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

HEAP 2019@ 蔵王温泉

December 2019

2

Index
•

IceCube Neutrinos•

Accretion Flow in AGN•

AGN Corona model Murase, SSK, Meszaros, arXiv:1904.04226••

LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

3

Index
•

IceCube Neutrinos•

Accretion Flow in AGN
Murase, SSK, Meszaros, arXiv:1904.04226•••

AGN Corona model
LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

4Detection of AstrophysicalNeutrinos
PTEP 2017, 12A105(a)

PTEP 2017, 12A105
(a)

•
•
•

IceCube experiment reported
detection of astro-ν (E ~ PeV) in 2013
Shower: good for spectrum
Track: good for source searchShower or Cascade (νe)(b)M. Ahlers and F. Halzen

Fig. 3. (a) Light pool produced in IceCube by a shower initiated byenergy is 1.14 PeV, which represents a lower
μ limit on the energy of t(b)dots represent sensors with no signal. For the colored dots, the colo
purple (late) following the rainbow, and size reflects the number otrack traverses the detector at an angle of 11◦ below the horizon. T
2.6 PeV.

Track (ν )signals are sorted into telltale patterns of light that reveal t
incident neutrino.
Even at a depth of 1450 m, IceCube detects a backgrou
originating in the Southern Hemisphere at a rate of 3000 per seneutrinos. Traditionally, neutrino searches have focused on tinteract primarily outside the detector to produce kilometerinstrumented volume. Although this allows the identification
detector, it is necessary to use the Earth as a filter in order to r
ray muons. This limits the neutrino view to a single flavor aICRC 2019 Results from IceCube

5

independent sample of starting events designed to study the inshows the best fit flux normalization and spectral index for t
starting event sample, and also the sample of contained cascad
no evidence of the prompt atmospheric neutrino component inlimit on the prompt neutrino flux is ongoing. A global fit of th

Neutrino Spectrum
electron and tau neutrinos
ICRC 2019 Results from IceCube

Dawn R. Williams

independent sample of starting events designed to study the inelasticity distribution [14]. Figure 1
shows the best fit flux normalization and spectral index for the through-going track sample, thestarting event sample, and also the sample of contained cascades presented in 2017 [15]. There isno evidence of the prompt atmospheric neutrino component in the sample; work to place an upper
limit on the prompt neutrino flux is ongoing. A global fit of the diffuse samples is forthcoming.

Figure 1: Best fit flux normalization and spectral index for a single(blue), contained cascades (green) and starting tracks + cascades (and 99% uncertainties respectively. Note that the all-flavor flux norm
divided by 3 to compare to single-flavor normalizations.

muon neutrinos•••

2.3 Tau Neutrino Searches

Assuming that the flavor ratio at the source is 1:2:0::ne :nphysics predicts a flavor ratio of 1:1:1 in IceCube. Even in extre
composition at the source, the flavor ratio in IceCube will stilFigure 1: Best fit flux normalization
spectraloscillations.
index for a single
poweranalyses
law fit toof
through-goingtracksi
afterandstandard
Severalthe flavor ratio(blue), contained cascades (green)+ cascades(yellow).
Inner/outer
68%
16,and17]starting
but a tracks
limiting
factor has
been thelack ofcontours
clearly are
identifi
and 99% uncertainties respectively. Note that the all-flavor flux normalization of the starting event sample isappear as single cascades in IceCube, and therefore the experim
divided by 3 to compare to single-flavor normalizations.
of neutral current events and ne charged current events. The si
current event is a cascade from the neutrino deep inelastic s
2.3 Tau Neutrino Searches
signature from the decay of the tau lepton. The tau lepton dec

Track analysis: flat spectrum (E > 200 TeV)
Cascade analysis: soft spectrum (E > 1 TeV)
Hint of 2 component??
Uncertainty of analyses??

6

Arrival Direction

The Astrophysical Journal, 809:98 (15pp), 2015 August 10•••

Isotropic → Extragalactic origin
No point-source detection so far
Neutrino flavor ratio is consistent withpion decay[(1:2:0) at Source]

7

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

νμμ
ν
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!"
Eγ’[GeV]though the equationphZ 1 q! !E! "
in
dN$

8

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

νμμ
ν
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!"
Eγ’[GeV]though the equationphZ 1 q! !E! "
in
dN$

9

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

νμμ
ν
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!"
Eγ’[GeV]though the equationphZ 1 q! !E! "
in
dN$

10

Point Source Constraint
see Murase & Waxman 16

104

102

103

RQ AGN

GC-accRL AGN10BL Lac

1
0.1
10 2
10 3

IceCube5 IceCube
20 IceCube

FSRQ

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

107
106
local rate density ˙ [Gpc 3 yr 1 ]eff [Gpc

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

effective local density

3

]

106

105
104
103
102

discovery potential fortransient source candidates
in 10 years
SNe & newborn pulsars
hypernovae
LL GRB

10
1
0.1
10 2

IceCube5 IceCube
20 IceCube

HL GRBjetted TDE

1046 1047 1048 1049 1050 1051 1052 1053 1054 1055
effective bolometric neutrino energy E [erg]Figure 2: Left: Comparison of the diffuse neutrino emission (solid magenta band) to the effecNodensity
point-source
detection
tive• local
and luminosityof extragalactic neutrino source populations. We indicate several candidate
populations(î)density
by the required
neutrino luminosity
to account for the full diffuse→ Highnumberof neutrino
sourcesflux [17] (see also [25]). The lower (upper) edge of the band assumes rapid (no) redshift evolu• The
IceCubealready
disfavors
luminous
sourcestion.
dark-blue-shaded
regionindicates IceCube’s
discovery
potential of the closest source2122of the (GRBs,
population (Ef⌃
10
TeV/cm/s in the Northern Hemisphere [26]). Right: The
Blazars,
Jetted TDEs)nµ +n̄µsame comparison for transient neutrino sources parametrized by their local density rate and bolo-

11

Gamma-ray Constraints
E 2 J [GeV cm 2 s 1 sr 1 ]

