活動銀河核の降着流からのニュートリノ放射東北大学（学振PD）

木村成生
References
1) Murase, SSK, Meszaros, arXiv:1904.04226
2) SSK, Murase, Meszaros, 2019, PRD, 100, 083014
3) SSK, Murase, Meszaros in preparation
see also: SSK, Murase, Toma, 2015, ApJ, 806, 159
Collaborators
Peter Meszaros (Penn State)
Kohta Murase (Penn State; YITP)
Kenji Toma
(Tohoku Univ.)

第32回理論懇シンポジウム@国立天文台�����2019年12月

Index2•

IceCube Neutrinos•

Non-thermal Processes in Accretion Flows•AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226•

LLAGN RIAF model SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

•

SummaryDetection of AstrophysicalNeutrinos

3

PTEP 2017, 12A105
(a)

(b)

Shower event (νe)
electron and tau neutrinos

Fig. 3. (a) Light pool produced in IceCube by a shower initiated by an electenergy is 1.14 PeV, which represents a lower limit on the energy of the neutri
dots represent sensors with no signal. For the colored dots, the color indicat
purple (late) following the rainbow, and size reflects the number of photon
track traverses the detector at an angle of 11◦ below the horizon. The depo
2.6 PeV.

••muon neutrinos
signals are sorted into telltale patterns of light that reveal the direc
incident neutrino.
Even at a depth of 1450 m, IceCube detects a background of at
originating in the Southern Hemisphere at a rate of 3000 per second. Tw
neutrinos. Traditionally, neutrino searches have focused on the obse

IceCube experiment reported detection of astro-ν (E ~ PeV) in 2013
Origins of astro-νs have yet to be understood

T
4

Neutrino Production Processw
S. R. KELNER, F. A. AHARONIAN, AND V. V. BUGAYOVsomby• pp inelastic collision

σ[mb]

νμ
μ-

π-

target
HEP

p

p

νμ
e-

νe

γ

π0

γ

νμ

μ+

π+

e+
νe

νμ

Kelner 06

Tthinel = (34 .3 + 1 .88 L + 0 .25 L
4 2
E th

2

)×

1−
, mb
Ep
FIG. 8. The multiplicity of photons and leptons produced
inEp’[TeV]
one interaction of a relativistic proton with 2.7 CMBR.

th

T

σ[mb]
• Photomeson production (pγ) FIG. 11.
Inelastic cross section of p-p interactions approxii.e

HEP

p

＜
〜
〜

target γ

Kelnerdata
08 are frommated by Eq. (79). The experimentalth
http:wwwppds.ihep.su:8001/c5-5A.HTML, the open points cor
respond to the cross sections which are used in the SIBYLL code
pr

γ
γ
γγ

π0

ν
νμμ
π+

μ+
νμ

e+
ννee

!
" !
"
cnH
E!
E!
q! !E! " # n~
"
m $
J m $
;
K! inel p K! p p K!

(77)

w
n. For the procedure described below, un~where K! # #=~p
and K! are free parameters.
"E
The emissivity of gamma rays is related to q! !Eas!"
E’[GeV]γ
though the equationphZ 1 q! !E! "
in
dN$
#2
(78)
p#################### dE! ;
ac

5

•

Source Candidate
Cosmic-ray accelerators
mainly pγ interaction• Active Galactic Nuclei (AGN)•

Cosmic-ray reservoirs
mainly pp interaction• Star Forming Galaxies (SFG)
Tamborra+14; Senno+ 15; Sudoh+17

Murase et al. 2014
Kimura et al. 2015
Inoue et al. 2019

• Galaxy Group / Galaxy Cluster• Gamma Ray Bursts (GRBs)

Murase et al. 2013

Murase+13
Fang & Murase 18

6Model-independent
Constraints
eff [Gpc

discovery potential for
steady source candidates
105
in 10 years
starburst & GC/GG-int
104

effective local density

3

]

106

102

103

RQ AGN

GC-accRL AGN10BL Lac

1
0.1
10 2
10 3

•

LL AGN

IceCube5 IceCube
20 IceCube

FSRQ

1039 1040 1041 1042 1043 1044 1045 1046 1047 1048
effective neutrino luminosity L [erg/s]

dark sources below 100 TeV not seen in g’s ?
νflux@10 TeV > γ flux@100 GeV
gamma rays cascade in the source to lower energy→ accompanying γ overshoot Fermi data
Figure 2:• Left:
Comparison ofdetection
the diffuse neutrNo
point-source→ ν sources should be opaque tive
to local density and luminosity of extragalactic
disfavors luminous sources
TeV γ rayseral candidate
populations(î)Jetted
by the TDEs)
required ne(GRBs,
Blazars,
Murase et al. 2013, 2016; Ahler & Halzen 2017
flux [17] (see also [25]). The lower (upper) edgetion. The dark-blue-shaded region indicates IceC

7•

Source Candidate
Cosmic-ray accelerators
mainly pγ interaction• Active Galactic Nuclei (AGN)•

Cosmic-ray reservoirs
mainly pp interaction• Star Forming Galaxies (SFG)• Non-beamed AGN
Tamborra+14; Senno+ 15; Sudoh+17

Murase et al. 2014
Kimura et al. 2015
Inoue et al. 2019

• Galaxy Group / Galaxy Cluster• Gamma Ray Bursts (GRBs)• Low-luminosity GRBs

Murase et al. 2013

Murase+13
Fang & Murase 18

8Hint of ν point sources

eft: The 2D distribution of events in one year of data for the final event selection as a function ofd declination and estimated energy. The 90% energy range for the data (black), as well as simulated
cal signal Monte-Carlo (MC) for an E 2 and an E 3 spectrum are shown in magenta and oranges a guide for the relevant energy range of IceCube. Right: The e↵ective area as a function of neutrino
IC86 2012-2018 event selection averaged across the declination band for several declination bins using
simulated data.

•

Point source searchwith 10-year data set

IceCube 2019

M77 (NGC 1068)

ap of -log10 (plocal ), where plocal is the local pre-trial p-value, for the sky between ±82 declination in
ordinates. The Northern and Southern hemisphere hotspots, defined as the most significant plocal inthat hemisphere, are indicated with black circles.

• Hottest Point (2.9σ) :M77 (NGC 1068; Seyfert 2)

GIC observations and about 4 hrs of H.E.S.S. observations [31, 39, 40] in Fig. 9.

Index9•

IceCube Neutrinos•

Non-thermal Processes in Accretion Flows•AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226•

LLAGN RIAF model SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

•

SummaryRadio-quiet AGNs

10

1 cm

1 mm

100 µm 10 µm

1 µm

1000 Å 0.1 keV 1 keV

QSO•

LLAGN x 104

•

44

42

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

38

9

10

11

12

QSO UV X-ray

13

14

15

16

17

QSO: Blue bump & X-ray→Optically thick disk + coronae
LLAGN: No blue bump & X-ray→Optically thin flow
[Radiatively InefficientAccretion Flow (RIAF)]

18

log (ν/Hz )

Opt.

~>

log (νLν /ergs s–1 )Ho 2008

ure 7

10 keV

LLAGN

corona

~~
~~
~~
>

~~
~~
~~
>
~~
~

mposite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm.
pted from L.C. Ho, in preparation.)ei (Ho 1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho & Peng
BH radio-loudness based on the relative strength of the radio and X-ray emission,
BH
1). Defining≡ ν Lν (5 GHz)/LX , Terashima & Wilson (2003b) also find that LINERs tend to be radiod, here taken to be RX > 10−4.5 . Moreover, the degree of radio-loudness scales inversely with/LEdd (Ho 2002a; Terashima & Wilson 2003b; Wang, Luo & Ho 2004; Greene, Ho & Ulvestad
6; Panessa et al. 2007; Sikora, Stawarz & Lasota 2007; L.C. Ho, in preparation; see Figure 10b).
n a parallel development, studies of the low-luminosity, often LINER-like nuclei of FR I radio
xies also support the notion that they lack a UV bump. M84 (Bower et al. 2000) and M87
bra et al. 2003) are two familiar examples, but it has been well documented that FR I nuclei
d to exhibit flat αox (Donato, Sambruna & Gliozzi 2004; Balmaverde, Capetti & Grandi 2006;
ozzi et al. 2008) and steep slopes in the optical (Chiaberge, Capetti & Celotti 1999; Verdoes
jn et al. 2002) and optical-UV (Chiaberge et al. 2002).

