Hadronic Gamma Rays fromthe Hot Accretion Flow in M87
FRIS, Tohoku University(JSPS Research Fellow)

Shigeo S. Kimura
References
SSK, Toma in preparation
see also:
SSK, Murase, Toma, 2015, ApJ, 806, 159
SSK, Murase, Meszaros in preparation
SSK, Toma, Suzuki, Inutsuka, 2016, ApJ, 822, 88
SSK, Tomida, Murase, 2019, MNRAS, 485, 163

Collaborators
K. Toma (Tohoku Univ.); K. Murase (Penn State; YITP)
P. Meszaros (Penn State); K. Tomida (Osaka Univ.)
T.K. Suzuki (Univ. Tokyo); S.-i.Inutsuka (Nagoya Univ.)

AGN Jet Workshop 2020/FRIS International Workshop 2020@SendaiContents

2•

Motivation for hadronic emission from RIAFs- Particle Acceleration in RIAFs- RIAFs as Neutrino Sources•

Hadronic Emission from MADs in Radio Galaxies•

SummaryM87 Observations

3

8

V. A. Acciari et. al.

E2 dN/dE [ergs cm-2 s-1]

−5

10

−2

10

4

10

10

10

7

10

10

Energy [eV]1013 1015

10−11 MAGIC Collaboration & Asano, Hada et al. 2020

10−12

10−13

M87 core data
Leptonic model, SSC
Hybrid model

How about the hot accretion flow?
10−14

108

1011

1014

1017

1020

1023

1026

1029

Frequency, ν [Hz]HE & VHE γ-rays are observed, butemission region & mechanism are not known• Leptonic jet model: too weak magnetic field (mG)et al. 2014, neglecting Fermi-II acceleration (see also, Asano &the electron injection spectrum are:45
the minimum and maximum•
Hadronicjetmodel:
very
high
power
(>10
erg/s)
Hayashida 2015, 2018). The code calculates the temporal evolution
electron Lorentz factors (and), thelocation of the break•

Figure 4. MWL SED of the radio core of M 87 compiled from quasi-simultaneous 2012–2015 observations (black points). VHE gamma-ray observations by
MAGIC are combined with HE gamma-ray data from Fermi-LAT, X-ray data from Chandra, NUV data from HST, radio data at 1.7 GHz and 5 GHz provided
by the EVN and at 43 GHz by VLBA. The models represent two possible scenarios: in the leptonic scenario (red solid line) the high-energy component isdominated by the SSC emission, whereas in the hybrid scenario (blue dashed line) the high energy emission is dominated by the synchrotron radiation ofrelativistic protons.

of the electron and photon energy distributions in the plasma rest

min

max

in the electron energy distribution ( br ), and power-law indices p1

Contents

4•

Motivation for hadronic emission from RIAFs- Particle Acceleration in RIAFs- RIAFs as Neutrino Sources•

Hadronic Emission from MADs in Radio
Galaxies•

SummaryAccretion flows in AGNs
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
10

11

12

13

14

QSO UV X-ray log (ν/Hz )

15

16

17

Opt.
~~
~~
~~
>

9

~>

38

~~
~~
~~
>
~~
~

log (νLν /ergs s–1 )Ho 2008

ure 7

10 keV

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

18

LLAGN

mposite SEDs for radio-quiet AGNs binned by Eddington ratio. The SEDs are normalized at 1 µm.
pted from L.C. Ho, in preparation.)coronaBH
ei (HoBH
1999b, 2002a; Ho et al. 2000) and a substantial fraction of Seyfert nuclei (Ho & Peng1). Defining radio-loudness based on the relative strength of the radio and X-ray emission,≡ 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