+
−
0
π
=
π
=
π
10 6

HESE (3yr)Fermi
gammas10

arXiv:1410.1749
Fermi IGRB (2014)

cosmicneutrinos

7E-2.1510 8

10 9

pp scenario
SFR evolution
10 3

•
•

10 2

0.1

1
10
E [TeV]

102

103

dark sources below 100 TeV not seen in g’s ?
gamma rays cascade in the source to lower energy

Astrophysical νs are accompanied with γ rays
ν intensity at 10 TeV > γ-ray intensity at 100 GeV→ accompanying γ rays overshoot Fermi data→ ν sources should be opaque to γ rays Murase et al. 2013, 2016
Ahler & Halzen 2017

12

Index
•

IceCube Neutrinos•

Accretion Flow in AGN
Murase, SSK, Meszaros, arXiv:1904.04226•••

AGN Corona model
LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

13

Active Galactic Nuclei(AGNs)•

Radio-quiet AGNs•

Radio-loud AGNs

M77 (NGC 1068): Wikipedia©

••

No prominent jet
90% of AGNs

M87 (NGC 4486): Wikipedia©

••

Powerful Jets
1-10 % of AGNs

Outlook on Results

14

shown are relative to the first
DOM hit according to the trackreconstruction, and earlier andlater times are shown with thesame colors as the first andlast times, respectively. The
total time the event took to
cross the detector is ~3000 ns.
The size of a colored sphere is
proportional to the logarithmof the amount of lightobserved at the DOM, withlarger spheres corresponding
to larger signals. The total
charge recorded is ~5800 photoelectrons. Inset is an overhead perspective view of th

Active Galactic Nuclei(AGNs)●

NGC106
8
GB6J
154
2+6
129
TXS0
506
PKS1424+240

•

NGC1
0
6
8a
l
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AGNs
NGC1
0
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. • Radio-loudconsistent with a zenith angle 5:7
degrees below the horizon.

Radio-quiet AGNs

þ0:50
"0:30

Multi-messenger campaign

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M77
(NGC
Ne
utrino1068)
Telescopes – March 22nd 2019•

Hottest Point in Northern Sky (2.9 σ )IC 170922 (3 σ)9

TXS 0506+056• 2. Fermi-LAT and MAGIC observations of IceCube-170922A’s
Fig.
location. Sky position of IceCube-170922A in J2000 equatorial coordinates
overlaying the g-ray counts from Fermi-LAT above 1 GeV (A) and the signal•
significance as observed by MAGIC (B) in this region. The tan squareindicates the position reported in the initial alert, and the green squareindicates the final best-fitting position from follow-up reconstructions (18).

2014-2015 Neutrino flare (3.5 σ)

pixel, usregion asmoothesignal si
a g-ray s
SourceOutlook on Results

15

Active Galactic Nuclei(AGNs)●

NGC106
8
GB6J
154
2+6
129
TXS0
506
PKS1424+240

•

NGC1
0
6
8a
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Radio-quiet AGNs

Multi-messenger campaign

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Telescopes – March 22nd 2019•

Hottest Point in Northern Sky (2.9 σ )9TXS 0506+056•

IC 170922 (3 σ)•2014-2015 Neutrino flare (3.5 σ)

16

15 July 2008

10 cm

46

1 cm

12:45

Radio-quiet AGNs

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

18

Opt.

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.)~~~~

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

log (ν/Hz )~>log (νLν /ergs s–1 )Ho 2008

ure 7

10 keV~~~~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

Hot plasma

cold disk

X-ray

X-ray
~~
~>

ARI

>
~
~
~

6-13

Hot plasma

RIAF
cold disk

1746

15 July 2008

10 cm

1 cm

12:45

Radio-quiet AGNs

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
–1 ≤ log L /L < 0
Protons in coronae& RIAFs[Radiatively
are collisionless
Inefficient–3 ≤ log L /L < –1
log L /L < –3
→ Non-thermal
protonAccretion
production
Flow (RIAF)]

ure 7

42

bol

40

bol

38

9

10

11

12

QSO UV X-ray

13

•

Edd
bol

Edd

bol

Edd

Edd

** electrons are collisional **
14

15

16

17

18

log (ν/Hz )

Opt.

~>

log (νLν /ergs s–1 )Ho 200844

10 keV>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.)