Hot plasma

cold disk

X-ray

X-ray
~>

10 cm

46

12:45

~~

15 July 2008

>

ARI

~~~

6-13

Hot plasma

RIAF
cold disk

Radio-quiet AGNs

11
15 July 2008

10 cm

46

1 cm

12:45

1 mm

100 µm 10 µm

1 µm

1000 Å 0.1 keV 1 keV

QSO

QSO: Blue bump & X-ray→Optically thick disk + coronae
LLAGN x 104• LLAGN: No blue bump & X-ray
log L /L
≥0
→Optically thin flow
< 0 RIAFs are collisionless–1 ≤ log L /L&
Protons in coronae[Radiatively Inefficient–3 ≤ log L /L < –1
log L /L < –3 proton production→ Non-thermal
Accretion Flow (RIAF)]
bol

40

bol

9

10

11

12

QSO UV X-ray

Edd
bol

Edd

bol

Edd

Edd

** electrons are collisional **log (ν/Hz )

13

14

15

16

17

Opt.

18

LLAGN

corona

~~
~~
~~
>

~~
~~
~~
>
~~
~

mposite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm.
pted from L.C. Ho, in preparation.)ei (Ho 1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho & Peng
BH radio-loudness based on the relative strength of the radio and X-ray emission,
BH
1). Defining≡ ν Lν (5 GHz)/LX , Terashima & Wilson (2003b) also find that LINERs tend to be radiod, here taken to be RX > 10−4.5 . Moreover, the degree of radio-loudness scales inversely with/LEdd (Ho 2002a; Terashima & Wilson 2003b; Wang, Luo & Ho 2004; Greene, Ho & Ulvestad
6; Panessa et al. 2007; Sikora, Stawarz & Lasota 2007; L.C. Ho, in preparation; see Figure 10b).
n a parallel development, studies of the low-luminosity, often LINER-like nuclei of FR I radio
xies also support the notion that they lack a UV bump. M84 (Bower et al. 2000) and M87
bra et al. 2003) are two familiar examples, but it has been well documented that FR I nuclei
d to exhibit flat αox (Donato, Sambruna & Gliozzi 2004; Balmaverde, Capetti & Grandi 2006;
ozzi et al. 2008) and steep slopes in the optical (Chiaberge, Capetti & Celotti 1999; Verdoes
jn et al. 2002) and optical-UV (Chiaberge et al. 2002).

Hot plasma

cold disk

X-ray

X-ray
~>

42

•

~~

ure 7

44

~>

log (νLν /ergs s–1 )Ho 200838

10 keV

>

ARI

~~~

6-13

Hot plasma

RIAF
cold diskc%0amplify the magnetic field and form two inward- andregions (sandwiched by the reddish curved
Azimuthally
averagedVφ , ⟨Voutward-flowing
Thewithwhitealines
iso-contoursof ⟨V
where planes).
Ωc% ¼ eColor
gridsize Δ was
set toφ ⟩L , on the R − φ plane.
φ ⟩L .
% B0 =m% c. The p
streams
highareplasma
density
andﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ
contours in the background at Y ¼ 1.91
and X ¼ 1.91
23=2 ðvt% =Ωp% Þ, where Ωp% ¼ 8πne2 =m% is the pairstrong electric current called the channel flow. The reddish
show the angular velocity vy in the local rotating frame. In
Vbul, φ as
background
velocitysandwiched
for analysesbyoftwo
the surfaces
test-particle
plasma frequency.
The Alfvén velocity
is thedefinedasregions
in Fig. 1(c) show the12p
ﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ
ﬃ
the early stage at T orbit ¼ Ω0 t=2π ¼ 0.31 in Fig. 1(a), thesimulations
SectionV A ¼ B= 8πm% n, so that the plasma
β is inequal
to 3.2.high-density channel flow with ρ ≥ hρi þ 2σ ρ where hρimagnetic field lines are parallel to thez axis, and the
4 plots
thedensity (upper) and the magnetic
3v2t% =V 2A . The parameters used were Fig.
ðV A =Ω
¼colormaps
25,
andofσthe0 Þ=Δ
ρ are the average density and standard deviation of
Keplerian motion or differential motion ofvy can be seen as
energythe equatorial
plane. Theinmagnetic
fields domain,areðvt% =Ωc% Þ=Δ ¼ 56.4, V A =c ¼ 6.25 ×
10−3 . (lower)
N x , N y ,onand
density distribution
the simulation
respectively.
the color contour at Y ¼ 1.91, whereNtheare
reddish(bluish)frozendifferentially rotating
fluid elements
fall to field
the stretched by thethe grid
sizes in the x, y,
and inz thedirections,
The amplificationof thethat
magneticz
region corresponds to a positive (negative)
toroidalandvelocBH.
structuremotion
as seenmay
in theWe bycanthe magnetic field
Keplerianbefigure.
balancedrespectively,
we assumed N x ¼
N z This¼ Ncreatesinspiralz ¼ 300the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ity. As time passes, the vertical magnetic fields start togetalso see that the fluctuation of the density is much smaller thandistorted due to the MRI, and they are stretched out in thethat of the magnetic field energy density. This implies that the fastれるor 降着流).
toroidal(降着円盤direction becauseof the Keplerian motion atmodes are a sub-dominant component in the MRI turbulence.
T orbit ¼ 6.89 in Fig. 1(b). This stretching motion can
To clarify the importance of the modes of the MHD waves (fast,
重力エネルギーが変換され,明るく輝き,
amplify the magnetic field and form two
inward- and 場合によってはジェット
slow, and Alfven), we evaluate the Pearson correlation coefficientsoutward-flowing streams with a high plasma density andbetween the fluctuations of the density, δρ(R, θ, φ) = ρ − ⟨ρ⟩L ,る.
strong electric current called the channel flow. The reddish
and themagnetic energy, δB 2 (R, θ, φ) = B 2 − ⟨B 2 ⟩L . According
巨大ブラックホール
regions sandwiched by two surfaces in Fig. 1(c) show theto the (太陽質量の１億倍)
linear MHD wave theory, the fast mode has a positive
ルギー天体のエネルギー源
high-density channel flow with ρ ≥ hρi þ 2σ ρ where hρicorrelation, the slow mode has a negative correlation, and the Alfven
and σ ρ are the average density and standard deviation ofmode has no correlation. We evaluate the correlation coefficientsas a function of R and θ, and average over them with weights
density distribution in the simulation domain, respectively.
ジェット(赤)The AstrophysicalJournal, 822:88 (11pp), 2016 May 10
Kimura et al.
associated with the area in the meridional plane. The resultingThe
amplificationofthemagnetic
field
stretchedby
the-1の想像図Table
2
coefficients
indicatethatthedensityandmagnetic
energyareweaklyKeplerian motion may be balanced by the magnetic field
Model Parameters and Physical Quantities
anticorrelated: the value of the coefficient is −0.22 in the disc円盤abc
tend a for dt
Dz b
Dy b resolution
q
A
Model
T
 gyrog0 θ ! 0.45|)
xCregionesc(|cos
run A. DThe
lower
runs
have D0
8
光る)
20Trot
3.4
´
10
1.59
´
10-4
4
2
416
50
2.3
17
1.6
2.38
0.30
A1
higher coefficients,i.e.
theanticorrelationsareweaker,butnorun
8.5 ´ 107
20Trot3.17 ´ 10-5
A2
1
2
5628
200
2.4
24
1.6
1.91
0.25
銀河(白)
has
a positive
Therefore,
fast modes
do 2.79
not play5.38 ´ 10-4
20Trot6.8 ´ 108 correlation.
A3
8
2
94
40
2.0 the 16
2.0
0.25
3.5 ´ 108role in
1.50 ´ 10-4
15TrotB1
4
2 important
447this system.
50
2.2 result
19 is natural
1.5
2.38
0.31
an
This
in
the
sub3.7 ´ 108
1.45 ´ 10-4
25Trot
C1
4
2
438
50
2.2
17
1.6
2.46
0.28
Alfvenic and sub-sonic turbulence.
85
et al.spectra
2016;2.0
et al. 5.83
2019
¥ Finally,¥ the azimuthal
20Trot3.4 ´we
10 discuss
´ 10
X1
4
200SSK3.1
19
0.969
0.54
power
ofSSK
the turbulence¥¥
8.5 ´ 107
20Trot2.88 ´ 10-5
1
200
2.5
20
1.5
1.31
0.46
X2
ブラックホール(cf.
Sorathia
et
al.
2012;
Suzuki&Inutsuka
2014;
seeParkin&58¥¥6.44 ´ 10
20Trot6.8 ´ 10
X3
8
200
3.5
21
3.0
0.654
0.58
(太陽質量の10倍程度)
see also
Lynn
et al.We2014
電波銀河
Bicknell 2013 for three-dimensional
power
spectra).
take theNotes.
Fourier transformationin the azimuthal direction,a
Hoshino 2013; Hoshino 2015
光で光る)
In unit of tgyro,0 .
の観測
b!
In unit of D Bohm .
1(colorTime evolution of the magnetorotational
instability. Panels (a) and (b) show the magnetic field lines (greenish
Riquelme et al. 2012; KuntzFIG.
et 1al.
2016
c online).
In unit of Dp .
Xm = at√Y ¼ y=λX¼exp(−imφ)dφ,lines) and angular velocities in the background1.91 and X ¼ x=λ ¼ 1.91 (color contour), and (13)panels (c)–(e) depict the
2π