X-ray

X-ray
~>

10 cm

46

12:45

~~

15 July 2008

>

ARI

~~~

5

6-13

Hot plasma

RIAF

cold disk

Coronae and RIAFs are collisionless→ protons are naturally accelerated

cold disk

まれる (降着円盤 or 降着流).

Particle Acceleration inする.
Accretion Flows
ルギー天体のエネルギー源
6

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

巨大ブラックホール(太陽質量の１億倍)

Particle-In-Cell Simulations
ジェット(赤)Hoshino 2013; Hoshino 2015
Riquelme et al. 2012; Kuntz et al. 2016

X-1の想像図

円盤
で光る)

銀河(白)

ブラックホール(太陽質量の10倍程度)
電波銀河
視光で光る)の観測
FIG. 1 (color online). Time evolution of the magnetorotational instability.
Panels (a) and (b) show the mlines) and angular velocities in the background at Y ¼ y=λ ¼ 1.91 and X ¼ x=λ ¼ 1.91 (color contour),• Magneto-rotational
instability (MRI) drives turbulencehigh-density regions as reddish curved planes. Panels (b) and (c) are at the same time stage. Panel (f): The eine). •Time evolution of the magnetorotational
instability.
Panels (a) and (b) show the magnetic field lines (great T orbitmagnetic¼ 0.31, 6.89, 7.18,
8.84, and 14.28. The dashed line is a Maxwellian fitting for T orbit ¼ 7.18.
Turbulence
triggers
reconnectionar 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) depi
• Distribution
Thermal→same
Non-thermal
tail
ons as reddish
curved planes.function:
Panels (b) and(c) are at thetime stage. Panel(f): The energy spectra during the061101-2
6.89, 7.18,
8.84,
and
14.28.
The
dashedline
isaMaxwellian
fittingforT¼ 7.18.
orbit in
• Magnetic reconnection accelerates CRs
plasma scale

061101-2

Particle Acceleration in
Accretion Flows

SSK et al 2016; SSK et al. 2019see also Lynn et al. 2014

MHD + Test Particle Simulations

Figure 7. Momentum distributions at t = 10tL in the lab frame (upper)
and the fluid flame (lower) for λini = 4. We can see a bulk motion in thelab-frame, while the bulk motion is not seen in the fluid frame.

ce
&
he

3)

en,

4)

nd
tic

The Astrophysical Journal, 822:88 (11pp), 2016 May 10

Figure 6. Orbits of test particles projected to the R − θ plane (upper panel)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.

where eφ is the unit vector of the φ direction and Vbul, φ is indeFigurependent
4. Colormaps
in theequatorial
plane
for runThe upperdistribution
and lowerof θ. Thebottompanel
shows
the A.
momentum• inshowpanelsdensityandthe magnetic
energyrespectively.
thethefluid
frame,
where
we can see
no density,bulk rotationalmotion. In
the following sections, we use the energy distribution in the fluid• frame. Note that the particle distribution is slightly anisotropic: theparticles tend to have positive p and negative pφ . This is because• the particles tend to move radiallyR outward along thespiral magnetic
field, as discussed above. This anisotropy becomes stronger in later• time and for higher energy particles(see Section
3.2.3). Since
this
MNRAS
485, 163–178(2019)
anisotropy appears in the particle simulations with all the MHD

Figure 4. Long-term evolution of the energy of CRs in the shear frame formodel A1. The thick solid line shows the evolution of energy for the mostenergetic CR at t = 400 tgyro,0 . The thin solid line shows that for the minimum
energy CR. The dotted line shows the average energy of CRs. The mostenergetic CR has about six times higher energy than the average value.

3.2.2. Evolution in Conﬁguration SpaceThe diffusion coefﬁcient in conﬁguration space is estimatedto be

MRI drives turbulenceádx ñ
Turbulence interacts with CRs
,
D =
(20)
2dt
Distribution function: broadened due
to wave-particle interaction
Figure 8. Energy distribution function at t = 4t , 10t , and 25 t in fluidwheredx = x (t + dt ) - x (t ) is the displacement of a particle
flame for λ = 4. The distribution function diffuses in the energy space.
MHD turbulence further accelerates( x = x, y , or z ), dt CRs
is the time span for estimating the
2
i p

xi

i

i

ini

i

i

L

L

L

diffusionandthe average
the CRs.
p denotes function
describes coefﬁcient,
the evolution oftheáñdistribution
for theoverparticles

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8

High-Energy Emission �
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�

Mahadevan et al. 1997; SSK et al. 2015

ν�

p

π

Turbulent#Field##!#proton#accelera4on

•Interaction between CRs and matter/photons→ efficient neutrino & gamma-ray emissions•

TeV-PeV γ-rays are absorbed by γγ pair production→ MeV-GeV photons from the RIAF

Contents

9•

Motivation for hadronic emission from RIAFs- Particle Acceleration in RIAFs- RIAFs as Neutrino Sources•

Hadronic Emission from MADs in Radio Galaxies•

Summary10

Detection of Astro-Neutrinos

High Energy Astrophysical Neutrino Flux Measurement Using Cascades

Hans NiederhauseGodzilla
100 m

IceCube 2013•••

 = .  ( ) . ±!!  " #  

IceCube experiment reported detection  of astro-ν (E ~ PeV) in 2013
Origin has yet to be determined  
Isotropic → extragalactic origin
Soft spectrum in TeV-PeV range   

  [•   ]

IceCube ICRC 2017, 2019

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11

•

Pre-IceCube Models
Cosmic-ray accelerators
mainly pγ interaction• Active Galactic Nuclei (AGN)•

Cosmic-ray reservoirs
mainly pp interaction• Star Forming Galaxies (SFG)

Blazars & luminous Seyfert galaxies

Loeb & Waxman 2006
Manheim & Biermann 1989