LLAGN

~~

~~

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

Hot plasma

cold disk

X-ray

X-ray
~~
~>

ARI

>
~
~
~

6-13

Hot plasma

RIAF
cold disk

Azimuthally
averagedVφ , ⟨Voutward-flowing
Thewithwhitealines
iso-contoursof ⟨V
where Ωc% ¼ e% B0 =m% c. The p
gridsize Δ was
set toφ ⟩L , on the R − φ plane.
φ ⟩L .
streams
highareplasma
density
andﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ
contours in the background at Y ¼ 1.91
and X ¼ 1.91
3=2
2
2 ðvt% =Ωp% Þ, where Ωp% ¼ 8πne =m% is the pairstrong electric current called the channel flow. The reddish
show the angular velocity vy in the local rotating
frame. InVasthebackground
velocitysandwiched
for analysesbyoftwo
the surfaces
test-particle
plasmafrequency.
The
Alfvénvelocityisdefinedasregions
in Fig. 1(c) show thebul,
φpﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃﬃ
the 18early stage at T orbit ¼ Ω0 t=2π ¼ 0.31 in Fig.
1(a), ﬃ thesimulations
inSection
3.2.
V A ¼ B= 8πm% n, so that the plasma β is equal tohigh-density channel flow with ρ ≥ hρi þ 2σ ρ where hρimagnetic field lines are parallel to thez axis, and the
4 plots
thecolormaps
density (upper) and the magnetic
3v2t% =V 2A . The parameters used were Fig.
ðV A =ΩÞ=Δ¼
25,
andofσthe0ρ are the average density and standard deviation of
Keplerian motion or differential motion of vy can be seen as−3
energythe equatorial
plane. Theinmagnetic
fields domain,areðvt% =Ωc% Þ=Δ ¼ 56.4, V A =c ¼ 6.25 ×
10 . (lower)
N x , N y ,onand
density distribution
the simulation
respectively.
the color contour at Y ¼ 1.91, whereNtheare
reddish(bluish)frozeninthedifferentiallyrotating
fluid
elements
that
fall
to
the
the grid sizes in the x, y, and z directions,
The amplification of the magnetic field stretched by thez
region corresponds to a positive (negative)
toroidalandvelocBH.
thespiralstructure
as seenmay
in theWe bycanthe magnetic field
Keplerianmotionbefigure.
balancedrespectively,
we assumed N x ¼
N z This¼ Ncreates
¼
300
in
z
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 resulting
The amplification of the magnetic field stretched by the-1の想像図Table
2
coefficients indicate that thedensityand magnetic energy are weakly
Keplerian 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
8-4
光る)
20Trot
3.4
´
10
1.59
´
10
4
2
416
50
2.3
17
1.6
2.38
0.30
A1
higher coefficients,
i.e. the anticorrelations are weaker, but no run
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
et al. 5.83
2019
¥ Finally,¥ the azimuthal
20Trot3.4 ´we
10 discuss
´ 10
X1
4
200SSK3.1
192016;
2.0
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
2012; Suzuki&3.5Inutsuka
Parkin &6.44 ´ 10-5¥
20Trot6.8 ´ 108et al. ¥X3
8200
21 2014;
3.0 see 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 background
1.91
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)affect
where
m
=
k
R
(k
is athewavenumber
theT orbit
φ direction).
at T orbit ¼ 0.31,
7.18,
8.84,
14.28.
The fordashedMaxwellian
fittinginfor¼ 7.18. Then,φ lineφ isscale 6.89,
that CRs
travel,weand
ignoreits effectsimplicity.
We 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 wave-particle interaction 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 fromsee 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, the19High-Energy Emission

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• Interaction between CRs and matter/photons
FIG. 2. Disk-corona SEDs and CR proton diﬀerential luminosities for L = 10 erg s , 10 erg s , 10 erg s ,10 erg s , 10 erg s (from bottom to top).
→ neutrino & gamma-ray emission
metrically thin, optically thick disk [61]. The spectrum
is• expected
to have an
exponential
cutoﬀ
at ε
≈
TeV-PeV
γ-raysarereprocessed
to MeV-GeV γ rays
2.8k T , where T≈ 0.49(GM Ṁ/16πσ R )
•

42

X

45

B diskdisk

disk,cut
3 1/4
SB S

is the maximum eﬀective temperature of the disk
˙

−1

46

−1

43

−1

44

−1

−1

TABLE I. Parameters used in this work. Units are [erg s−1 ]for LX and Lbol , [M⊙ ] for M , [cm] for R, [cm−3 ] for np , and

20

Index
•

IceCube Neutrinos•

Accretion Flow in AGN
Murase, SSK, Meszaros, arXiv:1904.04226•••

AGN Corona model
LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

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’
E’ > E0ていく。
E’< E0E’
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。
＜〜〜〜

21

ことなく粒子はエネルギーを増加させており、電磁場との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)

According to quasi-linear theory,
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 results inの方程式を
Equation Pと呼ぶこともある。
Dp Fokker-Planck∝ pq (power-spectrum
k = P0k-q)
ここで p は運動量の大きさであり、超相対論的を仮定して
2

Fi
布は

を行

と表
分
でき
Dp

green squares are located.
"3p@ 22 2@Fp!

Fp"p D " p+Fp+ Ḟp,inj ,
"p
@"p
tcool
tesc
(1)
where
Fp is thedistribution
momentum function
distributionfunctionthemomentumfor2
turbulence.
We compute steady state CR spectra by solv10-4(dN/d"p = 4⇡p
2 Fp /c), D"p is the di↵usion coefficient,/d"4⇡pcoolingFp the/c),time,
D"pt is isthedi↵usionp =ingfollowing
Fokker-Planck
equation(e.g., [75–78]),cf.) Dermer et al. 2012, SSK et al. 2015tcool
is thetheescape time,
and
esc
•
Stochastic
Acceleration(SA)
RQ AGN
is is
the
ool
esc
! protons,"
F˙p,inj
thecoolinginjection timeterm toforthestochastic tacceleration.
-5
3
10
ε
e Considering
time, andresonant
Ḟ
is
the
injection
funcscatterings
with
Alfven1 ∂∂Fp waves,pthe
Fp∂Fp,inj
p
2
εpas
Das
=is represented+
Fp −+ Ḟp,inj ,ε[144–146]
p D
di↵usioncoefficient2↵usioncoefficient
canbewritten⇡"∂tεp ∂εp p∂εpp tp−cool
tesc
total
q 2 2
2
-6
/H) "p , where A = B/ 4⇡mp c Np is10(1)✓ ◆2 ⇣ ⌘q 2 rL =locity, H ⇡ where
R/2 is⇣c
theis
scale
2
VA
rheight,
L distribution
2
F
the
CR
function,
D
≈
ε
p
ε
p
p /tacc
D "p ⇡"p ,(9)e Larmor radius, ⇣ isHthe cturbulent
H strength,−1
is the diﬀusion coeﬃcient in energy space, t−1
=
t
+
10-7
pp
p−coolpower-law index ofthe −1
turbulence power
−1
−1
−1
−1
−1
t
+t
+t
is
the
total
cooling
rate,
t
=
t
+t
esc
BH Ḟ0 p−syn
fall
diﬀ]. We set Ḟpγ("p"p,inj ) withp,inj =
2is the
rate,
and
10-8
theescapevalue of"p,injhasḞp,inj
no is the injection function
p c . Note that
reacceleration?
Appendix[79]).
The
stochastic acceleration time is
he resulting (seespectrumas longas theinjec22−qs muchlowergiven
than: bythecuto↵
energy.
Ḟ
is• Escape
t
≈
η(c/V
)
(R/c)(ε/eBR), where VA
Diffusive
escapeA 0& infall
p to SMBH
R acc
10-9
o that the condition
L"p d"velocity
istheAlfvénand
10-4 10-3 10-2 10-1 1
p = ⌘p ṁL
Eddηisis the inverse of the turbu• Coolings:
pp
inelasticcollision,
photomeson
production
1lence strength
[80,
81].
We
adopt
q
=
5/3,
which
is conre L"p = tloss "p dN/d"p is the proton lumiFIG. 1. Schematic
picture of the AGN disk-corona scenario.
++e
1 recent1 MHD
Protons
accelerated by(p+γ→p+e
turbulence generated
by -thesynchrotron,
Bethe-Heitler
process) MRI
sistentproton
with[56], aretogether
he injection parameterandthetloss= tcool +tesc1 simulations
in coronae, and produce high-energy neutrinos and cascaded
FIG.
3. EGBand
ENB spe
with
η
=
10.
Becausethe
dissipation
rate
in
the
coronaess rate.
We use&the
Chang-Cooper
method
gamma
rays
via
interactions
withmatterandradiation.
• MuonPion Coolings are negligibly inefficient
The data are taken from Sw
is expectedi↵usion equation[56, 57].to be proportional to LX , we assume that theloon-)[87] (blue), SMM [88]HEC,γ-raysareabsorbed
byscalestarget(γ+γthe→ e++einjectionfunctionas photons
LX . To explainB • andwe generalize
thelinearly
acceleraFermi-LAT [90] (orange), a
44
−1
ENB,
the
CR
pressurerequired
for
L
=
10
erg
s
X
by a power-law
injection
term, Ṅ"pemit,inj = high-energy→ electron& positron
gamma-rays
metricallythin, optically andthick upgoingdisk [61]. muon
The spectrum
track (
turns
out
to
be
∼
1%
of
the
thermalpressure,
which
is
sinj
is expected to have an exponential
cutoﬀofatreaccelerate
εdisk,cut ≈
exp(
"
/"
),
where
"
is
the
cutcontribution
p
p,cut
cut
−1
4
3
−1
→ CalculatecascadesR plot εp Lεp ≡ 4π(εreasonable.electromagneticWe/c )Fp 2.8k
V(tBesc+t, p−coolTdisk
where )Tdisk ≈ 0.49(GM Ṁ/16πσSB RS3 )1/4
p
nd Ṅ0 is normalized by "p Ṅ"p ,inj d"p =is the maximum eﬀective temperature of the disk
in Fig. 2, where V is the volume.
ut is obtained by equating the infall time,(e.g., [62]). Here, M is the SMBH mass, Ṁ is the mass
E2  [GeV cm-2 s-1 sr-1]