Particle Acceleration in
Accretion Flows

ラックホール降着円盤とは？
Particle-In-Cell Simulations

MHD + Test Particle Simulations

Numerical
Simulations

turbulenthigh-densityAlthough
regions as
planes.
Panelsandthe(c)length
are at the same time stage. Panel (f): The energy spectra during the MRI
thereddish
Lorentzcurved
contraction
mayalso (b)affectwhere
mfor= simplicity.
kφ Rline(kφ isis athewavenumber
theT orbit
φ direction).
at T orbit ¼ 0.31,
7.18,
8.84,
14.28.
The
dashedMaxwellian
fittinginfor¼ 7.18. Then,
scale 6.89,
that CRs
travel,weand
ignoreits effectWe note that the boundary wecondition
for averageCRs describedtake theof the power spectrum over the disc region:
above is different from that for the MHD "simulations. The latter
use the Galilean transformation, since the ﬂuid
is non061101-2|2 RdRdθ
|Xmvelocity
relativistic, G - 1  1. We cannot,(14)= the"Galilean transforPm use
RdRdθ
mation as the boundary condition for CRs because
the velocityof CRs after the Galilean transformation could exceed thethe
Lx ~ c +1.5integration
WK L x > c . region is set to be 0.1Rc ≤ R ≤ 0.6Rc andspeed of light, vy¢ = vy + 1.5WK where= min, DtWewherethe power spectra, mPm , for the magnetic
We set the time step as Dt|cos
θ|(D≤tgyro0.45.
cell ), plot
Dtgyro = Csafe Ep,0 (ceBmax ) and Dtcell = Csafe Dx c . Note thatthe velocities of CRs are always almost the speed of lightbecause we focus on the ultrarelativistic particles. We use themaximum value of the magnetic ﬁeld in the box Bmax toestimate Dtgyro , and set Csafe = 0.01.

Non-thermalbility. Panels (a) and (b) show the magnetic field lines (greenish
3.2. Resultsd X ¼ x=λ ¼ 1.91 (color contour), and panels (c)–(e) depict thethe same time stage. Panel (f): The energy spectra duringWe
theshow
MRIthe results of the simulations and discuss theevolution of the distribution function. The parameter sets are
Maxwellian fitting for T orbit ¼ 7.18.
tabulated in Table 2. The letters A, B, and C represent the

1-2

turbulent

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

Energy change

Figure 3. Evolution of the energy of a CR in the shear frame (thick solid line)
and the box frame (dashed line) for model A1. The thin solid and dotted lines
show the position x of the CR and the box boundary, respectively. The particle
energy jumps in the box frame, while it smoothly changes in the shear frame.
These jumps coincide with the CR’s crossing the box boundaries in the x
direction.

MNRAS 485, 163–178 (20Magnetic reconnection or MHD turbulence accelerates CRs
different times T of snapshot data, and we use the numbers 1, 2,
and 3 to distinguish between the initial energy  gyro. Group X
will be discussed in Section 4.2. We use Np = 215 = 32768 CR
particles and calculate their orbits until half of CRs escape fromthe system, the times of which are tabulated in Table 2 in units

see two jumps of energy at t ~ 240 tgyro,0 and t ~ 243 tgyro,0 ,which coincide with the CR’s crossing the box boundary(shown by the dotted lines) in the x direction. The CR position
is represented by the thin solid line. On the other hand, theenergy measured in the shear frame does not have such jumps

High-Energy Emission

13

Mahadevan et al. 1997; SSK et al. 2015

RIAFs in LLAGNs

Murase, SSK et al. 2019

Coronae in QSOs

 �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�

Hot#Plasma
BH

2

γ�

ν�

p

π

p

Turbulent#Field##!#proton#accelera4on

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

Fairly high photon & proton densities in Coronae & RIAFs
FIG. 2. Disk-corona SEDs and CR proton diﬀerential lumi• Interaction between CRs and matter/photons
nosities for L = 10 erg s , 10 erg s , 10 erg s ,10 erg s , 10 erg s (from bottom to top).
→ neutrino
gamma-rayemission
metricallythin, optically thick &disk [61].
The spectrum
•

42

X

45

is expected to have an exponential cutoﬀ at εdisk,cut ≈• B Tdisk , where Tdisk ≈ 0.49(GM Ṁ/16πσSBRS3 )1/4
2.8k
is the maximum eﬀective temperature of the disk(e.g., [62]). Here, M is the SMBH mass, Ṁ is the massaccretion rate, and RS = 2GM/c2 is the Schwarzschildradius.
Assuming a standard disk, we use Ṁ ≈

−1

46

−1

43

−1

44

−1

−1

TeV-PeV γ-rays are reprocessed
to MeV-GeV
γ raysTABLE I. Parameters
used in this work. Units
are [erg s ]−1for LX and Lbol , [M⊙ ] for M , [cm] for R, [cm−3 ] for np , and[%] for the ratio of the CR pressure to the thermal pressure.
log LX log Lbol log M ΓXθe

τT log R log np PCR /Pth

Index14•

IceCube Neutrinos•

Non-thermal Processes in Accretion Flows•AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226•

LLAGN RIAF model SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

•

Summary15

Basic Equations

MULTIMESSENGER TESTS OF
COSMIC-RAY ACCELERATION…
PHYS.
REV. D
MULTIMESSENGERTESTS OF COSMIC-RAY
ACCELERA

E2  [GeV cm-2 s-1 sr-1]

