−
transrelativistic
shocksare• p + γ → p + e + e driven by choked jets.
•

we focus on scenarios which may produ
8
Under the current constraints imposed byanalyses mentioned above, such LL GRBs a
the originators of the diffuse VHE neutrino f• Choked GRBs

Hidden Neutrino Source Candidates• AGN Core
Kimura et al. 2021

Orphan Neutrinos

See Murase-san’s talkPr

Senno et al. 2016

Stall Radius

Choked Jet

ExtendedMaterial
CE•

Most luminous steady source in the Universe•

Source of Cosmic X-ray background•+

γ+γ→e +e−

Progenitor
Core

GRBs failed to penetrate stellar envelope
FIG. 1. Left panel: The choked jet model for je• Stellar envelope absorbs γ-rays
jet is hidden, and such objects may be observ+−
transrelativistic
shocksare• p + γ → p + e + e driven by choked jets.
•mission from NGC 1068

IceCube (this work)

Electromagnetic observations (26)

Evidence of Neutrinos from Seyferts

9

Downloaded from https://www.science.org at Tohoku UniverE 2 Φ [TeV cm− 2 s− 1 ]esolution scan around the most sigTheoretical ν model (52,55)
0.1 to 100 GeV gamma-rays (40,4cation in the Northern Hemisphere
Theoretical
ν
model
(53)
>
200
GeV
gamma-rays(42)n Fig. 2A, with NGC 1068 located−968% confidence region. The posi10
IceCube 2022
þ22
^ ns ¼ 79#20 more
C 1068 produced mn expected from the atmospheric−1010
Neutrino
astrophysical neutrino backgrounds.
X-ray
shows the distribution of the anguion of these events from NGC 1068.
10 −11
Op cale 79 most contributing events, 63
γ-ray
ded in a previous analysis (23). The−1210^ ns is ~2 events (26).
uncertainty on m
þ0:2^
ured spectral index is g ¼ 3:2#0:2
Radio
−13
Y. Inoue et al. 202010
timated systematic uncertainty of, consistent with previous results
Murase, SSK, Meszaros 2020−14
stimate these systematic uncertain10
M77 (NGC 1068)
lyzing simulated data, assuming a
10 −1510 −12
10 −9
10 −6
10 −3
10 0
10 3
10 6
h flux equal to the one measured forEnergy[GeV]
but varying assumptions about the• Point
sponse(26). Systematic
uncertainties
sourcesearchwith
10-year
data
set
with
an
improved
analysis
methodcant location. (A) High-resolution scan around the Fig. 4. Multimessenger spectral energy distribution of NGC 1068. Gray points show multifrequency
ycontours
fromthemodelingofthephotonshowingits
68%(solid)and(dashed)result:• Cataloged source95%
observations(data
sources—>listed4.2σ
in tableS1). Dark and light green points indicate gamma-ray observations
search2.9σ(2020)(2022)
n in1068,
the andglacialNGC
the redice—e.g.,
circle is itsscattering
angular size in
at 0.1 to 100 GeV (40, 41) and >200 GeV (42), respectively. Arrows indicate upper limits, and error bars
2
quareddistance,, between
NGC
1068
and
tion—andefficiency
with
which
Fν the≫Fγ y^—>
Hiddenneutrino
source• angular
are 1s confidence intervals. The solid, dark blue line shows our best-fitting neutrino spectrum with theackground(orange)andthesignal(blue)
fromre detected by the IceCube opticaldark blue shaded region indicating the 95% confidence region. We restrict this spectrumto (NGC
the rangetrum•at theposition
of
NGC
1068.
The
superposition
M77
1068)
γ-ray,
CR&ν
produc
on
sites
are
under
debates.
Let’s
discuss
possibilies.
ystematic
uncertainties
aresmallerlack. This representation of the result ignores the
ti

ti

ti

between 1.5 and 15 TeV, where the flux measurement is well constrained (26). Two theoretical predictions

Index• Introducon• Neutrino emission models in Seyfert galaxies• Neutrino emission from AGN coronae• Sub-GeV gamma rays from Seyfert galaxiesti• Summary

10Neutrino & gamma-ray production sites• Large emission regions6≳ 10 RG• Compact emission regions11
2≲ 10 RG••

Mini jets∼10 - 100 pc
Michiyama et al. 2022

ti

•

−5

∼ 10

Starburst nuclei∼ 0.1 - 1 kpc
Loeb & Waxman 2006
Murase et al. 2013

AGN winds−3− 10 pc

Inoue S. et al. 2022•

Accre on shocks
−5

∼ 10
•

−4

− 10 pc

Inoue Y. et al. 2021

AGN coronae−5∼ 10 pc
Murase, SSK, Meszaros 2020
Kheirandish, Murase, SSK 2021
Eichmann et al. 2022

Starburst nuclei

Yoast-Hull et al. 2014
Wang et al. 2018•

Cosmic-ray produc on@SNR• Neutrino produc on@ISM

Mini jets Michiyama et al. 2022
Cosmic-ray produc on@jets• Neutrino produc on@jets or ISM

•

•

ν

ti

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γ

ti

•

ti

ppppp

Large emission regions

ν

γ

12

Starburst nuclei

Yoast-Hull et al. 2014
Wang et al. 2018•

Cosmic-ray produc on@ SNR• Neutrino produc on@ ISM

Mini jets Michiyama et al. 2022
Cosmic-ray produc on in jets
• Neutrino produc on in jets or ISM••

Hadronuclear interac on —> γ-rays consistent with GeV γ-ray data
ν

ti

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γ

ti

•

ti

ppppp

Large emission regions

ν

γ

13

γ-ray constraints
4
NGC 1068 should be hidden sources—> demands compact emission sites• EM cascade modeling with γ-ray data:—> Emission region: R ≲ 100RS

10- starburst nuclei (R ≳ 10 RG)
5- radio jets (R ≳ 10 RG)
as neutrino emission sites

30 RS (γ)
300 RS (γ)
410 RS (γ)

minimal pp

-7

10
-2 -1

E FE [GeV cm s ]6

Murase 2022

-6•

• This constraint rules out

14

ν

IceCube

-8

10

-9

10

Fermi

-10

MAGIC

10

-11

10

-3

10

-2

10

-1

10

10

0

1

2

10 10
E [GeV]

10

3

4

10

10

5

10

6

Neutrino & gamma-ray production sites
15• Large scale emission regions

❌

•

•

Mini jets∼10 - 100 pc

ti

Michiyama et al. 2022

AGN winds−5∼ 10

Starburst nuclei∼kpc
Loeb & Waxman 2006
Murase et al. 2013•

• Compact emission regions−3− 10 pc

Inoue S. et al. 2022•

Accre on shocks
−5

∼ 10
•

−4

− 10 pc

Inoue Y. et al. 2021

AGN coronae−5∼ 10 pc
Murase, SSK, Meszaros 2020
Kheirandish, Murase, SSK 2021
Eichmann et al. 2022

AGN Accre on Flows

10 cm

46

1 cm

1 mm

100 µm 10 µm

1 µm

1000 Å 0.1 keV 1 keV

10 keV

SED: Blue bump & strong X-rays→ Op cally thick disk + Corona

QSO

Wind?