Basic Equations

23

Target Photon Field
•

•
•
•
•

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

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
Higher Lopt/Lx for higher Lx AGNs
Softer spectra for higher Lx AGNs

24

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 canHigher 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

25

HE particles from
Nearby Seyfert Galaxies
4

10-6
-7

-2 -1

E FE [GeV cm s ]10Generated +e
Cascade 
Thermal e 

LX=1044 erg/s
d=30 Mpc
reacceleration?

oeASTROGAM•

A typical Seyfert at 100 Mpc•

pγ neutrinos are detectable
by IceCube-Gen2
MeV γ-rays can be detectedby future satellites.

IceCube (=30 )

10-8

10

•

Bethe-Heitler

-9
AMEGO

pp+p

10-10
10-11 -4 -3 -2 -1 0
10 10 10 10 10 101 102 103 104 105 106 107
E [GeV]MeV
γ-ray
is determined
production
FIG. 4.• Point
sourcefluxesluminosityof all flavor neutrinos
and gamma by B-H pair
FIG. 1. Schematic picture of the AGN disk-corona scenario.
Protons are accelerated by turbulence generated by the MRI
ays from a→nearbyAGN
withastandard
disk.
A
possibleefin coronae,
and producefieldhigh-energy neutrinos and cascaded
Ratio
of
γ
to
ν
flux
is
fixed
by
the
observedphotongamma rays via interactions with matter and radiation.
ect of secondary reacceleration is indicated (thin solid). For→ We
robustly
test
our model
ASTROGAM[102] canand AMEGO[103]
sensitivities,
the ob- by future experiments
6
ervation time of 10 s is assumed. The IceCube eight-year
metrically thin, optically thick disk [61]. The spectrum

denser x-ray target. At the same time
at a lower energy than for the lower luminosity AGN.
break 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
Z
L
z
1
dn0
dV
max
max
diff
effect of decreasing the magnetic fiel!!;ob"E!;ob# $!!;ob;dLX
dz
f"z#
4" Lmindz
dLX
0source and, hence, the synchrotron bre
energy as LX increases. However, the o(36)
44
AGNs with Lx ~10 erg/s provide
contribution
interactions
stays the same as MBH i
3 the dominant
e.g.,
Ueda
et
al.
2014
where dn0 =dLX 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
AGN

(this
work)
luminosity and evolution functions of Ref. [9] (model
panel of Fig. 5.
AGN (cascade)  (this work)I). In a Friedmann-Robertson-Walker
universe the coAGN (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
remax
FIG. 1. Schematic picture
theAGN disk-corona
scenario.
E, andprotons
can escape from thepof
Protonsareacceleratedbyturbulence
generated
by
the
MRI
•
We
choosetheinjection
efficiency
high-energy neutrinos and cascadedz# &and produceing
spect to redshift z is "dt=dz#!1 $ !H0 "1in% coronae,
to the observed cosmic ray 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
10 erg s , 10 erg s (fromlowluminosity
AGNs
would 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)!8!2 sR!1
TABLE I. Parameters used inand the source frame are related by E! $
E
"1
%
z#,
GeV
cm
[15].
However,10
isthemaximum
eﬀectivetemperature
of
the
disk
• Energetically!;obfor L andL , [M ] for Mth,reasonable:
X

26

Extragalactic γ & ν
Backgrounds•

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

10-4
10-5
10-6

42

X

45

10-7

B disk10

reacceleration?