Stecker et al. 1991
Halzen & Zas 1997

• Galaxy Group / Galaxy Cluster• Gamma Ray Bursts (GRBs)
Waxman & Bahcall 1997
Dermer & Atoyan 2003
Guetta et al. 2004

Murase et al. 2008
Kotera et al. 2009

12

Model-independent
Constraints
eff [Gpc

Ahler & Halzen 2017

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

Murase & Waxman 2016
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

13•

Models for 10 TeV Neutrinos
Cosmic-ray accelerators
mainly pγ interaction• Active Galactic Nuclei (AGN)• Non-beamed AGN•

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

Kimura et al. 2015
Inoue et al. 2019

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

Murase et al. 2013; Boncioli et al. 2018

Murase+13
Fang & Murase 18

AGN Core Model

14

Mahadevan et al. 1997; SSK et al. 2015, 2019, in prep.
Murase, SSK et al. 2019••

RIAFs in LLAGNs

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below
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estimatedattenuation
becomes
We plotelectron-positron
proton-induced
cascadegamma-rayin to be inef
V. CASCADE GAMMA-RAY EMISSION
decay and process.
high-energy
charged pion decay and the Bethe-Heitler
The electron/positron pairs through π
Fig. 3. A sufficiently developed cascadeZemission
generates
electron-positronpair pro
Hadronuclear
processes
producechargedpion
decay andthe Bethe-Heitler process. The
VHE gammaand
raysphotohadronic
are absorbed bysoft photons
throughthedna flat spectrum below the critical energy at which
γγγ
very-high-energy(VHE) gamma
through
pion
VHE
gamma
rays
areattenuation
absorbed bysoftphotons
through
the
τ
ðε
Þ
≈
R
KðxÞ
dε
;
ð3
γγ → eþ e− process
in therays
RIAF,
and neutral
produce
additionalγγγγ
becomes ineffective. The optical depth
to the
dε
þ
−
þ
−
γ
decay
and high-energy
electron/positronpairs
τγγ ðεγ Þ ≈γγ The→e high-energy
ethroughprocess ein
andpair
produce
additional
high-energy
electron/positronpairs.
ethe RIAF,BH
electron-positron
productionisp
estimatedto beþ
−chargedand the Bethe-Heitler
process. electron/positronThe
high-energy
pairs. The high-energy e e
Turbulent#Field##pairspionalso decay
emit gamma-rays
throughsynchrotron
processes,
Z gamma-ray !#proton#accelera4on
where
ε
is
the
energy,KðxÞ ¼ 0.652σ T
VHE gamma rays are absorbed by soft photons
through
the
γ
pairs
also
emit
gamma-rays
throughsynchrotron
processes,
dninverseComptonscattering,andbremsstrahlung,leadingtoγ
εγ ð36Þ
2 is the gammaÞ ≈ R leadingKðxÞγγ → eþ e− process in the RIAF, and produce
additional scattering, andðxbremsstrahlung,− xτγγ−2ðεÞ γlnðxÞHðx
− 1Þ,dεtoxdε¼ γ ;εwhereγ εγ =ðme c Þ, and HðxÞinverse theComptonelectromagnetic
cascades.
We
calculatecascadeemis−2γ
ðx −Wex tabulateÞ lnðxÞHðx
1Þ,FIG.high-energy
1. Schematic
picture of thepairs.
AGNThedisk-corona
scenario.
electron/positron
high-energyeþ e−cascades. We the
Heavisidestep function[161].
the−valu
electromagnetic
calculatethe cascade
emissionbysolvingthekinetic
equationsofphotons
and-4turbulence.
We
computesteadystate
CRspectraby
solvProtons
areacceleratedbyturbulence
generated
by
the
MRI
pairs also emit gamma-rays through synchrotron
thestepfunctiεγγand, atKðxÞ
which
τRQ
1T in
Table
II. W
sion
byprocesses,
solving thewherekinetic
equationsof10energy,
photons
εγ ofisthe
thecritical
gamma-rayenergy,¼ Heaviside
0.652σ×γγ ¼AGN
electron/positron
pairs
(see
Refs.
[87,159,160]):

(this
work)
in coronae,
andproducehigh-energy
neutrinosandcascadeding
thefollowing
Fokker-Planck
equation(e.g.,[75–78]),inverseComptonscattering,andbremsstrahlung,leading
to+of
thecriticalenergy,εγγ t• Transport equations for−2
2AGN
primary
protonsandsecondaryee&γelectron/positron
pairs
(see
Refs.
[87,159,160]):
can
see
flat
spectrabelow
thecriticalenergy.is (thisNotework)that
RQ(cascade)ðx−xÞlnðxÞHðx
−
1Þ,
x
¼
ε
ε
=ðm
c
Þ,
and
HðxÞ
gamma
rays
via
interactions
withmatterandradiation.
γ
γ
e
!
"
electromagnetic cascades. We calculate the cascade emiscan RQ
see
flat
spectrabelow
AGN(thermalusinge) 3evaluesare CR
approximately
calculated
FIG.
2. tabulated
Disk-corona
SEDs
andlumitheHeavisidestep function
Weprotontabulatediﬀerential
theAGNvalues10-5 [161].
∂Fp of εphotonsFsion by∂Fsolving
and
RL
(in
clusters)pp ∂nεe 1 the∂∂ kineticp
2 equations
43
44 values−1 are tabulated
εpICDþ
= 2þ
+Pff þ P
FeεepÞneεe−%∂+theḞp,inj, Lenergy,½ðP∂nnositiesfor= 10ε42erg
s−1 , 10erg
s−1
, calculations
10
erg
sblazars), are ap
fittingwhile
thecascadeεp Psyn þ
X formula,ofcritical
,
at
which
τ
¼
1
in
Table
II.
We
Cou
RL
AGN
(w.
p
γγ
γγ
e
electron/positron
pairs
(see
Refs.
[87,159,160]):
45
−1
46
−1
∂t ∂t εp ∂ε∂εptp−cool þ t½ðP∂εpe
Psyns þ ,P10
Þn%
esc 10fitting
formula,whilethIC þ erg
ff þ PCouεerg thes thebottomtop).
e(fromtotal to Note
formedcross
section.
Wethatoverplot
spectrawith
below
criticalenergy.
the the Ferm∂εe can see flat∂t
-6exacte
10
formedwith theexactregicronεe(1)LATðγγÞinjeetabulatedvaluesareapproximately
calculatedusingasensitivity
curve
in
the
high
galactic
latitude
metrically ∂nthin,
optically
thick
disk
[61].
The
spectrum
∂
n
_
_
¼
n
−
þ
n
;
ð31Þ
εe
εe
εe ε2 /tðγγÞDe
LAT sensitivity
curve inwhere
Fp ½ðPisεe theþtCR
Psyn distributionþPþ PCouatÞnfunction,acc_ εe ≈−εp ≈þpn_ inj; with a while
ð31Þ
fitting
the exposure
cascadecalculations
are perIC
ffcutoﬀ
εεe % ¼ nεeformula,
10-yearobtained