. Astrophys. 2008.46:475-539. Downloaded from www.annualreviews.org
by Osaka University on 12/18/13. For personal use only.

•

Coronae?

ti

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O

Accre

log (νLν /ergs s–1)

Ho 2008

S

Q

16

44

42
log Lbol /LEdd ≥ 0–1 ≤ log Lbol /LEdd < 040
PHYSICALREVIEW LETTERS
125,
0111
–3 ≤ log L /L < –1
bol

Edd

log Lbol /LEdd < –3
38

9

10

11

12

13

14

15

16

17

18

log (ν/Hz)
Figure 7

Composite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm
(Adapted from L.C. Ho, in preparation.)

nuclei (Ho 1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho &2001). Defining radio-loudness based on the relative strength of the radio and X-ray emi
RX ≡ ν Lν (5 GHz)/LX , Terashima & Wilson (2003b) also find that LINERs tend to be r
loud, here taken to be RX > 10−4.5 . Moreover, the degree of radio-loudness scales inversely
Lbol /LEdd (Ho 2002a; Terashima & Wilson 2003b; Wang, Luo & Ho 2004; Greene, Ho & Ulv
2006; Panessa et al. 2007; Sikora, Stawarz & Lasota 2007; L.C. Ho, in preparation; see Figure
on Shock?

42

νµ /dεν∼ 3 × 10 erg/s in this range.
er limits for gamma rays above 0.2 TeV
dels in which TeV gamma rays and neue source with similar flux [42].
reInoue S.Some
et al. 2022
• Two-zone model
nvoke proton acceleration and neutrino
t coronal
regions
near (theBH whereX≲ 100R
region)• InnerG
via thermal→ Comptonization,
ν produc on & γ aeither
enuaac-on
nae [43, 44] or accretion shocks [45, 46],•
Outer
region
(
pc)
≳
0.1
4
nying gamma rays would be significantly→ γwithproducinteractions
AGNon
photons [10].
roximating
Lp ∼ Lpicture/Ep,min), top ln(Ep,maxose
analternative
where
protons• Merit:
42
FIG.
1.
Schematic
sketchofthemodel.
The
accretion
disk
(ε
=
1
TeV)
∼
10
erg/s
as
observed,
ν
n the
inner
regionsoftheonewindrelatively- Explainγ&νbyscenarioaround
the
black
hole
(BH)
drivesanoutflowing
wind.
In6
44
Lνµ ln(10) which∼ 10 has
erg/s,
which
is someNGC
1068,
variousadvantages
ner region: winds from the inner disk dissipate their kinetic
Outer
wind
can
explainγ-rayscregionbut withina plausible
range,
discussed
a aswell
estabmodels[10]. DSA,
energy via shocks near the BH, caused by failed line-driven
withnaturalparametersets
fixed, εν,maxisisassumed.
plausibly
mtherforparameters
particle acceleration,
winds that fall back. Protons undergo diffusive shock accelerInner
Inner
"be
=
andwithηg =1.
Somewhat
lower
• 0.5
Demerit:
B identified
ation (DSA) and pγ interactions with photons from the diska “failed” wind that isherηg radiative,
give
lower toεν,maxthat wind
may bemore
Outer cascadeand corona, inducing neutrino and electromagnetic- neededuse unnaturalvalues:ed inline-driven
modelsththecurrent
IceCubedata[44].
emission,
modulatedbyγγinteractions.
Outerregion:suc8
10s corresponding
to
NGC
1068
[47].
Neuvw ∼ 10 cm/s ≪ vesc ∼ 10 cm/s cessful winds from theOuterouter disk propagate farther, partially
depth
for
γγ
interactions
withcoronaphoy generated via pγ interactions with the−1
impact
the
torus
and
triggershocks.
Protons
undergoDSA—>
slow
acceleraonv∼v
so gammarays" 410γγ(ε/1GeV)R̄1 [10],
w
esc
while
interactions
mediatethe associand pp interactions with the torus gas, inducing gamma-ray
neutrinos
are
above
a
few
4 attenuated2
eithemission,which
weevaluate
acrosstheηacc ≳ 10 : tacc = ηacc(rL /c)(c/vsh) emission, affected by γγ interactions with photons from theo theregionmodels.
m.
Forcoronal
the GeVgamma
rays, we invoke atorus. Indicated scales are only approximate.
ti

ti

tt

ti

ti

AGN wind scenario17

Accretion Shock Scenario

Galaxies 2021, 1, 0

CR produc on @ accre on shock• Neutrino produc on using AGN photons• Consider weak B- eld (B∼10-100 G)•

10 8CoronaDisk BH

E 2dN/dE [erg/cm2/s]

ν

ti

ff

fi

ti

ti

Shock

UV

Disk

GRAMS
(35 days)

10 9p+γ→p+π
π±→ 3ν+e±π0→2γX: Corona: Corona (Screen)

Inoue Y. 2020

Cf. Stecker et al. 1991

ti

p

18

γ

: Corona (Uniform)
IceCube
4FGL
3FHL
MAGIC

10 10

10 11
AMEGO
(5 yrs)

10 12
GRAMS
(3 yrs)

10 13

10 14

10 15
106

107

108

109

1010

1011

1012

1013

1014

1015

1016

Energy [eV]

Figure 11. The gamma-ray and neutrino
spectrumto
of reproduce
NGC 1068. The• Possible
ν circle,
datasquare, andR∼The
10-30Rgpoints are from
Fermi-LAT
collaboration [135], Ajello et al. [136], and Acciari et al. [16without
overshoongγ
data
The green shaded regions represent the 1, 2, and 3s regions on the spectrum measured b• γ-rayscome
di erent
regionThe expected gamma-ray and neutrino
spectrum
fromfrom
the corona
are shownfor 30 
The darker region corresponds to lower
hg . starburst
The blue region
shows the expected neutr
e.g.,
or
jets
See also Inoue Y. et al. 2019
The orange and magenta shaded region shows the gamma-ray spectrum for the uniformDemerit of Accretion Shock Scenario
10-2

Narayan et al. 2012100

SSK et al. 2014

100

10-1

-2

10

ν
-3

10

Downloaded from http://mnras.oxfordjournals.org/ at Osaka University on January 16, 2015

Existence of shock is unclear(Magneto-)hydrodynamic simula ons
do not nd any shock structure

•

10

19

10

100

Figure 3. Radial distributions of (a) the surface density Σ, (b) the specific angular
Figure 8. Similar to Fig. 7, but for time chunks M2 (top left), M3 (top right), M4SSK(bottomet
left),(bottom right) of the ADAF/MAD
Note that and dotted lines are for B2 (fv = 0.9), B1 (fv = 0.3), a
al.M52019
B. simulation.
The solid, dot–dashed,in chunk M5 (lower right) rstrict and rloose both lie outside the plotted area (see the numerical values given in Table 2).
Keplerian angular momentum. (d) Radial distributions of the integrated pressure fline depicts PTP for A1 (no HEPs) for reference.