-9
-4

10

-3

10

-2

10

-1

10

10

0

1

2

3

10 10 10
E [GeV]

10

4

10

5

6

10

•

10

7

46

−

−1

disk,cut
3 1/4
S

SB

X

-8

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]. Thebol⊙[%] 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 a

27

Index
•

IceCube Neutrinos•

Accretion Flow in AGN
Murase, SSK, Meszaros, arXiv:1904.04226•••

AGN Corona model
LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

thermal
protonsso thatluminosity
! the
The
dotted line
represents
LX,obs= Lproton
X,calc , and cyan band3
ndicates
L@F
<"L
X,obs
X,calc < 1.7LF
X,obs
nt28
of theluminosity:
p/1.7
p , in which all thepaccretion
2fraction"p D+ Ḟp,inj ,reensquaresare +located. Fp
"p
Z@"p
tcool
tesc

C.

 �
p+p ! p+p+Aπ0+Bπ±�
p+γ ! p+p+π0�
p+γ ! p+n+π+�
π0 ! 2γ, π± ! 3ν+e±�γ�Power-law injection models ν� (Bp and C)

Basic Equations

Hot#Plasma
L"p d"p = ✏p ṁLEdd ,(1) (10)For models B and C, we considerπ a generic accelere momentum distribution function for
BH
3tion mechanism, and the steady-state
proton spectrump
where Fp is2 the 1 momentum distribution function
4⇡p
Fp2ptF/c),
Disthethedi↵usion
pdN/d""p isdi↵erential=="p dN/d"istheprotonluN"p = dN/d"p , is obtained by solving the transport equ
=
4⇡p
/c),
D
di↵usioncoefficient,
p
p
"
loss
turbulence.
We
computesteadystate
CRspectrabysolv1
is1 isthefor
protons,
tesc and1time
10-4
thet cooling
cooling
time,
tesc
is totaltheescape
time,
tion:
tcool
=
+
t
is
the
loss
rate)
and
RQ AGN  (this
work)
esc
loss • cooling the following
Fokker-Planck
equation(SA)(e.g., [75–78]),•
Power-law
Injection(PL)˙
Stochastic
AccelerationFistheinjection
term
to
the
stochastic
acceleration.
p,inj and Ḟp,inj is the injection funcRQ AGN (cascade)  (this work)
ime,
!
" We
on-thermal
protonproduction
efficiency.
RQ AGN (thermal e) ◆
-5 ✓3 waves, theConsidering resonant
scatterings
with
Alfven
10
ε
1
∂
∂F
F
∂F
RL AGN (in clusters)  N"
p asthep
ion coefficientcanbe εwritten
Dp "pequation⇡ − p + Ḟ ,2 to solve
d
"p
hang
&coefficient
CooperpD= ismethod+Fdi↵usion
representedas[144–146]
ε
p
p,inj
p
RL AGN(w. blazars) 
p
2 ∂ε
N
=
Ṅ
, conside
(1
∂t
ε
∂ε
t
t
"injection"function.
q 2 2
where Ṅ"p ,inj is theWeppp−coolesc
p
p ,inj2 N steady
p
total
)and
"calculate
,
where
=
B/
4⇡m
c
is
the
time
evolution
until
d"
tcool
tesc
A
p
p
p
10-6p injection(1) power-law
withanexponential
cuto↵:✓ ◆2⇣2ity, H ⇡ where
R/2 Fisthescaleheight,r=⌘
L
Dε ≈ εp /tacc
p is the⇣c
VA CR distribution
rL q 2 2 function,D" ⇡"strength,(9)✓◆ sinj✓◆−1-7p,armor radius,⇣diﬀusionis
turbulentis thespace, t−1
=
t
+
10
H the
c coeﬃcient
H in energyppp−cool
"p
"p
−1
−1
−1
−1
−1
−1
Ṅ"p ,inj = Ṅ0
exp,(
tpγ +tBHof+tp−synthe total coolingrate, tesc = tfall +tdiﬀ
wer-law index
the isturbulence
power
"
"
p,cut
p,cut
-8
is
the
escape
rate,
and
Ḟ
is
the
injectionfunction
10p,inj
We set Ḟp,inj = Ḟ0 ("p"p,inj ) with
reacceleration?
(see Appendix [79]). The stochastic acceleration time is. Note thatthe value of " (R/c)(εhas/eBR)no 2−q , where V
where Ṅ-90 is the normalization factor, sinj is the ingiven by tacc ≈ η(c/VA )2p,inj
p
A
10
1
5
7
resulting isspectrum
as longandasη the"p,cut
is3 thecuto↵
energy.
the Alfvén velocity
is theinjecinverse of the turbu- tion spectral
10-4 10-3index,
10-2 10-1and
100 10102 10104 10106 10Escape: Infalllence
strength[80, 81].to
We SMBH
adopt qḞ= 5/3,
which is con- normalize the injection by
E [GeV]uch•lower
than
the
cuto↵
energy.
is0
R the recent MHD simulations [56], together
sistent withat the= ⌘dissipation
is in the coronae• condition
"p d"
p the
p ṁLEdd
Coolings:
pp
inelasticcollision,
photomeson
FIG. 3. EGB production
andZENB spectra in our RQ AGN core model.
with η = 10. L
Becauserate
1is expected to be proportional to LX , we assume that the
The data are taken from[86] (green), Nagoya
bal"p ṄSwift-BAT.
("p ,inj d"p = ✏p ṁLEdd
L"p = tloss"p dN/d"p is the proton lumi+
loon
[87]
(blue),
SMM[88](purple),COMPTEL[89](gray),
protonsynchrotron,
Bethe-Heitler
process(p+γ→p+e+e)
injection function
linearly
scalesthe1
1 as LX1. To explain
Fermi-LAT[90](orange),and
IceCube[5]
forshower(black)44
njection parametertloss required= tcoolfor+tLesc
ENB, the CRandpressureerg s−1
X = 10and upgoing muon track (blue shaded) events. A possible
Muon&toChang-Cooper
Pion
Coolings
arenegligiblyWe
can get an analytic solution of the transport equatturnstheoutbe ∼ 1%of the thermalpressure,
which is inefficient
rate.• Weuse
methodcontribution of reaccelerated pairs is indicated (thin solid).
−1
4 3
−1
reasonable. We plot εp Lε ≡ 4π(εp /c )Fp V(tesc +tp−cool ) (cf., [149]):
usion• equation
[56,
57].
in Fig.
2, where
V is thevolume.
HEγ-raysareabsorbedby target photons (γ+γZ → e++e-)
Whilegeneralize
the CRs are theaccelerated,
they interact withand C, weaccelera1
t
cool
0
0
0
and radiation
modeledin theprevious
section,→matter
electron&term,
positronemithigh-energy
gamma-rays
N
=
d"
Ṅ
("
)
exp
G("
,
"
"
p,inj
p
a power-law
injectionṄ=ppp) , (
p
"
,inj
AGNs.
Results
are
shown
in
Fig.
3.
Our
RQ
AGN
core
p
"p " pand produce secondary particles. Following Ref. [82, 83],
model
can
explainthe ENB at ∼ 30 TeV energies if thejsolvethewherekineticelectro-magnetic
equationstakinginto accounteleccascades
exp( →"pwe/"Calculate"cut is thecutp,cut ),R In this work, secondary injections CR pressure is ∼ 1% of the thermal pressure.
tromagnetic cascades.
Ṅ0 is normalized
by " Ṅ"p ,inj
d"p = are approxIn the vicinity of SMBHs,Z high-energy neutrinosby the Bethe-Heitlerp andpγ processes
0
"2
d"
2
BH
are
produced
by
both
pp
and
pγ
interactions.
The
t
cool
p
s obtainedby equating
infalltime,
imatelytreated astheεe (dṄe /dε≈e )|ε =(m /m )εE2  [GeV cm-2 s-1 sr-1]