fromRef. [126].
Tesc
is expected∂tto þhaveanexponential
disk,cut∂εet−1 escwitha
10-yearexposure−1-7
formedthe exactfluxes
cross
Wethanoverplot
the Fermi- curve
diﬀusioninṀ/16πσenergy space,t= tppwith+
3 1/4
10 section.
predicted
areinlower
thesensitivity
fore coeﬃcient
2.8kB Tdiskis
, the
where
Tdisk
≈γ 0.49(GMSB RS ) p−coolTABLE
I.
Parameters
used
this
work.
Units
are
[erg s−1] lowenpredicted
fluxes
are
ε
ðγγÞ
γ
γ
inj
e
LAT
sensitivity
curve
in
the
high
galactic
latitude
region
−1
−1
−1
−1
−1
−1
−3 Array (CTA) has¼ γ n_ εBH−þnεn_ is;the total coolingthe Lcases.
The Cherenkov Telescopenε temperature
γtesc
t∂n
+t
rate,nð31Þ
=γ tfor+tis the maximum
eﬀectiveofγ ðsynÞ
thedisk
εγp−syneðffÞ ∂ninj n
and[cmcases.
] for Thenp , andpγ ε+tX diﬀ
bol , [M⊙ ] for M , [cm] for R,thefallL
Cherenkotesc− γ eþ n_ ðICÞ
ε
ε
ε
with
a
10-yearexposure
obtainedfromRef.
[126].
The
_
_
_
þ
n
þ
n
þ
n
;
ð32Þ
¼
−
ðICÞ
ðffÞ
ðsynÞ
γ
γ
γ
inj
ε
ε
ε
ε
γ
γ
γ
γ
bettersensitivity
above
30
GeV
than
LAT,
but
the casca-8(e.g., [62]).is∂t
Here,
M
is
the
SMBH
mass,
Ṁ
is
the
mass
[%]
for
the
ratio
of
the
CR
pressure
tothethermalpressure.
_
_
_
_
−
þ
n
þ
n
þ
n
þ
n
;
ð32Þ
¼
−
ε
ε
ε
ε
tγγ trate,
the escapeandḞp,inj is
the injection
function
γ
γ
γ 10
γ
esc
bettersensitivity
above 3predicted
fluxes
are
lower
than
the
sensitivity
curve
for
all
2
∂t
t
t
γγ
esc
accretionγ rate,
and
R
=
2GM/c
is
the
Schwarzschildgamma-rayflux
isconsiderably
suppressedinthe VHreacceleration?
Sγγ[79]).
The
stochastic
acceleration
time
isThe
flux
the
cases.
a /Pthis cons∂nεγ (see Appendix
nεγ nεγlog
LX log
LbolCherenkov
log M ΓXTelescopeθe τT Array
loggamma-ray
R(CTA)
log nphas
PCR
ðICÞ
ðffÞ
ðsynÞ
inj
_ object thrange due to the γγ attenuation. For a lower-m
2 _ we use
radius.
Assuminga standard≈i__ disk,_ Ṁ 2−q

8.23
8.68
15
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)SSK, Murase, Meszaros in prep.
where ⇢H↵ is the H↵ luminosity function, see⌧i,IGMis Murase, Toma 2015;
also SSK,
the optical depth in intergalactic medium, and we use
SSK, Murase, Meszaros 2019;1
1
Murase, SSK, Meszaros 2019
H0 ⇠ 70 km s
Mpc , ⌦M ⇠ 0.3, and ⌦⇤ ⇠ 0.7.
H↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡s1(⇢⇤ /L⇤ )/[(L+ (LH↵ /L⇤ )s2 ],•where ⇢⇤ ' 4.11 ⇥H↵ /L⇤ )Preliminary
Seyfert: X-ray & TeV νeson interactions
10 5 Mpc 3 , L⇤ = 3.26 ⇥ 1041 erg s 1 , s1 = 2.78,
We calculate theand s2 = 1.88. We extrapolate thisfunc• luminosity
LLAGN:MeV γ & PeV ν
ng the formalism
38
1tion to Lmin = 10 erg s , below which the Palos by p interac• GeV
mar survey finds a hint of a flattening[69]. γs
Thearesur-attenuatedon given in Ref.
vey also indicates a correlation between in
LX theand source
LH↵ fortic field strength
LLAGNs: LX ⇡ 5 7LH↵ [69]. We use a correctione e↵ect of meson
factor X/H↵ = LX /LH↵ = 6.0. Then,→theconsistentluminosity with Fermi data
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 evoluonic interactions
tion [70–72],
we assume noredshift
evolutionof the lu• AGN cores
can accountfora broadrange
of γ & ν bkgd.
n/positronpairs,
minosity function. The mass of SMBHs in local Seyferts. We calculate
galaxies does not show any correlation with X-ray luetic equations ofpTurbulent#Field##!#proton#accelera4onContents

16•

Motivation for hadronic emission from RIAFs- Particle Acceleration in RIAFs- RIAFs as Neutrino Sources•

Hadronic Emission from MADs in Radio Galaxies•

Summary17

Accretion Flowin FR-I Galaxies

b Jet-mode AGN

3094
McKinney,
A.A.Tchekhovskoy
and R.R.D.
3094 J. C.
J. C.
McKinney,
Tchekhovskoy
D.Blandford
Blandford
Talks by and
Mizuno-san,Nakamura-san, Kino-san

Dominant
radio jetBlack hole

Truncated

Dust?

thin disk

McKinney et al. 2012•
Advection-dominatedinner accretion flow

Powerful jet → Strong magnetic field→ Magnetically Arrested Disks (MAD)Weak narrow-line
regionFigure 4. Evolved snapshot (see Supporting Information for the movie) of the fiducial model at t ≈ 15612rg /c showing log of rest-mass density in colou
the 4.
legendon theright-handin both thez–x plane forat ythe= 0movie)(top left-handpanel) modeland theaty–x
at zg /c= showing
0 (top right-handpanel). The
blackinlines
FigureEvolvedsnapshot(seeside)
Supporting
Informationof the fiducialt ≈plane
15612rlog of rest-mass
densitycolo
field lines,
where
the thicker
blacklinestheshow
is lightly
mass-loaded.
Theandbottompanel
has atthree
The top panel).
subpanelTheshowsṀ linththe legendon theright-handside)in bothz–xwhereplane field
at y =
0 (top left-handpanel)
the y–xplane
z =subpanels.
0 (top right-handblack
BHwhere(Ṁ H ), theout thicker
in the jet(Ṁ j ,lines
at r =
50rgwhere), and fieldout inisthemagnetized
wind (Ṁ, at r =50rghas) withlegend.
The middle
subpanel
shows
ϒ for
mw,o
fieldthelines,
blackshow
lightly
mass-loaded.
Thebottompanel
threesubpanels.
The topsubpanel
shows
Ṁsit
conditions.
Theinbottomthe),efficiency(η) for similar conditions. The horizontal lines of the same colours show the averages over the averthe BH(Ṁ H ), out
the jetsubpanel(Ṁ j , at showsr = 50r
g and out in the magnetized wind (Ṁ mw,o , at r = 50rg ) with legend. The middle subpanel shows ϒ forperiod, The
whilebottom
the square/triangle/circle
are placed
the giventime andThevalues.
In summary,thetheefficiency
is highshow
at η the∼ 200
per cent.
dconditions.
subpanel shows thetickers