Weak B elds in accre on shock scenario- accre on phenomena driven by B elds
- MHD Weinstabili
esamplifyB
eldsthen define two criteria, one ‘strict’ and one ‘loose’, to estimate
the 3
radius range over4 which the flow has achieved inflow equilibrium:B ∼ 10 − 10 G

tvisc (rstrict ) = tchunk /2 = ttot /4,(9)

tvisc (rloose ) = tchunk = ttot /2.

(10)
fi

ti

fi

ti

ti

fi

fi

Balbus & Hawley 1991, 1998

ti•

profile v r (r) of the gas within one scale height of the mid-plane(the restriction to one scale height is to enable us to focus on theaccretion flow rather than any mass outflow or jet). From this, weestimate the viscous time as a function of radius r in the standard
way:
r
.
(8)
tvisc (r) ≡|vr (r)|fluctuations. The strict criterion has ttot = 2tchunk = 4tvisc , which
For D3 (fv = fc = 0.9), the injection from the viscousis a fairly safe and conservative choice, while the loose criterion
dissipation mainly energizes HEPs because PTP is so small thattakes a more optimistic view of how soon inflow equilibrium isthe injection from the compressional heating is inefficient. On
achieved. Note that Penna et al. (2010) defined inflow equilibrium
the other hand, the injection from the compressional heating isby the condition ttot = 2tvisc , which is the same as our present loose
for D1for the varioustime(fv = fc = 0.3) because PTP is large enoughcriterion. The values of tchunk , rstrict and rloosedominant
Qvis< QV,TP . Both the viscous dissipation and theandsatisfy
rloose are
shown
chunks are listed in Tables 1 and 2, and rstrictto
make nearly the same contribution to theas circular solid lines in Figs 7 and 8. It willcompressional
be noticed thatheatingtheinjectionwith
for values
D2 (fv = fc = 0.6).
objectively determined rstrict and rloose are compatible
one might deduce by visual inspection of Fig. 6.
3.3.
Luminosities of Escaping Particles
In Figs 7 and 8, the time-averaged velocity streamlines are
wellbehaved within the respective inflow equilibrium regions of the four
We also calculate luminosities of escaping gamma rays,
panels. Note also that the steady-state zone is much more extended
neutrinos, neutrons, and protons. We define the luminosities
in the MAD simulation compared to the SANE simulation. For

Neutrino & gamma-ray production sites
20• Large scale emission regions❌

•

Starburst nuclei∼kpc
Loeb & Waxman 2006
Murase et al. 2013

Mini jets∼10 - 100 pc

Michiyama et al. 2022

ti•

• Compact emission regions△•

AGN winds
−5

∼ 10

−3

− 10 pc

Inoue S. et al. 2022•

Accre on shocks
−5

∼ 10

•

−4

− 10 pc

Inoue Y. et al. 2021

AGN coronae−5∼ 10 pc
Murase, SSK, Meszaros 2020
Kheirandish, Murase, SSK 2021
Eichmann et al. 2022

Index• Introducon• Neutrino emission models in Seyfert galaxies• Neutrino emission from AGN coronae• Sub-GeV gamma rays from Seyfert galaxiesti• Summary21

AGN Accre on Flows

10 cm

46

1 cm

1 mm

100 µm 10 µm

1 µm

1000 Å 0.1 keV 1 keV

10 keV

SED: Blue bump & strong X-rays→ Op cally thick disk + Corona. Astrophys. 2008.46:475-539. Downloaded from www.annualreviews.org
by Osaka University on 12/18/13. For personal use only.

•

QSO

ti

ti

O

Coronaelog (νLν /ergs s–1)

Ho 2008

S

Q

22

44

42
log Lbol /LEdd ≥ 0–1 ≤ log Lbol /LEdd < 040
PHYSICALREVIEW LETTERS
125,
0111
–3 ≤ log L /L < –1
bol

Edd

log Lbol /LEdd < –3
38

9

10

11

12

13

14

15

16

17

18

log (ν/Hz)
Figure 7

Composite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm
(Adapted from L.C. Ho, in preparation.)

nuclei (Ho 1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho &2001). Defining radio-loudness based on the relative strength of the radio and X-ray emi
RX ≡ ν Lν (5 GHz)/LX , Terashima & Wilson (2003b) also find that LINERs tend to be r
loud, here taken to be RX > 10−4.5 . Moreover, the degree of radio-loudness scales inversely
Lbol /LEdd (Ho 2002a; Terashima & Wilson 2003b; Wang, Luo & Ho 2004; Greene, Ho & Ulv
2006; Panessa et al. 2007; Sikora, Stawarz & Lasota 2007; L.C. Ho, in preparation; see Figureラックホールに引きつけられたガスは, 回転しながらブラックホールに

Par cle Accelera on in Turbulence

23

込まれる (降着円盤 or 降着流).

スの重力エネルギーが変換され, 明るく輝き, 場合によってはジェット

出する.

巨大ブラックホール
Particle-In-Cell Simulations in shearingbox
Particle-In-Cell Simulations with turbulence(太陽質量の１億倍)
エネルギー天体のエネルギー源
Hoshino 2013, 2015; Riquelme et al. 2012; Kuntz et al. 2016

Comisso &
2018,
2019;
Zhdankin
etDecember
al. 2018
TheSironi
AstrophysicalJournal,
886:122(29pp), 2019
1

The Astrophysical Journal, 886:122 (29pp), 2019 December 1

Comisso & Sironi

ジェット(赤)座X-1の想像図

降着円盤
X線で光る)銀河(白)ブラックホール(太陽質量の10倍程度)(可視光で光る)電波銀河の観測

The Astrophysical Journal, 886:122 (29pp), 2019 December 1

Comisso & S

Stochastic
turbulence
FIG. 1 (color online). Time evolution of the magnetorotational instability. Panels (a) and (b) show the magnetic field accelerationlines (greenish

水曜日

lines) and angular velocities in the background at Y ¼ y=λ ¼ 1.91 and X ¼ x=λ ¼ 1.91 (color contour), and panels (c)–(e) depict thehigh-density regions as reddish curved planes. Panels (b) and (c) are at the same time stage. Panel (f): The energy spectra during the MRI
Non-thermal
tailat T orbit ¼ 0.31,
6.89, 7.18, 8.84, and14.28. The dashed line is a Maxwellian fitting for T orbit ¼ 7.18.