p

p

p

p

p

e

e

p

p

Turbulent#Field##!#proton#accelera4onWe consider high-energy neutrino emission from the29
We consider a RIAF of size R andvicinityof supermassive black holes (SMBHs) in active
rate Ṁ around a SMBH of mass MBH .
galactic nuclei (AGNs) [72–78]. A luminous AGN hosts
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 [82–14 et al. 2015
R = RRS ' 2.95 ⇥cf.
10SSK
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 ratetion rate is normalized by the Eddington
has not been supported bynumerical
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 accretionrate
2
Low-luminosity AGNs (LLAGNs), however, have dif1
1/2
9
1
C
⇡
V
'
3.4
⇥
10
R
cm
s
s
K
erent spectral energy distributions, in which an UV
1
2bump is absent [88]. This indicates that there is an op1•
Poorofobservationaldata
H ⇡ R ' 1.5 ⇥ 1014 R1 M8 cm
tically thin,Low-luminosity
hot accretion flow→
insteadan optically2
thick disk. →Remarkably,plasma properties
hot AGN
Formulation
basedof ontheoryṀ
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 M8

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30

Rates & CR Spectrum
•
SA

PL

•
•

Infall is dominant
Neutrino is mainly
produced by pp
For SA, acceleration rate is much
lower than PL model �
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

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

31

HE Particle Spectrum•

Stochastic Acceleration (SA)NGC 4258
Thermal γ•

6Power-law Injection (PL)NGC 4258
Thermal γ
p

p

νμ

νμ
Cascade γ

Cascade γ
pγ-νμ

pγ-νμ

FIG.
3. Spectral energy
distributionsof gamma-ray(dashed by thermal
electrons; dott• pp-neutrino
dominant,pγ contribution
is comparable
at high-energy
and proton (dot-dashed) fluxes for Model A (stochastic acceleration; left panel), Mod• Cascade γ-rays @ E<1 GeV: too faint to detect by LAT & CTA
middle),and Model C (power-law injection with sinj = 2.0; right), respectively. The
NGC3516, NGC
4258,@
andNGC 3031,
respectively.
thin-dot-dashedlines in mid• Thermal
γ-rays
E<MeV:
detectablebyThefutureMeV satellites
spectrum of the protons. The thin-solid lines are the p neutrino flux. For NGC 3031
end of the figure.

32

Number of Muon Tracks

FIG. 5. The expected number of through going track events from3031 (right panel) for models A (solid), B (dashed), and C (doand for IceCube (thin lines). The dot-dashed lines show the exp

Murase & Waxman
11 2016

A Bright LLAGN (NGC 4258)

Stacking 10 brightest LLAGNs
backgroundbackground

IceCube-Gen2
IceCube-Gen2

IceCubeIceCube• IceCube cannot distinguish signals
IceCubecannotfrom
detectneutrinogh• going
track events
NGCa3516(left panel), NGC 4258 (middle panel), and NGC
from thebackground(thick lines)), •B (dashed), and C (dotted) with a ten-year operationof IceCube-Gen2
IceCube-Gen2 can detect• IceCube-Gen2 can detectdashed lines show the expected background.