turbulence.
We computesteady stateCR
spectra by
solv- We plot EMISSION
proton-induced
cascadeV. CASCADEGAMMA-RAY
EMISSIONV.
CASCADEGAMMA-RAY
10-4
R
ing the following Fokker-Planck equation (e.g., [75–78]),Fig. 3. A sufficiently developed cascF• Stochastic
Acceleration(SA)
forprotons
RQ
AGN(c
Hadronuclear
andphotohadronic
processes
produce
Hadronuclear
and photohadroniccf.)…
Dermeretal.processes
2012, SSK
etproduce
al.Dthe2015!
"
MULTIMESSENGER TESTSOF COSMIC-RAY
ACCELERATION
PHYS.
REV.
100,
0830
a
flat
spectrum
below
criticaRa-5
3
very-high-energy
(VHE)
gamma
rays
throughneutralpion
10εvery-high-energy
(VHE)
gamma
rays
throughneutral
pion
1 ∂
∂F
F
∂Fp
R
p
p
p
attenuation
becomes ineffective. Thaε2p Dεpelectron/positron−= 2high-energy+decay andFppairs
+
Ḟ
,
p,inj
decay
and
throughRhigh-energy
electron/positronpairsthrough
∂t
ε
∂ε
∂ε
t
t
We plot
proton-induced cascade
gamma-ray
V. CASCADE pGAMMA-RAY
electron-positronpair productionise
p EMISSION
p−coolesc
total
charged pionp decay and the Bethe-Heitler
process.
The-6charged pion decay Fig.
and3. theBethe-Heitler
The
A sufficiently
developed
cascade10 process.
(1)
Z emissio
Hadronuclear
andphotohadronic
processes
produce
VHE
gamma
rays
are
absorbedbysoft
photons
throughthe
VHEfunction,
gamma raysare≈by softphotons
the at2 spectrum• Protona absorbed
flat
below
the throughcritical energy
Coolings:
whereis thegamma
CR distribution
D
ε
/t
ε
acc
þ F
− p (VHE)ppvery-high-energyraysthroughneutralpion
τγγ ðεγadditionalÞ≈
R optical
KðxÞdþ −γγ → e e process in the RIAF,γγand
produce
additional
attenuation
becomes
ineffective.
The
→
e
e
process
in
the
RIAF,
and
produce
−1
−1
0
+
-7
þ=
decay
and
pairsspace,throughp+γ→pelectron/positron(n)coeﬃcient+electron/positron
π (πpairs.
isthehigh-energy
diﬀusionin)energy
tp−cool10þ e−high-energyThehigh-energy
eelectron-positrone−tpp + pairs. The
pair
production is eestimatedthigh-energy
electron/positron
high-energy
−1
−1
−1
−1
−1
−1
charged
pion
decay
and
the
Bethe-Heitler
process.
The+ is- totaltpγalso+t→+tthecoolingrate,
tesc= tfall +tdiﬀ
pairs
emit
gamma-rays
throughsynchrotron
processes,
p+γ
p
+
e
+
e
p−syn
BH
pairs
also
emit
gamma-rays
throughprocesses, ener
Z gamma-ray
where synchrotron
εγ is theVHE gamma rays are absorbed by soft photons through thew-8dnγ
inverseComptonscattering,andbremsstrahlung,leadingtoistheescape
rate,
and
Ḟ
is
the
injectionfunction
10
þ
−
p,inj
−2
inverseComptonscattering,andbremsstrahlung,leading
to
KðxÞ
→p+
πsWe
γγ → ep+p
e processin p
the+RIAF,
and produce additionalðx − xτγγ ðεÞ lnðxÞHðx
− 1Þ,
x dε
¼ γε; γ εγð=γÞ ≈ R
reacceleration?
electromagnetic
cascades.
calculatethecascadeemisdε(see
Appendix[79]).
The
stochastic
acceleration
time
is
γ
electromagnetic
thestep
cascadeemishigh-energy
electron/positronpairs.
The
high-energy eþ2−q
e− cascades. Wethecalculate
Heavisidefunction[161].t
p
with
B
→
p
+
γ
(synchrotron)
2
sion
by
solvingthekinetic
equationsofphotons
and
-9
given
by
t
≈
η(c/V
)
(R/c)(ε
/eBR)
,
where
V
acc
A synchrotron
p processes,
A equations 10
pairs also emit gamma-rays
throughsion bysolving thewherekineticofphotons
andofthecriticalenergy,
ε
,
at
which
ε
is
the
gamma-rayenergy,
KðxÞ
γγ
-4
-3
-2
-1 ¼ 0oγ
electron/positron
pairs
(see
Refs.
[87,159,160]):
isthe Alfvénvelocityandη electron/positronisthe inverse
the turbu10
10
10
10
10
inverseComptonscattering,andbremsstrahlung,leading toofpairs
(see
Refs.
[87,159,160]):−2 can see flat spectra below the 2critica• Proton
Escape:ðx−xÞlnðxÞHðx
−
1Þ,
x
¼
ε
ε
=ðm
c
Þ,
anc
γ γ
e
lence strength[80,
We adopt
q = 5/3,
electromagnetic
cascades.
We81].
calculatethe cascadeemis-which is contabulated
values [161].
are approximate
theHeavisidestep function
We tabulate∂n+eεwith
Infallescapese andsion bysolving
the∂thekinetic
equationsphotonssistent
recentMHD ofsimulations[56],
togethere diffusive
FIG.
1.
Schematic
pictureoftheAGN
disk-corona scenario.te∂n∂½ðPIC þ Psyn þ Pff þ PCouεeÞnεe % Protons
þfitting
formula,ofcoronae
theenergy,εγγEGB, atwhilewhich
τγγby¼cascade
1 MRI
in T
e and theacceleratedbyturbulence
the
FIG.
3.
ENB
spectrelectron/positron(see
Refs.the[87,159,160]):½ðP
þ criticalPþPþ in thewith ∂tη = pairs
10.
dissipation rate
∂εBecausef
IC aresyn þ Pff
Cou Þngenerated
εe %
e
in
coronae,andproducehigh-energy
neutrinosandcascaded∂ε∂tformed
with
the
exact
cross
section.
• Electromagnetic
can
see
flat
spectrabelow
thecriticalenergy.
Nfore
The data are taken from Swiftf
e Cascade tois expected to benproportional
LXsecondaries, we assume
that
the
gamma
rays
via
interactions
with matter and radiation.
ε
ðγγÞ
inj
e
e
e
tabulatedvaluesareapproximately
calculateLATsensitivity
curve SMM
in the
high
g
loon
[87]
(blue),
[88]
(pu
n
∂n
¼∂ n_function
þ n_ εe ; scales asð31Þ
εeinjection
linearly
To explain
the
εe
inj
εe −
e LX . ðγγÞL
½ðPIC þtesc
Psyn þ Pff þ PCou Þnεe % ¼ n_ εe −44fittingþþ n_ formula,;
ð31Þ
the exposure
cascadecalculation
εe withFermi-LAT[90](orange),anda while10-yearobtained
−1
ENB,
the
CR
pressurerequired
for
L
=
10
erg
s
∂ε
∂t
t
X
e
w
esc
formed withtheandexact
cross aresection.
We
overplotupgoing
muon
track
(blu
predicted
fluxes
lower
than
the
e
turns
out to be γ∼inj 1% of the thermal pressure, which isγ
ðγγÞ γnεe
LATsensitivity
curve
in diskthehighgalactic
latitppmetricallythin,
optically
thick
[61].
The
spectrum
contributionofreaccelerated
−1
4
3
−1
the
cases.
The
Cherenkov
Telesco
∂n
n
n
_
_
¼
n
−
þ
n
;
ð31Þ
ε
ε
ε
ε
ε
γ /c
γ inj
γ+t
ðICÞ p ≡ðffÞðsynÞ)Fp V(tγγ ereasonable.
We)have anexposurepn
esc
isnexpectedtoexponential
cutoﬀ
atfromεdisk,cut≈tp−cool∂nn__− plotþ nε_ εpγLεþn_ ε4π(ε
þ
þ
n
;
ð32Þ
¼ e − tγesc
witha10-year
Ref.
ε
ε
ε
ε
ε
ðICÞ
ðffÞsensitivityðsynÞ obtainedγγγ
inj 30γγγ
better
above
GeV
tha
3
1/4
tescV is the volume. ¼ −in∂tFig. 2, twhere−2.8kpredicted
n_ ε,γ whereþ
n_ εTγdiskþ≈lowern_ εγ0.49(GMþ n_ εṀ/16πσ; sensitivity
ð32Þ
Tdisk
γγ
Bþ
SB RS ) cγ
fluxes
are
than
the
b
∂t
t
t
gamma-ray
is considerablysu
γγ interact
is escthe maximum
eﬀective flux
temperatureof the disk
γ
γ
γ
While
the
CRs
are
accelerated,
they
with
the[62]).
cases.
The
Cherenkov
Telescope
Array
(C
∂nεγ
nεγ nεγ
g
ðICÞ
ðffÞ
ðsynÞ
inj
(e.g.,
Here,
M
is
the
SMBH
mass,
Ṁ
is
the
mass
range
due
to
the
γγ
attenuation.
Fon_ ε þ n_ εmodeledþ n_ ε inþ n_theð32Þ
¼ − i −
matterand þradiationsection,ε ; previous

16Target Photon Field
Photon SED
High Lx•

Luminous objects→ Rich observational data
→ We can use empirical relation
based on observations•

Opt-UV photons from accretion disk•

X-rays from hot coronae above thin disk•

Softer spectra for higher Lx AGNs
Low proton cutofff energy for high Lx AGNs

Low LxPringle 1981, Ho 2008, Hopkins 2007
Bat AGN Spectroscopic Survey 2017, 2018,
Mayers et al. 2018

Proton SED

High Lx
Low Lx•scales as LX [see Eq. (22)], the flux starts to fall as E!
at a lower energy than for the lower luminosity AGN.

17Extragalactic γ & ν
Backgrounds

•
-4

10

E2  [GeV cm-2 s-1 sr-1]