efficiency(η) foratsimilar
conditions.
horizontal
lines ofsame colours
averages
overAlso,
the avplenty(uptheto 10times around t ∼ 8500r
of same-signed
magnetic
flux surrounding
the BH,the magnetic
flux reaches
stablevag /c) areperiod,while
square/triangle/circle
tickers
placed at thepolarity
given timeand values.
In summary,the efficiency
is highat η ∼ a200
per saturated
cent. Also,

Heckman & Best 2014

18

Accretion Flowin FR-I Galaxies
3094
McKinney,
A.A.Tchekhovskoy
3094 J. C.
J. C.
McKinney,
Tchekhovskoyand
andR.
R.D.
D.Blandford
Blandford

Figure 20. The final energy partition between background electrons andions, versus σ i .

•Downloaded from https://academic.oup.com/mnras/article-abstract/473/4/48

Ball et al. 2018

McKinney et al. 2012

Powerful jet → Strong magnetic field→ Magnetically Arrested Disks (MAD)Reconnection with strong magnetic field→ most of released energy may
go to non-thermal protons
Downloaded from himportant step in estimating the energy budgets of astrophysicalphenomena, based on intensity of observed radiation from, e.g. systems like BH-powered (blazar) jets and accreting BH coronae (e.g.
•Rees et al. 1982; Rees 1984).
Our simulations are sufficiently large that over 90 per cent of
the dissipated magnetic energy goes to background particles, withless than 10 per cent absorbed by initially drifting particles in the
Harris sheets. Ions (collectively) gain more energy than electrons,
Figure 4. Evolved snapshot (see Supporting Information for the movie) of the fiducial model at t ≈ 15612rg /c showing log of rest-mass density in colou
the 4.
legendon theright-handin both thez–x plane forat ythe= 0movie)(top left-handpanel) modeland theaty–x
at zg /c= showing
0 (top right-handpanel). The
blackinlines
FigureEvolvedsnapshot(seeside)
Supporting
Informationof the fiducialt ≈plane
15612rlog of rest-mass
densitycolo
although the difference disappears as σ i increases, so that the scale
field lines,
where
the thicker
blacklinestheshow
is lightly
mass-loaded.
Theandbottompanel
has atthree
The top panel).
subpanelTheshowsṀ linththe legendon theright-handside)in bothz–xwhereplane field
at y =
0 (top left-handpanel)
the y–xplane
z =subpanels.
0 (top right-handblackseparation between electrons and ions becomes negligible (Fig.field20).
theBHwhere(Ṁ H ), theout thicker
in the jet(Ṁ j ,lines
at r =
50rgwhere), and fieldout inisthemagnetized
wind (Ṁ, at r =50rghas) withlegend.
The middle
subpanel
shows
ϒ for
mw,o
lines,
blackshow
lightly
mass-loaded.
Thebottompanel
threesubpanels.
The topsubpanel
shows
Ṁsit
conditions.
Theinbottomthe),efficiency(η) for similar conditions. The horizontal lines of the same colours show the averages over the aver
the the
BH(Ṁ H ), out
the jetsubpanel(Ṁ j , at showsr = 50r
g and out in the magnetized wind (Ṁ mw,o , at r = 50rg ) with legend. The middle subpanel shows ϒ for
As σ i is lowered, transitioning into the semirelativistic regime,
period, The
whilebottom
the square/triangle/circle
are placed
the giventime andThevalues.
In summary,thetheefficiency
is highshow
at η the∼ 200
per cent.
dconditions.
subpanel shows thetickers
efficiency(η) foratsimilar
conditions.
horizontal
lines ofsame colours
averages
overAlso,
the av
ion-to-electron energy gain ratio increases (hence the electron period,
heatplenty(uptheto 10times around t ∼ 8500r
of same-signed
magnetic
flux surrounding
the BH,the magnetic
flux reaches
stablevag /c) arewhile
square/triangle/circle
tickers
placed at thepolarity
given timeand values.
In summary,the efficiency
is highat η ∼ a200
per saturated
cent. Also,scenarios as the origin of these neutrinos, including
starburst galaxies [51,56–64], galaxy clusters [51,65–69],
19
and radio galaxies [70,71].
We consider high-energy neutrino emission from thevicinity of supermassive black holes (SMBHs) in active
galactic nuclei (AGNs) [72–78]. A luminous AGN hosts
a geometrically thin, optically thick accretion disk that
produces copious UV photons [79–81], and the ratio of the•observed UV to x-ray luminosity is very high [82–84].
Such target photon fields lead to a hard neutrino spectrum•
at PeV energies [85,86]. The accretion shock has been
considered, but the existence of such a shock has not been•
supported by numerical simulations so far. On the other
hand, recent studies on magnetorotational instabilities
suggest that particle acceleration via magnetic reconnections and turbulence is promising in AGN coronae, and•Ref. [87] showed that the mysterious 10–100 TeV component in the diffuse neutrino flux can be explained by theb Jet-mode AGN
AGN coremodel of radio-quiet AGNs. It was found that the
Dominant
radio jetBethe-Heitler process is critically important, which led torobust predictionsof MeV gamma rays � viaγ�proton-induced
p+p ! p+p+Aπ0+Bπ±�cascades.
ν�0�pp+γ !AGNs
p+p+π (LLAGNs), however, have
Low-luminosity
+�
p+γ !
p+n+π
hole
different Black
spectralenergydistributions, in which a UV0
± ! 3ν+e±�
π
!
2γ,
π
bump is absent [88]. This indicates that there is an optiDust? Hot#PlasmaTruncatedcally thin, hot accretion
flow instead of an optically thickthin disk
disk. Remarkably, plasma properties of π
hot AGN coronaeand radiatively inefficient accretion flows (RIAFs [89,90])
in LLAGNs seemBH
similar in the sense that the plasmas