Power-law tail061101-2

Figure 2. 3D plots of different ﬂuid structures in fully developed 3D turbulence (at ct/l=2.7) with σ0=10, δBrms0/B0=1, and L/de0=820 (with l=L/4). The
Figure6. leftTime
evolutionof thedN/dln(γ−1)forof thedensityJz along themean magnetic
displayed quantities
are (fromto right,
top to bottom)
the particle
ﬂuctuationspectrum
magnetic energy
density in units
B20/8π, the current
Figure7. Particle
spectradN/dln(γ−1) at late times for simulationsin
Figure2.
At
latetimes,
thespectrumdisplaysapower-law
tail
c,
the
bulk
dimensionlessfour-velocity
Γβ,
and
the
particledensity
ratio
n/n
.
Note
that
the
color
bars
for
Γβ
and
n/n
are
in
logarithmic
scale.
ﬁeld in units of ensimulation000magnetization σ0=40, systemAn size L/de0=3280 (with l=L/8),
different
at https://doi.org/10.7916/d8-prt9-kn88.
animation showing
the current
Jzdinlogp
N d x-ylogslices(H can1) be_ found2.9. About
16% of the particles
withindex density

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Magne
reconnec
on → rela vis c par cle produc on(a) and (b) show the magnetic
field clines
(greenish
1.91 (color contour), and panels (c)–(e) depict theInteracon
with
Turbulence→
furtherenergizaonstage. Panel (f): The energy spectra during the MRI

different values of initial ﬂuctuations EBrms0 B0  {1, 2, 4}. For thehave γ 15 at ct/l=12 (twice the peak of the particle energy spectrum),with larger initial ﬂuctuations, the late-time particle spectrum displays a po
which gives an indication of the percentage of nonthermal particles. The insetpd0.6
log. AtN d log (H  1) _ 1.9 , and about 31% olawattailwithH
1 _index
Hth0 
1
Note also that showsin the 3D
case the magnetic
shown
byatheblue linefunctionof peaking
the power-law
index penergy
and thedecays
cutoff faster
Lorentz factor
γc as
γhas
25decayed,
at ct/l=12
1 when2lateHtimes,
most of the particles
turbulent haveenergythe (twice the peak of the particle enthan in the 2D thecasemagnetization(compare insets
Figuresand indicates4). We willThedashed3 line
the scaling
expected
σ0. ofc r T0spectrumatthattime),
which
spectrum stops evolving (orange and red lines): it peaks atgives an indication of the percentag
show that this forleads
a reduced particle
acceleration
at latea σto0-independent
domainsize L/de0rate
=820.
nonthermal
particles.
γ−1∼5and extends
well beyond
the peakinto a nonthermal
Figure11. Relation
betweenparticle
injectionand electric current density fromtimes.
tail of ultrarelativistic particles that can be described by a power
the 3Dwith σ0=10, δBrms0/B0=1, and L/de0=820. Top panel:
lawsimulation
power-law index p for increasing magnetization
σ0 (see also
γ (bottom inset in Figure 6), deﬁned as the Lorentz factor wh

CR Acceleration Theorye.g.) Fermi 1949

MHD + Test Particle Simulations
SSK+ 2016 ApJ, 2019 MNRAS; Sun & Bai 2021

＜〜〜〜
’

’

MRI
turbulence

〜〜〜＜
’

’

where the integration region is set to be 0.1Rc ≤ R ≤ 0.6Rc andFigure 6. Orbits of test particles projected to the R − θ plane (upper panel)|cos θ| ≤ 0.45. We plot the power spectra, mPm , for the magnetic
and the R − φ plane (lower panel) for λini = 4. The initial and final positionsof the particles are shown by the stars and circles, respectively. In the bottom
panel, the cyan circle and black arrows indicate the initial ring R = Rini andthe rotation direction, respectively.

∂F

∂F
1 ∂
=
E D
∂t
E ∂E (
∂E )
ti

ti

e

e

l

l

ff

E

c

E

c

e

i

e

i

t

t

<

v

>

v

r

r

a

a

where eφ is the unit vector of the φ direction and Vbul, φ is independent
p of θ. The bottom panel2 shows the pmomentum distribution
in the fluid frame, where we can see no bulk rotational motion. InE2
the following sections, we use the energy distribution in the fluid
frame. Note that the particle distribution is slightly anisotropic: thealab-frame, while the bulk motion is not seen in the fluid frame.

where m = kφ R (kφ is the wavenumber in the φ direction). Then,
we take the average of the power spectrum over the disc region:"|Xm |2 RdRdθ
"
,
(14)
Pm =
RdRdθ

Some gain E, others lose E→di usion in E space

a

E

w

p

E

EE

w

p

EE

Stochas c Accelera on by MHD Turbulence

Downloaded from https://academic.oup.com/mnras/article-abstract/485/1/163/5305870 by guest on 29

frozen in the differentially rotating fluid elements that fall to the
BH. This creates the spiral structure as seen in the figure. We canalso see that the fluctuation of the density is much smaller thanthat of the magnetic field energy density. This implies that the fastmodes are a sub-dominant component in the MRI turbulence.
To clarify the importance of the modes of the MHD waves (fast,slow, and Alfven), we evaluate the Pearson correlation coefficients
between the fluctuations of the density, δρ(R, θ, φ) = ρ − !ρ"L ,and the magnetic energy, δB 2 (R, θ, φ) = B 2 − !B 2 "L . According
to the linear MHD wave theory, the fast mode has a positive
correlation, the slow mode has a negative correlation, and the Alfven
mode has no correlation. We evaluate the correlation coefficientsas a function of R and θ, and average over them with weights
associated with the area in the meridional plane. The resulting
coefficients indicate that the density and magnetic energy are weakly
anticorrelated: the value of the coefficient is −0.22 in the discregion (|cos θ ! 0.45|) for run A. The lower resolution runs havehigher coefficients, i.e. the anticorrelations are weaker, but no run
has a positive correlation. Therefore, the fast modes do not play
an important role in this system. This result is natural in the subAlfvenic and sub-sonic turbulence.
Finally, we discuss the azimuthal power spectra of the turbulence(cf. Sorathia et al. 2012; Suzuki & Inutsuka 2014; see Parkin &Bicknell 2013 for three-dimensional power spectra). We take theFourier transformation in the azimuthal direction,! Figure 7. Momentum distributions at t = 10tL in the lab frame (upper)
1X
exp(−imφ)dφ,
Xm = √
andthe fluid flame (lower) for λini(13)= 4. We can see a bulk motion in the
2π

24

Figure 4. Colormaps in the equatorial plane for run A. The upper and low
panels show the density and the magnetic energy density, respectively.

Stochastic
acceleration

MNRAS 485, 163–178 (201PHYSICALLETTER
MULTIMESSENGER TESTSREVIEWOF COSMIC-RAY ACCELERATION…

(1)✓ ◆2 ⇣ ⌘2
q2
where ⇣c
Fp is
the
CR
distributionfunction,
D
≈
ε
VA
rL
εp
p /tacc2
D "p ⇡
"p ,
(9)
−1
H
c
H
is the diffusion coefficient in energy space, tp−cool = t−1
pp +
−1
−1
−1
−1
−1
−1
tpγ +tBH +tp−syn is the total cooling rate, tesc = tfall +tdiff˙E  [GeV cm s sr ]QSO
here Fp turbulence.
is the momentum
distributionfunction
We compute steady state CR spectra by solv2
dN/d"p = 4⇡p Fp /c), D"p is the di↵usion coefficient,ing
thefollowing
Fokker-Planck
equation(e.g.,[75–78]),
•
Equa
ons
for
cosmic-ray
protons
ool is the cooling time, tesc is the escape time, and! stochastic acceleration. "F˙p,inj is the injection term to the3ε1 ∂∂Fp waves,pthe
Fp∂Fp2
onsidering resonant
scatterings
with
Alfvenεp Dεp= 2+
Fp −+ Ḟp,inj ,i↵usion coefficientas [144–146]∂t is εrepresented∂εptp−cool
tesc
p ∂εp

-1

IG. 2. Relationship between the observed X-ray luminosy, LX,obs , and the X-ray luminosity obtained by the model
alculation, LX,calc . The green squares are LLAGNs withm > 10 3 , while the blue circles are those with ṁ < 10 3 .
he dotted line represents LX,obs = LX,calc , and cyan bandndicates LX,obs /1.7 < LX,calc < 1.7LX,obs , in which all thereen squares are located.