denser x-ray target. At the same timebreak occurs at an increasingly lower emagnetic field increases as L1=2
X [seeB. Diffuse neutrino flux
effects are visible in Fig. 5.
The cumulative neutrino flux from all radio-quiet
In the case of variable MBH (top pan
AGNs in the Universe is obtained by convolution of theluminosity for a fixed accretion efficien
observed point source neutrino flux with the luminosity
by an increase of the black hole mass,function, taking into account its cosmological evolution,
disk and black hole region becomes larand photon density actually decrease as
Z LmaxZ zmax1dn
dV0effect of decreasing the magnetic fiel!diff!!;ob;
dLX
dz
f"z#!;ob"E!;ob# $4" Lmindz
dLX
0source and, hence, the synchrotron bre
energy as LX increases. However, the o(36)
AGNs with Lx ~1044 erg/s provide
thedominant
contribution
interactions
stays the same as MBH i3e.g.,XUedaet al. 2014
where dn0 =dL
describes
the present day x-ray lumiconsequence the proton energy at whic
nosity function of the sources and f"z# is its cosmologthe energy at which the neutrino flux sta
ical evolution. We have used the broken power-law
stays the same. The two effects are agai (this work)
luminosity andAGN
evolutionfunctions of Ref. [9] (model
panel of Fig. 5.
AGN (cascade)  (this work)I). In a Friedmann-Robertson-Walker
universe
thecoAGN (thermal e) Using Eq. (19), one sees that for low
2moving volume element is dV=dz $ 4"DL cjdt=dzj="1 %LX < 1041 erg s!1 , &X
p% < 1 for all proz# and the derivative of the cosmic time t with
reFIG. 1. Schematic picture
theAGN disk-corona
scenario.
Ebymax, andprotons
can escape from thepof
Protonsareaccelerated
turbulence
generated by the MRI•
We
choosetheinjection
efficiency!1high-energy neutrinos and cascadedz# &and produceing
spect to redshift z is "dt=dz# $ !H0 "1in% coronae,
totheobserved
cosmicray spect
gamma rays via interactions with matter and radiation.
p!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!!
FIG. 2. Disk-corona SEDs andmodel
can explain
quence,neutrinos
producedin the bl"1 % "m z#"1 % z#2 ! "# "2z % z2 #, where so
we thathave ournosities for L = 10 erg s
erg s , 10 erg s (fromlow luminosity AGNs 10would
in princi
used a standard #CDM cosmology with "m the$ 0:3MESE
andexcess.
metrically thin, opticallythick disk [61]. The spectrum
an exponential
cutoﬀ at ε≈
Waxman-Bahcall
bound: E2! dN"# $ 0:7. The neutrino energy in the observer’sis expected
frameto have the
2.8k T , where T≈ 0.49(GM Ṁ/16πσ)
TABLE I. Parameters used inand the source frame are related by E! $% maximumz#,
GeVreasonable:
cmof!2thesR!1
However,
10!8is theeﬀectivetemperature
disk [15].
• EEnergetically!;ob"1for L andL , [M ] for Mth
,

-5

10

10-6

42

X

45

10-7

B diskreacceleration?

-9
-4

10

-3

10

-2

10

-1

10

10

0

1

2

3

10 10 10
E [GeV]

4

10

10

5

6

10

•

10

7

46

−

−1

disk,cut
3 1/4
S

SB

X

-8

10

10

disk

−1

(e.g., [62]). Here, M is the SMBH mass, Ṁ is the massaccretion rate, and RS = 2GM/c2 is the Schwarzschildradius.
Assuming a standard disk, we use Ṁ ≈Lbol /(ηrad c2 ) with a radiative eﬃciency of ηrad = 0.1.
Although the spectra calculated by Ref. [60] extend to
low energies, we only consider photons with εdisk > 2 eV
because infrared photons would come from a dust torus.
X rays are produced via Compton upscattering by
thermal electrons with Te ∼ 109 K. The spectrum canbe modeled by a power law with an exponential cutoﬀ.
The photon index, ΓX , is correlated with the Eddington ratio as ΓX ≈ 0.167 × log(λEdd ) + 2.0 [63]. The
cutoﬀ energy is also given by εX,cut ∼ [−74 log(λEdd ) +bol⊙[%] for the ratio of the CR pres

PCR/Pth ~ PCR/PB ~ 0.01
123001-8
Cascadeemission provides 10 - 30
% of MeV γ-ray backgroundlog LX log Lbol log M ΓX θ
42.0
43.0
6.51 1.72 0.2
43.0
44.2
7.25 1.80 0.2
44.0
45.4
8.00 1.88 0.2
45.0
46.6
8.75 1.96 0.1
46.0
47.9
9.49 2.06 0.1

LX , β and r are the only reare also constrained in ations [68, 69] and numerica

18

HE particles from Nearby AGNs
A typical Seyfert at 30 Mpc

10-410-6

AGN (thermale)  
CascadeThermal-HWWHG AGN (in clusters) e 
-HWWHG AGN (w. blazars) 

10-7

reacceleration?

-8

10 10

-2 -1

E2  [GeV cm-2 s-1 sr-1]
E FE [GeV cm-2 s-1]

NGC 1068
10-6

LX=10 erg/s
AGN  (this work) +e
d=30 Mpc AGN (cascade)Generated
 (this work)

IceCube (=30o)eASTROGAM

-7

•

44

10-510-7-64

E FE [GeV cm s ]

•

total

Bethe-Heitler

-9

10 10

AMEGO

pp+p
-8 -10Generated +e
Cascade 
Thermal e 

NGC 1068
LX=1043 erg/s
d=12.7 Mpc
eASTROGAM

Bethe-Heitler-8

10

IceCubereacceleration?

pp+p
AMEGO

-9

10

Fermi LAT

starburst(MW16)

-10

10 10

10
reacceleration?

-11
10-910 -4 -4-3 -3-2 -2-1 -10 0 1 1 2 2 3 3 4 4 5 5 6 6 7 7
10 10 10 10 10 10 10 10 10 10 10 10

10

10

10

10

10

10 10 10
E [GeV]E [GeV]

10

10

10

・pγ neutrinos are detectable by

10

FIG.
Pointandsourcefluxesof allfrom
flavorAGNs,
neutrinos
and gamma
G.
S5.4. EGB
ENB
spectrain whichconaysfromfroma nearbya standard
disk.are
A possible
utionsboth AGN
AGNwithcoronaeand jets
included.efIceCube-Gen2of secondary
reaccelerationis [48]indicated(thinsolid).balFor
eect
data
are taken from
Swift-BAT(green),Nagoya[102] andAMEGO COMPTEL[103] sensitivities,
the obnASTROGAM[49] (blue), SMM[50] (purple),[51] (gray),
6・MeV
γ-rayscanbedetectedbyervation [52]time(orange),of 10 s andis assumed.
Theshower
IceCubeeight-year
mi-LAT
IceCube for(black)[43]
ensitivity
[104]
and
the
5
times
better
case
[105]
are
shown.
upgoing muon track (blue shaded) [44] events. A possifutureofsatellites.
contribution
reaccelerated electron and positron pairs)t are hadronically injected is indicated (red thin solid).
e MeV-TeV EGB and TeV-PeV ENB can universally be

10-11 -4 -3 -2 -1 0
10 10 10 10 10 101 102 103 104 105 106 1
E [GeV]

tofluxes
explain
IceCubedata
FIG.・Possible
S6. Point sourceof all flavor
neutrinos
and gamrays from a nearby Seyfert galaxy, NGC 1068. A possibl
fect ・Consistentof secondary reacceleration
is indicatedwith Fermi
data(thin solid).
eASTROGAM [60] and AMEGO [61] sensitivities, theservation time of 106 s is assumed. The ten-year IceC
due to γγ pair-production
data [58] and the Fermi gamma-ray data [55] are shown.
CR pressure is set to 30% of the thermal pressure, and ⌘is adopted. For comparison, the prediction of the starb
scenario presented in Murase and Waxman [31] is shown

19

HE particles from Nearby AGNs

E2  [GeV cm-2 s-1 sr-1]
E FE [GeV cm-2 s-1]

10-510-7
-8

10 10
-7

•

NGC 1068
10-6

LX=10 erg/s
AGN  (this work) +e
d=30 Mpc AGN (cascade)Generated
 (this work)
44

AGN (thermale)  
CascadeThermal-HWWHG AGN (in clusters) e 
-HWWHG AGN (w. blazars) 

10-7

reacceleration?

IceCube (=30o)eASTROGAM

-6

4

total

Bethe-Heitler

-9

10 10

AMEGO

pp+p
-8 -10-2 -1

10-410-6

A typical Seyfert at 30 Mpc

E FE [GeV cm s ]•Generated +e
Cascade 
Thermal e 

NGC 1068
LX=1043 erg/s
d=12.7 Mpc
eASTROGAM

Bethe-Heitler-8

10

IceCubereacceleration?

pp+p
AMEGO

-9

10

Fermi LAT

starburst(MW16)

-10

10 10

10
reacceleration?