a few neutrinos of E > 10 TeV

several neutrinos of E > 30 TeV

8.68
33
9.14

268.77 -1.02 1.30
452.72 -0.57 0.91

3.25
2.14

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

γ�

> 1.5 ⇥ 104 GeV
26, 67])
Hot#Plasma
ZZ
eV for ṁ ⇠ 10 3 ,BHcdz
L"i ⌧i,IGM
um appears at apdLH↵ ⇢H↵
e
,
i =
4⇡H0
"i
(1 + z)3 ⌦m + ⌦⇤ectral index and(3)
ls B and C, the41 erg/s
where⇢H↵
the H↵
luminosity
function,contribute⌧i,IGM is to ν
LHα
< is4x10equally
entical•to AGNs
the in- withthe optical depth in intergalactic medium, and we use
minates over the
41
1
1 erg/s mainly contribute to MeV• LLAGNsHwith
L
~
4x10
Hα
⇠
70
km
s
Mpc
, ⌦M ⇠ 0.3, and ⌦⇤ ⇠ 0.7.
0nge.
H↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡LLAGN
explain(⇢⇤ /L⇤ )/[(LH↵ /L⇤ )s1 + (LH↵ /L⇤•)s2 ],where ⇢⇤ 'can
4.11 ⇥
eson interactions
1
10 5 Mpc 3 , L⇤ = 3.26 ⇥ 1041 ergs, s1 = ν2.78,
TeV-PeVand MeV γ
We calculate theand s2 = 1.88. We extrapolate this luminosity funcng the formalism
bkgrdssimultaneously
tion to Lmin = 1038 erg s 1 , below
which thePalos by p interacmar survey finds a hint of a flattening
Theattenuatedsur• GeV[69].
γs areon given in Ref.
vey also indicates a correlation between LX and LH↵ fortic field strength
at RIAFs
in LLAGNs
LLAGNs: LX ⇡ 5 7LH↵ [69]. We
use a correctione e↵ect of meson
factor X/H↵ = LX /LH↵ = 6.0. Then,the luminosity→ consistent
with Fermi data
neutrinos. Then,integration is performed in the range of 1038 erg s 1 ) = (1, 2, 0) at
~ 10
0.1P
41 th for
LH↵  ⌘rad ṁLEdd /(X/H↵ bol/X• ) 'PCR
4.2 ⇥erg s 1 .SA,o the neutrino osPCR ~redshift
0.4Pthevolufor PL
Since dimmer AGNs tend to have weaker
onic interactions
tion [70–72], we assume no redshift→evolution
the lu- spectrum
Needof hardn/positron pairs,
minosity function. The mass of SMBHs in local Seyferts. We calculate
galaxies does not show any correlation with X-ray lu-

ν�

p

π

p

Turbulent#Field##!#proton#accelera4on

8.68
34
9.14

268.77 -1.02 1.30
452.72 -0.57 0.91

3.25
2.14

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

γ�

ν�

p

> 1.5 ⇥ 104 GeV
26, 67])
Hot#Plasma
ZZ
eV for ṁ ⇠ 10 3 ,BHcdz
L"i ⌧i,IGM
um appears at apdLH↵ ⇢H↵
e
,
i =
4⇡H0
"i
(1 + z)3 ⌦m + ⌦⇤ectral index and(3)
ls B andC,the
41 erg/s equally contribute to ν
• LLAGNswhere
with
L
<
4x10
Hα
⇢
is
the
H↵
luminosity function, ⌧i,IGM isH↵
entical to the inthe optical depth in intergalactic
medium, and we use
41
minates• overthe
LLAGNsHwith
LHα~s 4x101
1 erg/s mainly contribute to MeV
⇠
70
km
Mpc
, ⌦M ⇠ 0.3, and ⌦⇤ ⇠ 0.7.
0nge.
H↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡(⇢⇤ /L⇤ )/[(LH↵ /L⇤ )s1 + (LH↵ /L⇤ )s2 ], where ⇢⇤ ' 4.11 ⇥• LLAGNeson interactions
can explain
5
3
41
10
Mpc , L⇤ = 3.26 ⇥ 10 erg s 1 , s1 = 2.78,
We calculate the
TeV-PeVν and MeV γ
and s2 = 1.88. We extrapolate this luminosity
funcng the formalism
tion to Lmin = 1038 erg s 1 , below which
Paloobs. thesimultaneously
s by p interacmar survey finds a hint of a flattening [69]. The suron given in Ref.
• GeV
vey also indicates a correlation between
LX andγsLH↵
areforattenuatedtic field strength
LLAGNs: LX ⇡ 5 7LH↵ [69]. We use
a correction
at RIAFs
in LLAGNse e↵ect of meson
factor X/H↵ = LX /LH↵ = 6.0. Then, the luminosity
neutrinos. Then,consistentintegration is performed in the range of→1038
erg s 1  with Fermi data
) = (1, 2, 0) at
LH↵  ⌘rad ṁLEdd /(X/H↵ bol/X ) '•4.2 ⇥ 1041 erg s 1 .
Po the neutrino osCR ~ 0.01Pth for PeV ν,
Since dimmer AGNs tend to have weaker redshift evoluonic interactions
PCR ~of0.1P
tion [70–72], we assume no redshift evolution
the thlu-for TeV ν
n/positron pairs,
minosity function. The mass of SMBHs in local Seyferts. We calculate
galaxies does not show any correlation with X-ray luπp

Turbulent#Field##!#proton#accelera4on

35

Index
•

IceCube Neutrinos•

Accretion Flow in AGN
Murase, SSK, Meszaros, arXiv:1904.04226•••

AGN Corona model
LLAGN RIAF model
Summary

SSK, Murase, Meszaros, 2019, PRD, 100, 083014
SSK, Murase, Meszaros in prep.

36

Summary

Accretion flows in AGNs are a feasible neutrino source• Seyfert galaxies can reproduce TeV-PeV νs without violating
Fermi constraints with observation based parameters.
• RIAFs in LLAGNs can explain MeV γ & TeV-PeV νbackgrounds simultaneously without violating Fermi constraints
ns (thick-solid lines), neutrinos (thick-dashed lines), photons by thermal electrons
c cascades (thick-dot-dashed lines). The thin dotted lines are the p neutrino• Future
observations
can robustly test both models:og L = 41.50(ṁ ⇠ 10 multi-messenger) and log L = 39.68(ṁ ⇠ 10 ), respectively.
Left andpectively. For models B and C, rhe injection proton spectrum (thin-solid lines)- IceCube-Gen2 can detect AGNs as point sources
3- AMEGO can detect MeV γ rays from AGNs•