Advection-dominatedp
inner accretionare expected
to flow
be collisionless forions.Turbulent#Field##
It is natural!#proton#accelera4on
to consider the same type of proton acceleration in both
Weak narrow-line
region Seyfert galaxies and LLAGNs. Reference [91] considered the stochastic acceleration expected in such RIAFs
of LLAGNs, and showed that the neutrinos producedby the accelerated protons can account for the diffuse

on the detection prospects of individual nearby LLAGNs.

Model

II. PHYSICAL QUANTITIES IN RIAFs

We consider a RIAF of size R and mass accretion ratec.f.)mass
Mahadevanal. 1997;_ around a SMBH ofM
M BH . Weet use
the notationx
SSK
et al.
2015, 2019,prep
Qx ¼ 10 in cgs units,
unlessotherwise
noted.inTo
represent
the physical quantities in the RIAF, it is convenient
to normalize R by the Schwarzschild radius: R ¼
RRS ≃ 2.95 × 1014 R1 M 8, where RS ¼ 2GMBH =c2 is the
Schwarzschild radius, G is the gravitational constant, and cis the speed of light. The mass accretion rate is normalized_ 2=LEdd , where_ ¼ Mcby the Eddington accretion rate: m
LEdd ≃ 1.3× 1046M 8 erg s−1 is the Eddington luminosity.
According to recent magnetohydrodynamic (MHD)
simulations (see e.g., Refs. [94–99]), the radial velocity,
the sound velocity, the scale height, the number density, themagnetic field, and the Alfven velocity in the RIAF are
estimated to be

Consider Steady & one-zone accreting plasma
Non-thermal protons emit gamma-rays through hadronic interactions
Thermal electrons emit soft photons through Synchrotron & SSC→ take γγ attenuation by soft photons into account
Secondary e+e- pairs emit soft radiations through Synchrotron

ive-mode and jet-mode AGNs (not to scale).
hin, optically thick accretion disk, reaching into the1
V R ≈ αV K ≃ 3.4 × 108R−1=2
α−1 cm s−1 ;12
1
cm s−1 ;Cs ≈ V K ≃ 3.4 × 109 R−1=2
1
21
H ≈ R ≃ 1.5 × 1014 R1 M8 cm;
2_M=2 −1 −1
_ −2 cm−3;
np ≈≃ 4.6× 108R−3
α−1 M 8 m
1
4πmp RHV R
sﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ
8πPg−1=2 1=2 −1=2
_ −2 β0.5 G;B≈
α−1=2≃ 2.6× 102R−5=4
1−1 M 8 mβ
B−1
V A ≈ pﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ ≃ 2.7× 109 R−1=2
β−1=21
0.5 cm s ;4πmp np

20

LERATION …

Soft photons in RIAFs
PHYS. REV. D 100, 083014 (2019)c.f.) Kawashima-san’s talk
Mahadevan et al. 1997;
SSK et al. 2015, 2019, in prep

•
•
•
•
•
FIG. 1. Soft photon spectra for NGC 3516 (red-solid line),
NGC 4203 (blue-dashed line), NGC 3998 (green-dotted line),and NGC 5866 (purple-dot-dashed).

Observations give us the x-ray luminosity, LX , which is_ in our model. Using the bolometricconnected to mcorrection factor, κbol=X , the x-ray luminosity is related
to the bolometric luminosity as
Lbol ≈ κbol=X LX :

ð2Þ

According to the x-ray surveys, κbol=X is higher for a higher•

Synchrotron peak at radio
SSC for IR to MeV gamma
high accretion rate
→ luminous & hardhigher black hole mass
→ luminous
Electron temperature:
always MeV range
Our model is consistent with
X-ray observation in terms of- softening in hard X-ray range- anti-correlation of λEdd - Γxεpstretched
N εpC. Power-law injection
models (B In
andC)turbulence
Refs.
[99,133]).
the
with
a pow
dependso¼ N_ εp ;inj −dN=dε; p , is obtained
ð11Þ by solving the−qtransport equation:
tons’ meofconsider
Pk ∝ka generic, the acceleration
parallel meanto freepastescmodels B and C, we
For
MHD
is
much
! steady-state" proton diffusion
coefficientare estimechanism, andd thespectrum,
N
N εp
εp ¼
εp perpendicular
stretched
i
Refs.
[99
_
−by solving¼
N εp Refs.
;
ð11Þ tons’ mean
εp[145,146,150,151]);inj −
dN=dεobtainedtheN transport
p , is ais the injection function. We
considerpowerof Pk ∝dεptcool(e.g.,
tescequation:
p ;injis perpendimuch! "q−1tion with an exponential cutoff:!
"rL H Refs.
N εp
εp
d
• Transport equation:_(e.g.,
Re_
N εp ;inj is the; [99,1− injection¼ N εp ;inj We
N εp function.
− consider; λk ≈a powerð11Þ!
"−s! where "3ζ rL of Pk ∝k
dεp
tcool
tesc
inj
εp
εpinjection
law
with
an
exponential
cutoff:N_ εp ;inj ¼ N_ 0
;
ð12Þ
exp −
perpendicu
εp;cutεp;cut! function.
"−s We!consider"a power-Dk (e.g., Refs
where N_ εp ;inj is the injection
inj
:
εp
εD
p⊥ ≈
2
_
_
N εwith; 1 þð12Þ
¼ N0
exp −
• Power-law injection:
ðλk =rL Þ
p ;inj an exponentiallaw
injection
cutoff:
ε
ε
p;cut
p;cut
0 is the normalization factor, sinj is the injection"−s" the RIAF is estimated!radius inndex, and εp;cut is the cutoff energy. We normalize ! εThe
Larmorinj
εpp is the injection_ 0 is the_normalization
factor,
s
where
N
;
ð12Þ
¼
N
exp
−
inj
εp ;inj
0
ion by • Normalization:
ε
ε
p;cut
−5=4
−1=2 −1=2The
1=2Larm−1
spectral index, and εp;cut
is≃1.3the cutoff
energy.
normalize
10 p;cut We
_
×
10
ε
R
α
M
m
β
r
L
p;15 1−1
8−2 0.5Z
the injection by_normalization factor, sinj is the injection_ whereεp N_ εp ;inj dεp ¼ ϵp mL
Edd : N 0 is theð13Þ
rL ≃1.3
with our fiducial parameter set (see
Ta

εp
d
−
N εp
dε
t
21 p
cool

Non-thermal Protons

Z
spectral index, and εp;cut is the cutoff energy. We normalize
The Larmo
¼
ε
=PeV.
Then,
we
obtainλ_εp N_ εεp ;inj
dε
¼
ϵ
mL
:
ð13Þ
p;15 p p p
k =rL ≃2.3
Eddwith ou
the injection by−9
get an analytic solution of the transport equation
ing to D⊥ =Dk ≃1.9 × 10 . Hence,rwe
ignoreLε≃1.3p;15 ¼×εZ[149]):
escapethis paper,i.e., weingusetes
We can get an analytic