10-11 photon
Ratio
of
γ
to
ν
flux
is
fixed
by
the
observedfield
10-4 10-3 10-2 10-1 100 101 102 103 104 105 106 1
10 10 10 10 10 10 10 10 10 10 10 10E [GeV]E [GeV]E [GeV]→ We can robustly
test our model by future experiments-1110-910 -4 -4-3 -3-2 -2-1 -10 0 1 1 2 2 3 3 4 4 5 5 6 6 7 7
10 10 10 10 10 10 10 10 10 10 10 10•・pγ neutrinos are detectable by

FIG.
Pointandsourcefluxesof allfrom
flavorAGNs,
neutrinos
and gamma
G.
S5.4. EGB
ENB
spectrain whichconaysfromfroma nearbya standard
disk.are
A possible
utionsboth AGN
AGNwithcoronaeand jets
included.efIceCube-Gen2of secondary
reaccelerationis [48]indicated(thinsolid).balFor
eect
data
are taken from
Swift-BAT(green),Nagoya[102] andAMEGO COMPTEL[103] sensitivities,
the obnASTROGAM[49] (blue), SMM[50] (purple),[51] (gray),
6・MeV
γ-rayscanbedetectedbyervation [52]time(orange),of 10 s andis assumed.
Theshower
IceCubeeight-year
mi-LAT
IceCube for(black)[43]
ensitivity
[104]
and
the
5
times
better
case
[105]
are
shown.
upgoing muon track (blue shaded) [44] events. A possifutureofsatellites.
contribution
reaccelerated electron and positron pairs)t are hadronically injected is indicated (red thin solid).
e MeV-TeV EGB and TeV-PeV ENB can universally be

tofluxes
explain
IceCubedata
FIG.・Possible
S6. Point sourceof all flavor
neutrinos
and gamrays from a nearby Seyfert galaxy, NGC 1068. A possibl
fect ・Consistentof secondary reacceleration
is indicatedwith Fermi
data(thin solid).
eASTROGAM [60] and AMEGO [61] sensitivities, theservation time of 106 s is assumed. The ten-year IceC
due to γγ pair-production
data [58] and the Fermi gamma-ray data [55] are shown.
CR pressure is set to 30% of the thermal pressure, and ⌘is adopted. For comparison, the prediction of the starb
scenario presented in Murase and Waxman [31] is shownIndex20•

IceCube Neutrinos•

Non-thermal Processes in Accretion Flows•AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226•

LLAGN RIAF model SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

•

Summary21

Basic Equations

MULTIMESSENGER TESTS OF
COSMIC-RAY ACCELERATION…
PHYS.
REV. D
MULTIMESSENGERTESTS OF COSMIC-RAY
ACCELERA

E2  [GeV cm-2 s-1 sr-1]

turbulence.
We computesteady stateCR
spectra by
solv- We plot EMISSION
proton-induced
cascadeV. CASCADEGAMMA-RAY
EMISSIONV.
CASCADEGAMMA-RAY
10-4
R
ing the following Fokker-Planck equation (e.g., [75–78]),Fig. 3. A sufficiently developed cascF• Stochastic
Acceleration(SA)
forprotons
RQ
AGN(c
Hadronuclear
andphotohadronic
processes
produce
Hadronuclear
and photohadroniccf.)…
Dermeretal.processes
2012, SSK
etproduce
al.Dthe2015!
"
MULTIMESSENGER TESTSOF COSMIC-RAY
ACCELERATION
PHYS.
REV.
100,
0830
a
flat
spectrum
below
criticaRa-5
3
very-high-energy
(VHE)
gamma
rays
throughneutralpion
10εvery-high-energy
(VHE)
gamma
rays
throughneutral
pion
1 ∂
∂F
F
∂Fp
R
p
p
p
attenuation
becomes ineffective. Thaε2p Dεpelectron/positron−= 2high-energy+decay andFppairs
+
Ḟ
,
p,inj
decay
and
throughRhigh-energy
electron/positronpairsthrough
∂t
ε
∂ε
∂ε
t
t
We plot
proton-induced cascade
gamma-ray
V. CASCADE pGAMMA-RAY
electron-positronpair productionise
p EMISSION
p−coolesc
total
charged pionp decay and the Bethe-Heitler
process.
The-6charged pion decay Fig.
and3. theBethe-Heitler
The
A sufficiently
developed
cascade10 process.
(1)
Z emissio
Hadronuclear
andphotohadronic
processes
produce
VHE
gamma
rays
are
absorbedbysoft
photons
throughthe
VHEfunction,
gamma raysare≈by softphotons
the at2
• Proton� γ�througha absorbed
flat
spectrumbelow
thecritical energy
Coolings:
whereis thegamma
CR distribution
D
ε
/t
ε
acc
þ F
− p (VHE)ppvery-high-energyraysthroughneutralpion0±�τγγ ðεγadditional
Þ≈
KðxÞd
þ −
p+p RIAF,! p+p+Aπ+Bπγγ → e e process in the RIAF,γγand
produce
additionalν� R opticalattenuation
becomes
ineffective.
The
→
e
e
process
intheandproduce
−1
−1
0
+
-7
0�
p
þ=
decay
and
pairsspace,through! p+p+πp+γ→pelectron/positron(n)coeﬃcient+electron/positron
π (πpairs.
isthehigh-energy
diﬀusionin)energy
tp−cool
tpp +
10
þ e−
high-energyThehigh-energy
eelectron-positron
e−p+γ
pair
production is eestimatedt+�
high-energy
electron/positronpairs.
The
high-energy
p+γ
!
p+n+π
−1
−1
−1
−1
−1
−1
charged
pion
decay
and
the
Bethe-Heitler
process.
The+ is- totaltpγalso+t→+tthecoolingrate,
tesc= tfall +t
±�
pairs
emit
gamma-rays
throughsynchrotron
processes,
p+γ
p
+
e
+
e
p−syn
π0 !
2γ,
π± !
BH
diﬀ
pairs
also
emit
gamma-rays
throughsynchrotron
processes, ener
Z gamma-raywhere
εγ3ν+eis theVHE gamma rays are absorbed by soft photons through thew-8dnγ
inverseComptonscattering,andbremsstrahlung,leadingtoistheescape
rate,
and
Ḟ
is
the
injectionfunction10
þ−Hot#Plasmap,inj−2
inverseComptonscattering,andbremsstrahlung,leadingto→p+πsWeγγ → ep+p
e processin p
the+RIAF,
and produce additionalðx − xτγγ ðεÞ lnðxÞHðx
− 1Þ,
x dε
¼ γε; γ εγð=γ Þ ≈ R π KðxÞreacceleration?
electromagnetic
cascades.
calculatethecascadeemisdε(see
Appendix[79]).
The
stochastic
acceleration
time
is
γ
electromagnetic
thestep
cascadeemishigh-energy
electron/positronpairs.
The
high-energy eþ2−q
e− cascades. Wethecalculate
Heavisidefunction[161].t
p
with
B
→
p
+
γ
(synchrotron)
2
sion
by
solvingthekinetic
equationsofphotons
and
-9
given
by
t
≈
η(c/V
)
(R/c)(ε
/eBR)
,
where
V
acc
A synchrotron
p processes,
A equations 10
pairs also emit gamma-rays
throughsion bysolving thewherekineticofphotons
and
BH
of
the
criticalenergy,
ε
,
at
which
ε
is
the
gamma-rayenergy,
KðxÞ
γγ
-4
-3
-2
-1 ¼ 0oγ
electron/positron
pairs
(see
Refs.
[87,159,160]):
p 10 Turbulent#Field##isthe Alfvénvelocityandη electron/positronisthe inverse
the turbu10
10
10
10
inverseComptonscattering,andbremsstrahlung,leading toofpairs
(see
Refs.
[87,159,160]):−2 can see flat spectra
2critica
below
the• Proton!#proton#accelera4on
Escape:ðx−xÞlnðxÞHðx
−
1Þ,
x
¼
ε
ε
=ðm
c
Þ,
anc
γ γ
e
lence strength[80,
We adopt
q = 5/3,
electromagnetic
cascades.
We81].
calculatethe cascadeemis-which is contabulated
values [161].
are approximate
theHeavisidestep function
We tabulatet∂neεwith∂ kinetic
Infalle andsion bysolving
equationsphotonssistent
recentMHD ofsimulations[56],
togethere the thee∂n∂½ðPIC þ Psyn þ Pff þ PCouεeÞnεe %þfitting
formula,ofcoronae
the criticalenergy,εγγEGB, atwhilewhich
τγγ ¼cascade
1 in Te
FIG.
ENB
spectrelectron/positron(see
Refs.the[87,159,160]):
½ðP
P
% the
þ in thewith ∂tη = pairs
10.
dissipation rate
∂εBecausef
IC þ Psyn þ P
ff þ 3.
Cou Þnεeande∂ε∂tformed
with
the
exact
cross
section.
• Electromagnetic
can
see
flat
spectrabelow
thecriticalenergy.
Nfore
The data are taken from Swiftf
e Cascade tois expected to benproportional
LXsecondaries, we assume
that the
ε
ðγγÞ
e
e
e
tabulatedvaluesare[87]
approximately
calculateLATsensitivity
curve SMM
in thehighgloon(blue),
[88]
(pu
nð31Þ
∂n
¼∂ n_function
þ n_ injεe
injection
linearly
To explain
the
εe
inj
εe ; scales asεe −e LX . ðγγÞL
½ðPIC þtesc
Psyn þ Pff þ PCou Þnεe % ¼ n_ εe −44fittingþþ n_ formula,;
ð31Þ
the exposure
cascadecalculation
εe withFermi-LAT[90](orange),anda while10-yearobtained
−1
ENB,
the
CR
pressurerequired
for
L
=
10
erg
s
∂ε
∂t
t
X
e
w
esc
formed withtheandexact
cross aresection.
We
overplotupgoing
muon
track
(blu
predicted
fluxes
lower
than
the
e
turns
out to be γ∼inj 1% of the thermal pressure, which isγ
ðγγÞ γnεe
LATsensitivity
curveThein thehighgalactic
latitppcontributionofreaccelerated
−1
4
3
−1
the
cases.
Cherenkov
Telesco
∂n
n
n
_
_
¼
n
−
þ
n
;
ð31Þ
ε
ε
ε
ε
ε
γ /c
γ inj
γ+t
ðICÞ p ≡ðffÞðsynÞ)Fp V(tγγ ereasonable.
We
)
pn
esc
p−coolt∂nnn__− plotþ nε_ εpγLεþn_ ε4π(ε
þ
þ
n
;
ð32Þ
¼ e − tγesc
witha10-yearexposure
from
Ref.
ε
ε
ε
ε
ε
ðICÞ
ðffÞsensitivityðsynÞ obtainedγγγ
inj 30γγγ
better
above
GeV
tha
tescV is the volume. ¼ −in∂tFig. 2, twhere− predictedþ n_ εγ þ
n_ εγ are
þ lowern_ εγ thanþ n_ εthe; sensitivity
ð32Þ
γγ
γ
fluxescb∂ttt
gamma-rayflux
isconsiderably
su
γγ
esc
γ
γ
γ
the
CRs are accelerated, they interact
with The Cherenkov Telescope Array (Cthe cases.
∂nεγ While
nεγ n
εγ
g
ðICÞ
ðffÞ
ðsynÞ
inj
n_ ε þ n_ εmodeledþ n_ ε inþ n_theð32Þ
¼ − i −
matterand þradiationsection,range due to the γγ attenuation. Fo
ε ; previous