-2 -1

AGN Corona Model

25
We pl
V. CASCADE GAMMA-RAY EMISSION
Murase, SSK, Meszaros 2020
Fig. 3. A
Hadronuclear and photohadronic processes produce
a flat sp
very-high-energy (VHE) gamma rays through neutral pionattenuatio
decay and high-energy electron/positron pairs throughelectronMULTIMESSENGER
TESTS
OFcharged pion decay and the Bethe-Heitler process. The
VHE gamma rays are absorbed by soft photons through theγγ → eþ e− process in the RIAF, and produce
additional GAMMA-R
V. CASCADE
MULTIMESSENGERTESTS OFpairs.
COSMIC-RAY
ACCELERATION…
high-energy electron/positron
The high-energyeþ e−Hadronuclear and photohadro
pairs also emit gamma-rays through synchrotron processes,
where raεγ
very-high-energy
(VHE)
gamma
V. CASCADEGAMMA-RAY
EMISSION leading to We plot−2pr
inverseCompton scattering,and bremsstrahlung,
decay
and
high-energy
electron/ðx− suffx Þelectromagnetic cascades. We calculate the cascade emis-Fig. 3. A
Hadronuclear and photohadronic processes
producecharged
pion
decay
and
the
Beth
the
Heav
sion by solving the kinetic equations of photons anda flat spectr
very-high-energy (VHE) gamma rays throughneutral pion
VHE
gamma
rays are absorbedof thebycriattenuationbelectron/positron pairs (see Refs. [87,159,160]):þ−
decay and high-energy electron/positronγγ →pairse e throughprocess inelectron-posit
thecanRIAF,
see fcharged pion edecay and the Bethe-Heitler
process.electron/positronThe
high-energy
pairs
tabulated∂nεe∂e
VHE gamma rays
soft
photons
throughthe
3 fo½ðPIC þ Pby
þ
P
þ
P
Þn
þ are absorbed
pairs
also
emit
throughfittingsynff
Cou εe % gamma-rays
e the RIAF, and produce additional
γγ → eþ e− ∂tprocess∂εininverse Compton
scattering,and breformedweþ−nεe pairs.
high-energy electron/positron
The electromagnetic
high-energy e e cascades.LAT
ðγγÞ
inj
sen
We
calc
-4 ¼ n
_
_
−
þ
n
;
ð31Þ
εe
εe
10
pairs
also
emit
gamma-rays
throughsynchrotron
processes,twithεγ equais1sion bysolving
the where
kinetic
esc electromagne• Equa ons forccascades
RQ
AGN
(this
work)
inverse Compton scattering,and bremsstrahlung,leading to−2
predicted
electron/positron
Refs.
[
RQ AGN (cascade) (this work) pairsðx(see
−
x
Þ
lnðx
γ
γ
γ We calculate
electromagnetic
cascades.
the (thermalcascade emisRQ
AGN
e) 
-5
the cases∂n10nεγ nεγ
εγ
ðICÞ
ðffÞ
ðsynÞ
inj
the
Heavisideclusters) ð32Þ
_ εγphotons− kinetic
þ n_ εγequations
þRLn_ εAGN
þofn(inþe n_ εγ ;and
¼ − thesion by solvingγ
better
sen
∂n
∂
∂t
t
t
ε
of
the
criticale
RL AGN (w. blazars)
electron/positron γγ
pairs esc(see Refs. [87,159,160]):
Psyn þ Pþ  ½ðPIC þgamma-r
total
∂εe
∂t
can see flat sp
-6
range
du
i
10
e
e
(i ¼ ðγγÞ
e or γ),wheretabulatedvaln∂nεenεi is∂ the differential number density
ε
inj
e
has
_ εe formu
þn; a higðxxÞ½ðPIC þ Psyn þ Pff þ PCou Þneεe % ¼ n_ εe − fittingþ
n_ εi ∂tis thetesc flux is to∂εeparticle source term from the process xx
formed with t-7e[xx ¼10IC(inverse
Comptonscattering),
γγ
(γγ
pair
pronεebesensitiv
challe
ðγγÞ
inj
LAT
γ
γ _ inj γ¼ n_synþ n_ εe ; or ff (bremsstrahlung)],
ð31Þ
εe −
duction),(synchrotron),is∂nncurrentaεγεγ N εni εγ withðICÞ
ðf
t
a
10-ye
esc SSK+ 2021; Kheirandish,
See
also
SSK+
2019;
SSK
_ εγ Sgr_ εγ−
þ n2021
þ nA*¼ −Murase,the injection
term
from
the
hadronic
interaction,
and
P
is
yy
-8
∂t
t
t

ti

ti

ti

2

FIG. 1. Schematic picture of the AGN disk-corona scenario.
Protons are accelerated by plasma turbulence generated in the

26

Mul -messenger Spectra from NGC 1068
-6

10

ν
Cascade γ

Murase, SSK+ 2020
-7

Cascade emission at 10 MeV—>Testable by MeV γ ray satellites

E FE [GeV cm s ]-2 -1

IceCubeeASTROGAM-810

AMEGO-X-9

Fermi LAT

10

-10

10

-11

fl

10

ti

•

Both pp & pγ (with X-rays) contributeto resul ng neutrino ux

ti•

CR accelera on is suppressed by
Bethe-Heitler process with UV photonsti•

Possible to explain IceCube data
without overshoo ng γ-ray data

ti

•

10

-4

-3

-2

-1

10 10 10 10

0

1

2

3

4

5

6

7

10 10 10 10 10 10 10 10E [GeV]

ESO 138-1
NGC 7582
NGC 1068
NGC 4945
NGC 424
UGC 11910
CGCG 164-019
NGC 1275

1
0.7
0.5
0.8
0.7
0.5
0.4
0.3

0.1 (0.02)
0.2 (0.04)
0.2 (0.05)
0.5 (0.2)
0.4 (0.2)
0.4 (0.4)
0.4 (0.3)
0.4 (0.4)

Nearby Seyfert galaxies

The Astrophysical Journal, 922:45 (17pp), 2021 November 20•

0.4 (0.3)
0.3 (0.08)
0.4 (0.3)
0.4 (0.2)
0.4 (0.4)
0.4 (0.2)
0.5 (0.4)
0.5 (0.4)
Kheirandish,Murase, SSK 2021
0.5
(0.4)
0.5 (0.4)
0.5 (0.5)
0.5 (0.5)
0.5 (0.5)
0.5 (0.5)
0.5 (0.5)
0.5 (0.5)