solutionprocessof the intransport
equation
to
D
_
⊥
_ Edd
εp N εp ;injvaluedεp ¼ofϵpDmL:
ð13Þdue withcould
belargertopossib
Z
our
(cf.
Ref.
[149]):•⊥
loss is dominantescape p∞tcool Infall
000
diffusion. To understand the behavior
εp;15
¼ εof
p=
dεp N_ p;inj ðεp Þ exp
ð−Gðεp ; εp ÞÞ; ð14Þ
value
of
Z ∞
•
pp
interaction
is
the
εp εp
tcoolanalytic 0solution
We can get an
theconfigurationtransport0 equation
protonsinspace,
muching
to Dmo0 ofdiffusion
_
⊥=
N
¼
dε
ðε
Þ
exp
ð−Gðε
;
ε
ÞÞ;
ð14Þ
N
p
p
p
ε
p;inj
p
dominant cooling process(cf. Ref.p [149]):
εp εpcalculations would be required, whichareescapepro
protonsZ ε02 tcool dεpscope of this paper (see Ref. [99] fora related
value
of D
Z ð15Þ
calculatiZ:Gðε1 ; ε2 Þ ¼
∞
t
0
0
cool
ε2 the
0 N
0 for
0 ÞÞ;
dεp ; εcooling
protonprocesses,
we taketcool
diffusion.
_ p;injAs
scope
of
dε
ðε
Þ
exp
ð−Gðε
ð14Þ
ε1 tesc εp N εp ¼
p
p
p
p
;
ε
Þ
¼
:
ð15Þ
Gðε
1 2pp inelastic collisions,
εp εp0photomeson
produ
protonsin
As forε1 tesc εp

SED for M87

22

Preliminary SSK & Toma in prep.
Thermal electrons π decay

Secondary•••

Roughly consistent with observed broadband spectrum
・ → MAD model achieves observed GeV flux
LCR = 0.1Mc
radio (<10 GHz) & TeV γ should be produced by other regions

SED for NGC 315

23

Preliminary
Thermal electrons

SSK & Toma in prep.

π decay

Secondary•

Rough agreement with observed dat
with the same parameter set of M87 except for MBH & ・Mbr making the most nearby
radio galaxies in the future,
41 and luminous examples of such sources

19.5
−1
FRI
must begin to turn overat aroundP1.4 GHz
mheatluminosity
5 MeV function
for all gamma-ray-loudradiogalaxiesas∼a10 W Hz , or
FRII
(including
Centaurus
Centaurus
B, andM than87)L40particularly
promising targets of study.
acedensityof radio sources
would exceed A,
the spacedensity of galaxies
brighter∗,model.
24 Then, we are able to define the average SED
typical hosts. Cattaneo & Best (2009) similarly argue for a turnover no fainter than 38
39
40
41
42
43
44
4
f the
gamma-ray-loudradio galaxies for all luminosities
19.2
−1
6
10
W Hz
by
comparison
with
the
local
space
density
of
SMBHs
above
10
M
.
⊙
ϵ ∝ ϵ −2.39 at ϵ >5 MeV
andanddN/dϵ ∝ ϵ −1.89 at From Radio
Rays
Galaxies log10(L5 GHz [erg/s])
observed2radioGamma
luminosity functions
now broadlyNeutrinos
constrain the entire distribution ofeV
by following Equation (1).
Figure 1. Gamma-ray luminosity at 0.1–10 GeV vs. radio luminosity at 5 GH
luminosities.
ver,
only
three
sourcesarecurrently
studiedwithmultiThe squareand triangle data represent FRI and FRII galaxies, respectively. T
GN populationasrevealedby radiosurveys is composedof unresolved
two distinct categories
of galaxies
The
sum
of
the
emission
from
all
radio
leads to a gamma-ray flux that
solid
line
is
gth
observationaldata.
We need to makefurtherstudies referred to as the fit to all sources.
population
with strong
QSO/Seyfert-likeemissionlines (historicallyis given
by: radioversion of this figure is available in the online journal.)
dual
gamma-ray-loudto understandtheir(referred(Ation sources)
and a population
withgalaxies
weak LINER-likeemission lines
to color
as lowoperties
inHinewide
luminosity
ranges.
examineother of the present reZ LaingZ
ources; e.g.,& Longair
1979,
al.We
1994).
In the language
2et
dF
dF
dL
d
V
modelsin Section
5.2.
areradiative-mode
and jet-mode= AGNs,dz where the high excitation sources⇢form(La small, z)(1 !(F (L , z))) exp( ⌧ (E , z)),ion of radiative-mode
AGNs that produce
powerfuldE
jets. Best
Heckman (2012) sepadE d⌦dz d⌦
L &log(10)Figure 1 shows the 5 GHz and 0.1–10 GeV luminosity relatio
dio luminosity function into these two radio source populations and showed that both3.verGAMMA-RAY
LUMINOSITYFUNCTION
of sources
Fermi gamma-ray-loud radio galaxies. Square and triang
2radiothe fullrange ofpowerstheyobserved(see Figure 11).
The jet-mode
where
d
V
/dz
d⌦
is
the
co-moving
volumeelement,
and/dE
is thespectrumofrespectively.
gamma- Th• dF• 95% of the radio-AGN population at all radio luminosities below Pdata
represent
FRI
and
FRII
radio
galaxies,=
about
1.4 GHz
1.−1Radio
andfrom
Gamma-ray
Luminosity
rays
a radiogalaxyofCorrelation
luminosity
L26 W
and
locatedredshiftz. The! represents−1shows
the fitting
line function
to all the data.
The function, but theradiative-mode
population
becomesdominant
above 10
Hzsolid. In line
the atHigh-Energy Backgrounds
from MAD in Radio Galaxies

Convert Radio luminosity
Radio luminosity function
to γ-ray luminosity
for “jet-mode AGN”
by accounts for the fact that resolved radio
Fermi’s point source detection efficiency [54],given
which

log 10(Lγ [erg/s])

Number density/Mpc–3 log10 (L)–1

imate the contribution of gamma-ray-loud radio galax10–4
EGRB,
we need do
to construct
a GLF. However,galaxies
not contributeto thebecausedi↵use gamma-ray background2 . The attenuation of thelog10 (Lγ ) = (−3.90±0.61)+ (1.16±0.02)
log10 (L
0.1
Astrophysical733:66to(9pp),
2011 May a20GLFz =Inouemall The
samplesize, itJournal,
isspectrum
difficultconstructusing
5 GHz ), (5gamma-rayfrom
scatteringwiththeextragalactic
backgroundlightischaracterized
10:48 10–5 data alone. Here, theM87-likegamma-rayof radio galaxspectrum
is given
by depth, ⌧RLF
48 the model of Ref. [55].
bytheoptical(E, z), (see,for whichadoptbeen extensively studied in previous works