We consider a RIAF of size R andvicinity of supermassive black holes (SMBHs) in active
rate Ṁ around a SMBH of mass MBH .
galactic
nuclei (AGNs) [72–78]. A luminous AGN hosts
22
tation Qx = 10x in cgs units, unless oa geometrically thin, optically thick accretion disk that
To represent the physical quantities in
produces copious UV photons [79–81], and the ratio ofconvenient to normalize R by the Schwa
the observed UV to X-ray luminosity is very high [82–
R = RRS ' 2.95 ⇥ 1014 R1 M8 , where R84]. Such target photon fields lead to a hard neutrino
is the Schwarzschild radius, G is the gr
spectrum at PeV energies [85, 86]. The accretion shock
stant, and c is the speed of light. T
has been considered, but the existence of such a shockHighaccretion rate
PhotonSED bytion rate is normalized by the Eddington
has not been
supportednumerical
simulations so far.
ṁ = Ṁ c2 /LEdd , where LEdd ' 1.3 ⇥ 10
On the other hand, recent studies on magnetorotational
the Eddington luminosity.
nstabilities Synchrotron
suggest that particle acceleration via magAccording to recent magnetohydrod
netic reconnections and turbulence is promising in AGN
simulations (see e.g., [94–99]), the radicoronae, and Ref. [87] showed that the mysterious 10 –sound velocity, the scale height, the num
100 TeV component in the di↵use neutrino flux can be exInverse
magnetic field, and the Alfven velocity iplained by the AGN core model of radio-quiet
AGNs. It
Comptonestimated to be
was found that the Bethe-Heitler process is critically important, which led to robust predictions of MeV gamma
1
1/2
8
↵V
'
3.4
⇥
10
R
↵ 1 cm sV⇡rays via proton-induced
cascades.
R
K
1
Low accretion rate2
Low-luminosity AGNs (LLAGNs), however, have dif1
1/2
C⇡
VK ' 3.4 ⇥ 109 R1
cm s 1serent spectral energy distributions, in which an UV
2bump is absent [88]. This indicates that there is an op1•
Low-luminosity→
Poor
observationaldata
H⇡
R ' 1.5 ⇥ 1014 R1 M8 cm
tically thin, hot accretion flow instead of an optically2
thick disk. →Remarkably,plasmaproperties
of
hot
AGN
Formulation based on theoryṀ
3/2
coronae and radiatively inefficient accretion flows (RIAF;
Hot#Plasma
np ⇡
' 4.6
⇥ 108 R1 ↵ 11
electronsin the
RIAFs
seed photons
4⇡mp RHVR
89, 90])• in Thermal
LLAGNs seem
similar insenseemitthat thesBHplasmas •areOur
expectedtobecollisionless
for
ions.
It
is
results are consistent with X-ray observations
8⇡Pg
5/4
1/2
natural to consider the same type of proton acceleration
B⇡' 2.6 ⇥ 102 R1 ↵ 1 M8n both Seyfert galaxies and LLAGNs. Ref. [91] considered the stochastic acceleration expected in such RIAFs
B
1/2
1/2
V ⇡p
' 2.7 ⇥ 109 R

Target Photon Field �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�

γ�

ν�

p

π

p

Turbulent#Field##!#proton#accelera4on

8.23
8.68
23
9.14

159.56 -1.48 1.79
268.77 -1.02 1.30
452.72 -0.57 0.91

> 1.5 ⇥ 104 GeV
eV for ṁ ⇠ 10 3 ,um appears at aectral index andls B and C, theentical to the inminates over the
nge.

4.04
3.25
2.14

41.15/40.97/41.7 13.9/9.3/48.4
41.60/41.43/42.1 11.3/7.2/41.1
42.05/41.88/42.6 7.7/4.1/28.6

Extragalactic γ & ν
Backgrounds

26, 67])

c
=
i
4⇡H0

Z

dz

Z

 �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�

Hot#Plasma

γ�

ν�

p

π

BH
L"i ⌧i,IGM
p
dLH↵ ⇢H↵e,3"(1 + z) ⌦m + ⌦⇤i(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
can explainH↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡PeV
ν
and⇥ MeV γ
s1s2(⇢⇤ /L⇤ )/[(L
/L
)
+
(L
/L
)
],
where
⇢
'
4.11
H↵
⇤
H↵
⇤
⇤
Preliminaryeson interactions
5
3
41
1
10
Mpc , L⇤ = 3.26 ⇥ 10 erg sobs.
, s1 simultaneously= 2.78,
We calculate theand s2 = 1.88. We extrapolate this luminosity func• GeV γs are attenuatedng the formalism
tion to Lmin = 1038 erg s 1 , below which
the Palos by p interacRIAFs
LLAGNsmar survey finds a hint of a flatteningat[69].
The insuron given in Ref.
vey also indicates a correlation between →LX consistentand LH↵ for with Fermi datatic field strength
LLAGNs: LX ⇡ 5 7LH↵ [69]. We use a correctione e↵ect of meson• PCR
0.01Pthfactor X/H↵ = LX /LH↵ = 6.0. Then,the ~luminosity
neutrinos. Then,integration is performed in the range of 1038 erg s 1 ) = (1, 2, 0) at
LH↵  ⌘rad ṁLEdd /(X/H↵ bol/X ) ' 4.2 ⇥ 1041 erg s 1 .
o the neutrino osSince dimmer AGNs tend to have weaker redshift evolu• With a similarparameter set, Seyfert contribute to TeV-PeV ν
onic interactions
tion [70–72], we assume no redshift evolution of the lun/positron→pairs,AGN corescanaccount
a broadenergyminosity
function.
Theformassof SMBHs
in localrange
Seyfertof ν backgrounds. We calculate
galaxies does not show any correlation with X-ray luetic equations ofpTurbulent#Field##!#proton#accelera4onthen they can be as bright as EEs owing to the highproduction efﬁciency. Also, neutrino ﬂuences from pro
emission can be higher than the plateau and ﬂares if G 1 30
realized.