X

3

edium (IGM),
turbulence. We compute steady state CR spectra by solvrays owing ing
to the following Fokker-Planck equation (e.g., [75–78]),
E , as shown
!
"
3
ε
1 ∂
∂Fp
Fp
∂Fp
p
ε2p Dεp= 2+
Fp −+ Ḟp,inj ,
∂t
ε
∂ε
∂ε
t
t
p
p
p−coolescpower luminosi(1)ies. Ref. [76]2where Fp is the CR distribution function, Dεp ≈ εp /tacc
To reproduce � γ�−1is the diﬀusion
coeﬃcient in energy space, t−1
lowernumberp+p ! p+p+Aπ+Bπν�
p−cool = tpp +p+γ ! p+p+π> 0.1 tis
−1
−1p
−1
−1
−1
−1
of
⌘p p+n+π⇠p+γ !
pγ +tBH +tp−syn is the total cooling rate, tesc = tfall +tdiﬀ
π ! 2γ, π pres! 3ν+e
smic-rayis the escape rate, and Ḟp,inj is the injection function
Hot#Plasmahich is at odds
π(see Appendix[79]). The stochastic acceleration time isfield generated
given by tacc ≈ η(c/VA )2 (R/c)(εp /eBR)2−q , where VA
BH
mic rays. Also, pTurbulent#Field##is the !#proton#accelera4on
Alfvénvelocity and η is the inverse of the turbuthe cosmic-ray
lence strength [80, 81]. We adopt q = 5/3, which is conrticles, leading
0

±�

0�

+�

0

±

±�

10-4

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

2

X

10-5
10-6

RQ AGN  (this work)
RQ AGN (cascade)  (this work)
RQ AGN (thermal e) 
RL AGN (in clusters) 
RL AGN (w. blazars) 
total

10-7
10-8
reacceleration?

10-9
10-4 10-3 10-2 10-1 100 101 102 103 104 105 106 107
E [GeV]37

Thank youfor your attention

38

Active Galactic Nuclei(AGNs)•

Accretion onto Supermassive black hole (MBH~108Msun)
gravitational energy → radiation or thermal energy•

SMBH paradigm is proved by Event Horizon Telescope
The Astrophysical Journal Letters, 875:L1 (17pp), 2019 April 10

maximum likelihood (RML; e.g., Narayan & Nityananda 1986;
Wiaux et al. 2009; Thiébaut 2013). RML is a forward-modelingapproach that searches for an image that is not only consistent withthe observed data but also favors speciﬁed image properties (e.g.,smoothness or compactness). As with CLEAN, RML methods
typically iterate between imaging and self-calibration, although
they can also be used to image directly on robust closure quantities
immune to station-based calibration errors. RML methods have been
extensively developed for the EHT (e.g., Honma et al. 2014;
Bouman et al. 2016; Akiyama et al. 2017; Chael et al. 2018b; seealso Paper IV).
Every imaging algorithm has a variety of free parametersthat can signiﬁcantly affect the ﬁnal image. We adopted a twostage imaging approach to control and evaluate biases in thereconstructions from our choices of these parameters. In
the ﬁrst stage, four teams worked independently to reconstruct
the ﬁrst EHT images of M87* using an early engineering data
release. The teams worked without interaction to minimize
shared bias, yet each produced an image with a similar
prominent feature: a ring of diameter ∼38–44 μas withenhanced brightness to the south (see Figure 4 in Paper IV).
In the second imaging stage, we developed three imaging
pipelines, each using a different software package andassociated methodology. Each pipeline surveyed a range ofimaging parameters, producing between ∼103 and 104 images

The EHT Collaboration et al.

EHT 2019

39

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

94.73 -1.93 2.32
40
159.56-1.48 1.79
268.77 -1.02 1.30
452.72 -0.57 0.91

⇥ 104 GeV •m ⇠ 10 3 ,ears at andex andnd C, theto the in- •over theteractions
culate theormalisminteracn in Ref.
strengthof mesonos. Then,2, 0) atutrino osteractions
ron pairs,

5.16
4.04
3.25
2.14

40.70/40.52/41.2 15.3/10.2/51.6
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

Luminosity Function

ν & γ intensities
from LLAGNs
Z
Z
26, 67])

c
i =
4⇡H0

dz
p
(1 + z)3 ⌦m + ⌦⇤L"i ⌧i,IGM
dLH↵ ⇢H↵e,"i(3)⇢H↵ lineis theLuminosityH↵ luminosityfunction
function, ⌧i,IGM isρwhereHα: Hαthe optical depth in intergalactic medium, and we use
H0 ⇠ 70 km s 1 Mpc 1 , ⌦M ⇠ 0.3, and ⌦⇤ ⇠ 0.7.
H↵ luminosity function is given by Ref. [68]: ⇢H↵ ⇡(⇢⇤ /L⇤ )/[(LH↵ /L⇤ )s1 + (LH↵ /L⇤ )s2 ], where ⇢⇤ ' 4.11 ⇥Hao 2005
10 5 Mpc 3 , L⇤ = 3.26
⇥ 1041 erg s 1 , s1 = 2.78,
We use higher LF
and s2 = 1.88. We extrapolate this luminosity function to Lmin = 1038 erg s 1 , below which the Palomar survey
finds
a 2007
hint of a flattening [69]. The surGreene& Ho
vey also indicates a correlation between LX and LH↵ for
LLAGNs: LX ⇡ 5 7LH↵ [69]. We use a correction
factor X/H↵ = LX /LH↵ = 6.0. Then, the luminosity
integration is performed in the range of 1038 erg s 1 
LH↵  ⌘rad ṁLEdd /(X/H↵ bol/X ) ' 4.2 ⇥ 1041 erg s 1 .
Since dimmer AGNs tend to have weaker redshift evolution [70–72], we assume no redshift evolution of the lu-

41Target Photon Field
•
•
•

•

Low-luminosity → Poor observational data
→ Theory driven formulation
Thermal electrons in RIAFs emit seed photons
Provide X-ray luminosity by observation→ Bolometric correction based on AGN survey→ Estimate mass accretion rate
→ Obtain physical quantities (ρ, B, n, Te) in RIAFs→ Calculate target photon spectrum by one-zone approximation
We do not adjust X-ray luminosity

42

EM Cascades in IGM
•
•
•

•

Cutoff energy by γγ pair production in RIAFs
Ecut ~ 0.1 - 100 GeV
τγγ,IGM ~1 for Eγ =100 GeV @ z=0.5→ We need to consider attenuation
γγ pair production: e+e- of γe~105 for Eγ~100 GeV→Eic ~ 4γe2 ECMB/3 ~ 10 MeV→ we can ignore EM cascade in IGM
τ << 1 for ν & MeV γ→ we can ignore attenuations

43

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