e.g., wewhere
errors show 1σ uncertainties. In the case of blazars, th!studyInoue Y. 201147
Inarecent[23],
werefinedtheempirical
betweentheLradio
and gamma−(p+1)/2& Peacockis a corslope ofcorrelation
the correlation
betweenϵ !Ifϵthere
ϵ al. 2001).
γ (>100 MeV), luminosit
10–6 1990; Willott et
br ,
dN/dϵ
∝
betweentheemission
radio
and detectedgamma-rayluminosities,
we are(1)(seeabove
MeV,
radio
luminosity
at 20thisGHz
is 1.07 ± 0.0
ϵ >radioϵbr ,ϵ −(p+2)/2 fromraygalaxies
also 46100
Refs.
[42,and43])
andcombined
informaconverttion
the
RLF to the GLF with that correlation. In(Ghirlanda
et al. 2011).
The correlationof gamma-ray
45
redshiftdistribution[56]slopes
to determine
10–7 with the measured radio luminosity function andwhere ϵbr itcorresponds
the IC photon
fromelectronsof blazars,
has been tosuggestedthat energy
there isa correloud radio galaxies are similar to those of blazars. This indicate
44 ⇢ (L , z), and ultimately the total conthe
gamma-rayluminosity
functionofradiogalaxies,withγ(Rybicki&
Lightman1979).
brHeckman& Best 2014
tween the radio
and gamma-ray
luminosities from thethat the emission mechanism is similar in gamma-ray-loud radiThefitting
for NGC
1275
and M87 shows
thatthe IC–8
10SED
tributionfromallunresolved
radio
galaxiestothe43di↵use
gamma-ray background. In doing
era peak(Padovani
et
al.
1993;
Stecker
et
al.
1993;
Salamongalaxiesand blazars.
Radio-loud,radiative-mode
AGN
energy in thegalaxy
rest framesurvey
is locatedat ∼5 MeVMauch(Abdo et al.
eer6-degree-field(6dFGS;&
Sadler
2007)
are
1994;
Ghisellini
et al. index,
2001;Γc ,Radio-loud,jet-mode
AGN
We 42
need to examine whether the correlation between th
2009b,Dondi
Incasethis study,
we1995;use
theZhangmeanthis
photon22009c).&
In
the
of
radio
galaxies,
calculation
depends
very
little on the precise form of !, as only a very
to &asTotani2006),
althoughhas also
been
10–90.1–10
radio
and
gamma-rayluminosities is true or not. In the flux
how
that
the
local
radio
function
of
AGNs
is
well
fittedΓ atGeVand set luminosity
aitpeakenergy,ϵbr , discussed
in the photon
41
22
23
24
25
26 flux27originates from abovesmall
fraction
of
the
total
gamma-rayor
near
thresholdsources. FRI
spectrumcannot
at 5 MeV
forall gamma-ray-loudradio galaxies
as acorrelation
befirmlyestablishedbecauseoffluxlimited
observations,theluminosities
of FRII
samples are strongl
–1
log
(L
/W
Hz
)
imilarin form
toEquation
6:to define the average SED
10 NVSS40
baselinemodel.
Then,
we
are
able
amples (Muecke et al. 1997). Recently, using the Fermi
correlated
This
might
resultinlu
38 with39redshifts.
40
41
42
43
44 a spurious
45
!
"
shape
of
gamma-ray-loudradiogalaxiesforallluminosities
24.95
P0=10
W/Hz,
β
Ghirlanda
et(P/Pal. (2010,2011) confirmedthat
there−1.89
is aminosity correlation.
Asloginprevious works on blazar sample−2.39)α +. radiative-mode(7)
ρ0 /function
as=dN/dϵ
∝ ϵof
>5
MeV
and
dN/dϵ ∝-5.33
ϵand jet-mode
at
0 at ϵ (P/P
0 ) into
10(L5 GHz [erg/s])
dio ρ
luminosity
radio-loud
AGNs
split
sources.
-3
ρ0=10
Mpc
on
betweenthetheradio
and Equation
gamma-rayluminosities.
(Padovani
1992; Zhang et al. 2001; Ghirlanda et al. 2011
ϵ !drawn
5 MeV
following(1). (2012),
largely
frombyresults of Best& Heckman
who split these two populations
Figureand
1. Gamma-ray luminosity at 0.1–10 GeV vs. radio luminosity at 5 GHz.
amine
a luminosity
correlation
iningamma-ray-loudradio
weperforma partial
correlation
analysisto test thenosityfunctionsusing
radio-selected
AGNs
the Sloan
DigitalSurveymain
galaxysample.
However,only
threesourcesarecurrently
studiedSky
withmultiThesquare andgive
triangle
data represent
FRI and FRII
galaxies, respectively.
The correla
valuesfromcombiningallfourspectroscopic
surveys
above
–
2
–
shown
by Gendreetal.the(2013),
theseresults
underestimate
theluminosity
function of
the linesolid
is the fit tothe
all sources.
wewavelength
firstderiveradio
andgamma-rayluminosity
observationaldata.
We
need to−1
makefurtherstudiestion
betweenradio and gamma-ray luminosities exclud−1
24.95±0.14

High-Energy Backgrounds
from MAD in Radio Galaxies

25

Preliminary
Thermal γ

Neutrinos
Hadronic γ•

MeV gamma-rays, GeV gamma-rays, & TeV-PeV
neutrinos can be explained simultaneouslywith the same parameter set

Contents

26•

Motivation for hadronic emission from RIAFs- Particle Acceleration in RIAFs- RIAFs as Neutrino Sources•

Hadronic Emission from MADs in Radio Galaxies•

Summary27

Summary•••

Cosmic-rays are naturally accelerated in hot accretion flows,
especially in MADs by magnetic reconnection
Radio-galaxies likely host MADs, leading to hadronic emissions
With a one parameter set, hadronic emission from MAD canreproduce the gamma-ray observations for M87 & NGC 315
With the same parameter sets, hadronic emission from MADs willproduce wide range of cosmic high-energy backgrounds.

AA52CH13-Heckman•

ARI

4 August 2014

10:48

a Radiative-mode AGN
Radio
jet

Dusty obscuring
structureb Jet-mode AGN
Dominant
radio jetType 1Preliminary

Preliminary
Black hole
Broad-line
regioned from www.annualreviews.org/16/20. For personal use only.

Accretion
diskRadio-loudTruncated

Radio-quiet

thin disk

Dust?

Black hole
Type 2

Narrow-line
regionAdvection-dominatedinner accretion flow
Weak narrow-line
region28

Thank youfor your attention