Detectability

24

3. Probability of Neutrino Detection•

nm -induced
The •expected
number offor
eventsis estimated t
Prospects
stacking
analysis

NGC 4565&m =Thermal γpò fn Aeff (d, En ) dEn ,where Aeff is the effective area. The effective areas of upg+horizontal and downgoing tracks for IceCube is show
Aartsen et al. (2017)IceCube-Gen2
as a function of Eν. For upg

Cascade γ

Preliminaryνμ
IceCube••Cascade γ-rays @ E<1 GeV: too
faint to detect by LAT or CTA
Thermal γ-rays @ E~MeV:
detectable by future MeV satellites•••

Atmospheric backgroundis negligible @ E > 100 TeV
IceCube cannot detect νs
IceCube-Gen2 will detect νsIndex25•

IceCube Neutrinos•

Non-thermal Processes in Accretion Flows•AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226•

LLAGN RIAF model SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

•

Summary26

Summary••••

Accretion flows in AGNs are a feasible neutrino source
Seyfert galaxies can reproduce TeV-PeV ν backgroundwithout violating Fermi constraints with observation based parameters.
RIAFs in LLAGNs can explain MeV γ &PeV ν backgrounds
simultaneously without overshooting Fermi constraints
Future multi-messenger observations can robustly test both models:- IceCube-Gen2 can resolve AGNs as point sources- Proposed MeV satellites can detect MeV γ rays from AGNs

Preliminary �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�

Hot#Plasma
BH

p

γ�

ν�

p

π

Turbulent#Field##!#proton#accelera4on27

Thank youfor your attention

Stochastic Acceleration
3.4. Evolution
momentum・Consider
plasma inwithturbulentspace
fieldsdN/dpi を図 8 示
E0E0 次に各方向の運動量の数スペクトルす。点線が初期の数スペクトル、実践が計算終了時の数ス＜〜〜〜
ペクトルである。初期に等方な速度分布を持った粒子群は
particlewave
particlewave
計算終了時もほぼ等方のままである。計算途中もずっと等方であり、時間とともに各方向の分散が同時に大きくなっ
E’ > E0
E’
E’< E0ていく。
E’
今回の実験で得られた分布関数の時間進化は、等方な運Some gain E, others lose E →diffusion in E space
動量空間の拡散方程式でよく記述される。
＜〜〜〜

28

相互作用によりエネルギーを得ている。一方、下図では粒子が境界を横切る度に階段的にエネルギーを増やしているのが見て取れる。こちらはローレンツ変換のせいでエネル
ギーを増やしている。今回の実験では電磁場によるエネル
ギー増加が支配的であると考えられる(議論の項を参照。 et al. 2015e.g.) Fermi 1949, Stawarz& Petrosian 2008, )
SSK

!
"
∂f
1 ∂
∂f
2
= 2
p Dp
+ f˙inj
∂t
p ∂p
∂p

(24)

これは一回の散乱でエネルギーの変化が少なく、多数回のAccordingto quasi-linear theory,
散乱を統計的に扱う見方が良い場合に有効な式である。こgyro-resonant
scattering resultsinの方程式を Fokker-Planck
Equationと呼ぶこともある。
ここで p Dは運動量の大きさであり、超相対論的を仮定してp ∝ pq (power-spectrum Pk = P0k-q)
p = γmc2 という関係で表される。今回の実験では、初期に同じエネルギーを持った粒子を用意した後は粒子の注入

Fi
布は

を行

と表
分
でき
Dp

の形
とき

T
29

Neutrino Production Processw
S. R. KELNER, F. A. AHARONIAN, AND V. V. BUGAYOVsomby• pp inelastic collision

σ[mb]

νμ
μ-

π-

target
HEP

p

p

νμ
e-

νe

γ

π0

γ

νμ

μ+

π+

e+
νe

νμ

Kelner 06

Neutrino spectral shapeis similar to that of protons
inel = (34 .3 + 1 .88 L + 0 .25 L
4 2
E th

2

T
th

)×

1−
, mb
Ep
FIG. 8. The multiplicity of photons and leptons produced
inEp’[TeV]
one interaction of a relativistic proton with 2.7 CMBR.

th

T

σ[mb]
• Photomeson production (pγ) FIG. 11.
Inelastic cross section of p-p interactions approxii.e

HEP

p

＜
〜
〜

target γ

Kelnerdata
08 are frommated by Eq. (79). The experimentalth
http:wwwppds.ihep.su:8001/c5-5A.HTML, the open points cor
respond to the cross sections which are used in the SIBYLL code
pr

γ
γ
γγ

π0

ν
νμμ
π+

μ+
νμ

e+
ννee

!
" !
"
cnH
E!
E!
q! !E! " # n~
"
m $
J m $
;
K! inel p K! p p K!

(77)

w
n. For the procedure described below, un~where K! # #=~p
and K! are free parameters.
"E
The emissivity of gamma rays is related to q! !Eas!"
E’[GeV]γ
though the equationphZ 1 q! !E! "
in
dN$
#2
(78)
p#################### dE! ;
ac

T
30

Neutrino Production Processw
S. R. KELNER, F. A. AHARONIAN, AND V. V. BUGAYOVsomby• pp inelastic collision

σ[mb]

νμ
μ-

π-

target
HEP

p

p

νμ
e-

νe

γ

π0

γ

e+
νe

νμ

Neutrino spectral shapeis similar to that of protons

νμ

μ+

π+

Kelner 06

inel = (34 .3 + 1 .88 L + 0 .25 L
4 2
E th

2

T
th

)×

1−
, mb
Ep
FIG. 8. The multiplicity of photons and leptons produced
inEp’[TeV]
one interaction of a relativistic proton with 2.7 CMBR.

th

T

σ[mb]
• Photomeson production (pγ) FIG. 11.
Inelastic cross section of p-p interactions approxii.e

HEP

p

＜
〜
〜

target γ

Kelnerdata
08 are frommated by Eq. (79). The experimentalth
http:wwwppds.ihep.su:8001/c5-5A.HTML, the open points cor
respondpto the cross sections which areγused in the SIBYLL code
pr

γ
γ
γγ

π0

ν
νμμ
π+

μ+
νμ

e+
ννee

E ~10 TeV �����E ~ 30 keV
Ep~10 PeV!
Eγ"~!30 eV"
cnH
EE"inel mp $ ! Jp mp $ ! ;K!
K!
K!
Neutrino
spectralshape

q! !E! " # n~(77)

w
n. For the procedure described below, un~where K! # #=~depends on photon spectrum p
and K! are free parameters.
"E
The emissivity of gamma rays is related to q! !Eas!"
E’[GeV]γ
though the equationphZ 1 q! !E! "
in
dN$
#2
(78)
p#################### dE! ;ac31

Rates & CR Spectrum
Lx=1044 erg/s•

Ep,max ~ 105 GeV by tacc = tBH•

BH suppresses ν production
at Ep ~ 3x104 - 3x106 GeV
Escape is inefficient

•

•
•
•

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

Hard spectrum due to SA
Pile up around Emax

metrically thin, optically thick disk [61]. The spectrum
is expected to have an exponential cutoﬀ at εdisk,cut ≈2.8kB Tdisk , where Tdisk ≈ 0.49(GM Ṁ/16πσSB RS3 )1/4
is the maximum eﬀective temperature of the disk(e.g., [62]). Here, M is the SMBH mass, Ṁ is the massaccretion rate, and RS = 2GM/c2 is the Schwarzschildradius.
Assuming a standard disk, we use Ṁ ≈Lbol /(ηrad c2 ) with a radiative eﬃciency of ηrad = 0.1.
Although the spectra calculated by Ref. [60] extend to
low energies, we only consider photons with εdisk > 2 eV
because infrared photons would come from a dust torus.
X rays are produced via Compton upscattering by
thermal electrons with Te ∼ 109 K. The spectrum canbe modeled by a power law with an exponential cutoﬀ.

Higher Lx → lower Emax
because of efficient cooling

FIG. 2. Disknosities for LX
1045 erg s−1 , 1

TABLE I. Par
for LX and Lb[%] for the rat

log LX log Lb
42.0
43.0
43.0
44.2
44.0
45.4
45.0
46.6
46.0
47.9

LX , β and r

32

Rates & H-E Spectrum
1

Infall is dominant• Neutrino is mainly
produced by pp• For SA, acceleration rate is much
lower than PL model•

celere thep,inj ,

(S1)
prooeffihe esming3 �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�

Hot#Plasma
BH

(S2)

•

⌘ =
s the
w inr the
nying

•

rates
and
r the

p

γ�

ν�

p

π

Turbulent#Field##!#proton#accelera4on

Hard proton spectrum for SA
Cutoff energy for SA is muchhigher than the critical energy(tacc = tfall) due to hard spectrum& gradual cutoff

33

Plasma Conditionsin Accretion Flow
•
•
•
•
•

To accelerate non-thermal particles,
relaxation time > dissipation time
For RIAFs in LLAGNs, tdis ~ tdyn ~ R/Vfall
For Coronae in QSO, tdis ~ H/VA
Protons are collisionless for both cases→ Non-thermal Proton
Electrons are collisional for both cases→ Thermal electrons only

34

Implications & Caveats•

Multi-messenger tests are promising:
Nearby LLAGNs are detectable by IC-Gen2 & e-ASTROGAM•High source number density (~10-3 Mpc-3)→ LLAGNs can avoid the point-source constraints•

Luminosity Function (LF) is very uncertain- LF by Hao et al. (2005) >> Greene & Ho (2007)- If we use Greene & Ho (2007), neutrino flux becomes too low
to explain TeV-PeV neutrinos