27• Stacking nearby Seyferts
Our model predicts Lν ∝ LX
Table
4
—>TeV
listenergies,
up bright
ν-source
candidates
the source
beneﬁts froma 100% visibility in

Prospects
Observation
KM3NeT. Therefore, the likelihoodfor itsfor
observation
is high, of nearby Bright Seyfert Galaxies in One Year of KM3NeT O

trophysical Journal, 922:45 (17pp), 2021 November 20t, we model neutrino production assuming the stochastic
ration scenario. As mentioned earlier, in this scenario,
utrino spectrum has a more complicated shape than apower law. Accommodating the IceCube ﬂux at TeV
es requires a relatively high normalization, while theum has to cut off fast enough that the spectrum drops
d 100 TeV. Such conditions would result in a high levelpressure in the corona model.
rder to maintain realistic scenarios, we restrict ourselvesrange of parameters for which the ratio of the CR
re (PCR) to the thermal pressure (Pth) is bound to less
.5. In this limit, the nonthermal energy is equal to half ofavitational binding energy at the coronal radius without
g room for thermal particles. Although the coronala may be heated more through magnetic ﬁelds connected
inner disk, we assume 0.5 as the maximal case in thisand the neutrino spectrum peaks at ∼5 TeV and falls
y around 20 TeV. We refer to this scenario as “High CR
re.”consider the second scenario for neutrino emission from1068 assuming coronal emission from stochastically
rated particles, where instead of matching the ﬂux at
we match the diffuse neutrino ﬂux at tens of TeV,
ated by the medium-energy excess in the neutrino
um. In this case, as shown previously (Murase et al.
we adopt parameters that can explain the high-energy
no ﬂux excess observed at medium energies (Aartsen
2020a). In this case, PCR/Pth is set to ;0.01. Here, theno spectrum peaks at ∼40 TeV, which corresponds to alevel of neutrino ﬂux compared to the previous scenario.
fer to this case as “Modest CR pressure” hereafter.
se results are compatible with the spectra presented
usly by Murase et al. (2020) where the CR pressure
ered to explain the medium-energy neutrino ﬂux and1068 are found at the level of ∼1% and ∼30% of theal pressure, respectively. Here, we allow the pressureo be as high as 50% to explain the soft spectrum reported
GC 1068 by the IceCube Collaboration (Aartsen et al.

Kheirandish, Murase, & Kimuraand can exceed 3σ in 3 yr of operation for the stochastic
acceleration scenario with High CR pressure.
p-value 1 yr (3 yr)
As the signal events from the rest of the sources in the list
Source fall short of yielding a statistical
Visibility
Stochastic(High
CR
Pressure)
Stochastic(ModestCR
Pre
signiﬁcance in 3 yr, we now
turn to the prospects for observation of neutrino emission in a−8
0.2 (0.07)
Cen A stacking analysis. We only consider
0.7 the Modest CR pressure0.001 (9.3 × 10 )−5
scenarioinstochastic
acceleration
since
emission
under
either0.2 (0.09)
Circinus Galaxy
1.0
0.008 (1.9 × 10 )of the other two scenarios should be identiﬁed by IceCube. In
ESO 138-1
1
0.1(0.02)
0.4(0.3)
addition to KM3NeT, we consider IceCube-Gen2 for the
stacking search in this scenario.
NGC 7582
0.7 Here, we assume that the 0.2 (0.04)
0.4 (0.3)
effective area for IceCube-Gen2 is ∼5 times larger than the
NGC 1068
0.5
0.2 (0.05)
0.4 (0.4)
current IceCube detector.
1. Modeled neutrino spectrum for NGC 1068 compared to the best-ﬁt
NGC 4945We present Figure
0.5 (0.2)
0.5 (0.4)
thep-values
expected
for theModest CR
ﬂux (yellowband) reported0.8
by
the IceCube Collaboration
10 yr point-source
study (Aartsen et al. 2020b). The red line shows the expected ﬂux in thepressurestochastic
accelerationfor
KM3NeT
stochastic
acceleration scenarioscenario
matching IceCube’s
best
ﬁt at TeVs. Thetogether
NGC 424with the ones forpurpleIceCube-Gen2
0.4
(0.2)
0.5
(0.4)
line depicts the ﬂux 0.7
that would give the medium-energy neutrino ﬂux,
inreported
Figure10.analysis
We(Aartsen
project thecompatible with the total neutrino ﬂux
in the cascade
et al. 2020a). The blue line presents the ﬂux expected for the magnetic
UGC 11910
0.5
0.4 (0.4)
0.5 (0.5)
prospects for identiﬁcation
reconnection scenario.of neutrino emission from the brightsources assuming Figure
an angularresolutionof 0°ﬂuxes. 3 (solid)
and 0°. 7
CGCG 164-019
0.4modeled neutrino
0.4 (0.3)
0.5 (0.5)
1 shows the three
from NGC(dashed) for each1068.
detector.
We should
that
ourbyestimation
We also projected
the best-ﬁt note
spectrum
reportedtheCollaboration.
The best-ﬁt power-law spectrum
NGC 1275
0.3
0.4
(0.4)
0.5
(0.5)
of the prospectsIceCube
for
identifying
Seyfertgalaxiesarequite
corresponds to the ∼51 excess neutrinos found from the
Figure 10. Prospects for observation of the bright Seyfert galaxies in the nextdirectionof NGC1068. The shaded
area shows of
the uncertainty
conservative, giventhat
an angularresolution
0°. 3 or better isgeneration neutrino telescopes: KM3NeT and IceCube-Gen2. The solidon the ﬁtted spectrum as reported by IceCube. As shown, all•
Futureν from
AGN
not that far-fetched
forneutrino
KM3NeT.
The theexpected
improvements(dashed)lines detectors
show expectationsshould
for 0°. 3 (0°. 7)detect
angular resolution
for themodeledspectra are within
68% uncertaintyof themeasured spectrum. in
TheIceCube-Gen2
parameters that we adapt
each make it
Modest CR pressure scenario. The thick lines show the prospects forin the angular reconstruction
willinalso
scenario for particle acceleration and interaction efﬁciency are—>
testablebyfutureexperiments
identiﬁcationof the 10 nearby
brightsources neutrino
in Table 2 in a stacking
analysis.
easier to identifypresented
thesein sources.
In
fact,
our
estimates
indicate
Table 1. The common parameters among different
The thin lines show the prospects for identiﬁcation of the sources in the
scenarios
are
the
same
as
in
Muraseetal.
(2020).
The
injected
that achieving ﬁner angular resolutions at ∼10–30 TeV is

TeV energies, the source beneﬁts from a 100% visibility in

RIAFs in LLAGN

LLAGN

log (νLν /ergs s–1)

ev. Astro. Astrophys. 2008.46:475-539. Downloaded from www.annualreviews.org
by Osaka University on 12/18/13. For personal use only.

RIAF

1 cm

1 mm

100 µm 10 µm

1 µm

1000 Å 0.1 keV 1 keV

44

42
log Lbol /LEdd ≥ 0–1 ≤ log Lbol /LEdd < 040–3 ≤ log Lbol /LEdd < –1
log Lbol /LEdd < –3

38

PHYSICAL REVIEW
9

10

11

12

13

14

15

16

ti

ti

fl

ffi

4
0

1

x

ti

ti

ti

N

ti

G

O

A

17

18

log (ν/Hz)

QSO

Figure 7

Wind?

Composite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm
(Adapted from L.C. Ho, in preparation.)

Turbulence?

nuclei (Ho 1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho & P
2001). Defining radio-loudness based on the relative strength of the radio and X-ray emis
RX ≡ ν Lν (5 GHz)/LX , Terashima & Wilson (2003b) also find that LINERs tend to be ra
loud, here taken to be RX > 10−4.5 . Moreover, the degree of radio-loudness scales inversely
Lbol /LEdd (Ho 2002a; Terashima & Wilson 2003b; Wang, Luo & Ho 2004; Greene, Ho & Ulv
Accre
on 2007;
Shock?
2006; Panessa et al. 2007; Sikora, Stawarz& Lasota
L.C. Ho, in preparation; see Figure 1
In a parallel development, studies of the low-luminosity, often LINER-like nuclei of FR I rgalaxies also support the notion that they lack a UV bump. M84 (Bower et al. 2000) and(Sabra et al. 2003) are two familiar examples, but it has been well documented that FR I n

Protons in coronae & RIAFs are collisionless → Non-thermal proton produc onS

10 keVHo 2008

LLAGN: No blue bump & X-ray→Op cally thin ow
Radia vely Ine cient Accre on Flow (RIAF)

•

10 cm

46

QSO: Blue bump & X-ray→Op cally thick disk + coronae

•

L

L

Q

28

Cosmic
High-energy
Background
from
RQ
AGNs
ger than t for " > 1.5 ⇥ 10 GeV
26, 67])
fall

p

29

PHYSICAL 4
REVIEW LETTERS 125, 011101 (2020)r "p > 5.1 ⇥ 103 GeV for ṁ ⇠ 10 3 ,the proton spectrum appears at adue to its hard spectral index and6, 61]. For models B and C, thenosity is almost identical to the incause the infall dominates overQSO
then all the energy range.

ti

tt

fl

dz

Z

L"i ⌧i,IGM
p
dLH↵ ⇢H↵e,"i(1 + z)3 ⌦m + ⌦⇤(3)
where ⇢H↵ is the H↵ luminosity function, ⌧i,IGM is
the optical depth in intergalactic medium, and we use
H0 ⇠ 70 km s 1 Mpc 1 , ⌦M ⇠ 0.3, and ⌦⇤ ⇠ 0.7.
LLAGN
G. 1. Schematic picture of the AGN disk-corona scenario. H↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡otons are accelerated by plasma turbulence generated in the
s1
s2
(⇢
/L
)/[(L
/L
)
+
(L
/L
)
], where ⇢⇤ ' 4.11 ⇥
⇤
⇤
H↵
⇤
H↵
⇤
ronae,
andproducehigh-energy
neutrinosandcascadedlisionsandphotomeson
interactions
5
3
41
1
mma rays via interactions with matter and radiation.
•
10
Mpc
,
L
=
3.26
⇥
10
erg
s
, s1 =
2.78, & 10 TeV neutrinos
X-ray⇤ used in this work, for LX ¼ 1042, QSO:
FIG.
2.
Disk-corona
SEDs
decay to neutrinos. We calculate the1043 , 101046 ergWe
s−1 (frombottom to top). See
text luminosity funcand
s442, 10=45 , and1.88.
extrapolate
thisom
pp collisions
using thefor details.
stributions(SEDs) are constructedfromformalism
the data and from
38
1
•
tion
to
L
=
10
erg
s
,
below
which
the PaloLLAGN:
MeV
γ
&
PeV
neutrinosminmpiricalrelations,
and
then
we
computeneutrinoand
For the neutrinos
by p interacCoronae
scade gamma-ray spectra by consistently
solving
particle emar survey finds a hint
γ bythermalof2021
a flattening [69]. The surSSK+
analytic
prescription
given
in Ref. of Fig. 2 (see Supplemental Material [68] for details), whichansportequations.
We demonstrate
the importance• LCopious
photonsvey
also
indicatesacorrelation
between
and
L
for
RIAFs
X
H↵
enablesustoquantitatively
evaluateCR,
neutrinoandturemoderate
MeV gamma-ray
observations
revealing thehemagnetic
fieldforstrength
ν byp [69]. We use a correction
cascade gamma-ray
emission.
LXnon-thermal⇡ 5 7LH↵
igin of IceCube neutrinos especially in the medium- LLAGNs:
we canignorethe e↵ect
meson
cient
γγ
—>
e+eNext we estimate the nucleon density np and coronal → eergy(∼10–100
TeV) range
and for of
testing
neutrino factor
field= B.
LX/L= 6.0.
Then,
the luminosity
CoronaeH↵
X/H↵
magnetic
strengthLet
usconsideracoronawith
mission1068 andneutrinos.
other AGN. Then,focusfrom
onNGC
sub-PeV38
1x
the
radius
R
≡
RR
and
the
scale
height
H,
where
R
is
integrationisperformed
in
the
range
of
10
erg
s

S
→
strong
GeV
γ
a
enua
on
We use a notation with Qx ¼ Q × 10 in CGS units.
atiois (⌫e , ⌫prescription(1, 2, 0) at
2 is theµ , ⌫⌧ ) =41
1
the normalized
coronalradius
and
R
¼
2GM=c
RIAFs
Phenomenologicalof AGN disk coronae.— L
S

⌘
ṁL
/(

)
'
4.2
⇥
10
erg
s
.
H↵
rad
Edd
X/H↵
bol/X
Schwarzschild radius. Then the nucleon density is → GeV
1)
on Earth,duea phenomenologicalto the neutrino
ose beginby providing
disk-corona
ux
below
the
Fermi
data
Since
dimmer
AGNs
tend
to
have
weakerredshift
evoluexpressed by np ≈ τT =ðσ T HÞ, where τT is the Thomson
odel
based
on
the
existing
data.
Multiwavelength
SEDs
agation. The hadronic interactions
optical[70–72],
depth thatwe
is typically
The standard
Seyfert galaxies have been extensively
studied,consist- tion
γ by EM
cascadesassume∼0.1–1.
no redshift
evolution of the luag ofrays
andelectron/positron• AGN cores can account for
accretion theory [69,70] gives
the coronal scale height
severalcomponents;
radio emission (seepairs,
Ref. [60]),
p
ffi
ffi
p
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
ffi
minosity
function.
The
mass
of
SMBHs
in
local
Seyfertfrared emission cascades.
from a dust torus
optical andH ≈ ðCs =V K ÞRRS ¼ RRS = 3, where Cs ¼ kB T p =mp ¼
omagnetic
We[61],
calculatepffiffiffiffiffiffiffiffiffiffiffiffiffi
pffiffiffiffiffiffi
does
not
show
any
correlation
with
X-ray
lutraviolet components from an accretion disk [62], and galaxies
keV-MeVγ&
TeV-PeVνbackground
¼
GM=R
¼
c=
6R
is
the
sound
velocity,
and
V
K
by
solvingthekinetic
equationsofpffiffiffiffiffiffirays from a corona [33]. The latter two components are minosityandH↵ luminosity
c= 2R is the
Keplerianvelocity. For an [73].
opticallyRef.
thin [74] provides alevantfor thiswork.
rs andphotons.
We approximately
corona, theelectron temperature
is average
estimated andbysampleof
LLAGNs,andthemedian values
The “blue” bump, which has been seen in many AGN, is
T e ≈ εX;cut =ð2kB Þ, and τT is empirically determinedfromSeealso Murase, SSK+ 2020 PRL; SSK+ 2019, PRD; SSK+ 2015
onprocesses
byBethe-heitler
protributed to multitemperature blackbody emission from a
ffi

c
i =
4⇡H0

Z

Index• Introducon• Neutrino emission models in Seyfert galaxies• Neutrino emission from AGN coronae• Sub-GeV gamma rays from Seyfert galaxiesti• Summary30

AGN-Starburst connec on
Many Seyfert galaxies are forming stars in central regions• Star-forma on ac vity can produce cosmic-rays,leading
to gamma-ray and neutrino produc ons
Taniguchi• Famous example:
NGC 1068, Arp 220, NGC 4945, Circinus galaxy•

270

0LQRU&0HUJHUTaniguchi 2013
6H\IHUW

8/,5*

4XDVDU

ti

ti

ti

ti

0DMRU&0HUJHU

6WDUEXUVW

31

Gamma rays from 3NGC 1068-610coronastarburst

Ajello, Murase, McDaniel 2023
-7

-2 -1

E FE [GeV cm s ]10

IceCubee-ASTROGAM•

Gamma rays by starburst ac vity:
low-E cuto at sub-GeV by pion decay•

Sub-GeV γ-ray spectrum in NGC1068:
extending to ward lower energies—> need addi onal component•

Gamma-ray ux consistent withhadronic cascade by corona modelcorona ν-8

10

AMEGO-X-9

Fermi-LAT

10

MAGIC
-10

10

-11

10

-3

-2

10

10

-1

10

0

10

1

2

10
10
E [GeV]

3

10

4

10

5

10

ti

ti

fl

ff

igure 2. Model spectra of MeV-TeV gamma-ray emissionom NGC 1068, compared to Fermi-LAT data obtained by

32

33

Gamma rays from NGC 4151?

Publi

4

•

UFO scenario•

Accre on-shock + Jet scenario
10 11
Jet
Corona
Total

Flux F (erg/cm2/s)

Peretti et al. 2023

10 12

Inoue et al. 2023

10 13

Neutrino
10 14

10 15 6
10
Figure 2. Observed W-ray flux compared with the multimessenger model•prediction. Data points (cyan) are shown along with the 1 f uncertainty
band. The thick blue line represents the predicted W-ray flux while the dotted•red curve is the associated per-flavor neutrino flux.

Gamma-ray

Corona :Fermi/LAT

107

108

109

1010 1011 1012 1013 1014 1015

Energy E (eV)

Perre et al. reported gamma-rays from NGC 4151
Fig. 2. Same as Fig. 1, but showing the energy range only in 10 –10 eV
Gamma-rays can be explained by the ultrafast ou
lowtogether with the Fermi-LAT data (Peretti et al. 2023b).
• The gamma-ray ux is low and neutrino cannot be detected
based
onother
accre
on (>
shock
scenarioNGC41512 . No
significant
5 f) emission
is found in the rest
flux above 10 TeV is an order of 10 14 erg cm 2 s 1 , which

ion of). Thele the
mond.

tf

fl

ti

ti

tti

of the map. Nevertheless, we update the SM by including all warm

6

15

34

Summary
QSO

FIG. 1. Schematic picture of the AGN disk-corona scenario.
Protons are accelerated by plasma turbulence generated in thecoronae, and produce high-energy neutrinos and cascaded
gamma rays via interactions with matter and radiation.

•

IceCube discovered evidence of neutrino signal from Seyfert galaxy•

2. Disk-corona SEDs used in this work, for LX ¼ 10
Accre on shock & failed wind scenario can explain ν data, FIG.
1043 , 1044 , 1045 , and 1046 erg s−1 (from bottom to top). See t
for details.
distributions(SEDs)are constructedfrom theand from
but they need to assume
ine
cient
accelera
ondataparameters

ti

fl

ti

ti

ti

ffi

•

ti

•

empirical relations, and then we compute neutrino andcascadegamma-rayspectrabyconsistently
solving
particleCoronae around SMBH
canexplain
ν
data
for
NGC
1068
without
overshoo
ng
γ
data
Fig.
2
(see
Supplemental
Material[68] for details), whi
transport equations. We demonstrate the importance ofus totest
quantitatively evaluate CR, neutrino a
MeV gamma-ray
observations
for will
revealingtheand future neutrinofuture& MeV
γ-ray observa
ons
provideaenablesrobustcascade gamma-ray emission.
origin of IceCube neutrinos especially in the mediumNext we estimate the nucleon density np and coronenergy(∼10–100
TeV)
range
and
for
testing
neutrino
Combining a contribu
on
from
LLAGN,
AGN
accre
on
ows
can
be
the
sourceofmagnetic field strength B. Let us consider a corona w
emission from NGC 1068 and other AGN.
xtheradius
R
≡
RR
and
the
scale
heightH,
whereR
S
the cosmic neutrino We
backgroundfor Qalltheenergy
range
(1
TeV
10
PeV)
use a notation with
¼
Q
×
10
in
CGS
units.
x
2 is tthenormalized
coronalradius
and
R
¼
2GM=c
Phenomenological prescription of AGN disk coronae.—
S
Schwarzschild radius. Then the nucleon density
We
begin
by
providingaphenomenologicaldisk-corona
Starburst ac vity can
explainγ-raysofE>
GeV,
expressed by np ≈ τT =ðσ T HÞ, where τT is the Thoms
model based on the existing data. Multiwavelength SEDs
but cannot explain ofneutrinos& havesub-GeVgamma-raydata optical depth that is typically ∼0.1–1. The standa
Seyfert galaxies
been extensively
studied, consistaccretion theory [69,70] gives
heig
ing of several components; radio emission (see Ref. [60]),
pffiffi the coronalpscale
ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi
infrared emission from a dust torus [61], optical andH ≈ ðCs =V K ÞRRS ¼ RRS = 3, where Cs ¼ kB T p =mp
pffiffiffiffiffiffiffiffiffiffiffiffiffi
pffiffiffiffiffiffi
ultraviolet components from an accretion disk [62], and
c= 6R is the sound velocity, and V K ¼ GM=R
ti

•

ti

LLAGN