High-energy Emissions
from GW sources
Penn State (IGC Fellow)

Shigeo S. Kimura
References
1) Kimura, Murase, Meszaros, Kiuchi, 2017, ApJL, 848:L42) Kimura., Murase, Meszaros, Bartos, in prep3) Murase, Toomey, Fang, Oikonomou, Kimura et al. 2018, ApJ, 854, 60
Collaborators
Peter Meszaros, Kohta Murase, Michael W Toomey (Penn State)
Kunihito Ioka, Kenta Kiuchi (Kyoto University)
Imre Bartos (University of Florida)
Ke Fang (University of Maryland)
Oikonomou Foteini (European Southern Observatory)
Hotokezaka Kenta (Princeton University)
Kashiyama Kazumi (Tokyo Unversity)

2

Outline•Introduction•

High-energy Neutrinos from SGRBs

SSK, Murase, Meszaros, Kiuchi 17•

Sub-photospheric Neutrinos from NS mergers
SSK, Murase, Meszaros, Bartos in prep•Long-duration HE-photons from NS mergers
Murase, +, SSK et al. 18•

Summary3

Introduction

15

–2

4

60

–5

Seconds–5

Seconds
Trigger 1606

Trigger 1425
40

40
Seconds
Trigger 1974

40

40

20

10

0
–5

30
20

10

20

20
10

20

–5

150
Seconds

Seconds

Binary neutron star mergers—> sources of gravitational wave & γ-ray
Remnant black hole (or magnetar) launches a relativistic jet—> internal dissipation produces high-energy particles
800

103 Counts/Second

60

40
20

0
–20

400
200

Seconds
40

103 Counts/Second

103 Counts/Second

くくGravitational Wave

25
20

Seconds

30
20

10

40

0.5 sec

60

40
20
0
0.0

0.5
Seconds

Fig. 20. Sample lightcurves of GRBs.
Source: [5].

eCentral Engine

Trigger 3152
80

Seconds

BH or NS
Jet Launch

Trigger 2812

0
–2

100

5

Seconds

Internal
Dissipation

10
–10

50

0
–2

5

103 Counts/Second

Trigger 2571

15

100

Seconds

30

BNS mergers

150

600

0
–2

80

Trigger 2514

Trigger 2151

103 Counts/Second

Trigger 2067

80

103 Counts/Second•30

see e.g. Berger
2014
0

0
–10

Seconds•

103 Counts/Second

103 Counts/Second

103 Counts/Second

Short Gamma-ray Bursts
30

Fig. 23 gives an example of GRB 990123 whose time integrated spectrum is well fit by the Band function [432].
The Ep distribution of GRBs is wide. While bright BATSE GRBs (a sample of 156 bursts with 5500 spectra) have Ep clu
around 200–300 keV range [433], lowereEp bursts are found by softer detectors such as HETE-2 and Swift. The distriof Ep seems to form a continuum from several keV to the MeV range, e.g. [434]. From hard to soft, bursts are som
also vaguely classified as gamma-ray bursts (GRBs, Ep > 50 keV), X-ray rich GRBs (XRGRBs, 30 keV < Ep < 50 keV
X-ray flashes (XRFs, Ep < 30 keV), with no clear boundaries in between [435]. For the bright BATSE sample, the two sp
indices have a distribution of ↵ ⇠ 1 ± 1 and⇠ 2+12 [433]. Such a distribution is also confirmed for the Ferm
INTEGRAL bursts [103,436,434].
Spectra for some GRBs can be fitted with a cutoff power-law spectrum, in the form

ObserverX-rays・γ-rays

N (E ) = A✓

E
100 keV

◆

ˆ

exp

✓

E
Ec

◆

5

Afterglow of SGRBslog(L [erg/s])

Example of
SGRB Light Curve

1051
Extended Emission(for SGRBs)

1049

see e.g. Nakar 07, Sakamoto+11

Plateau Emission
1048

promptBurst

X-ray flare

Dashed line:
Standard classical afterglow
1

•
•
•
•

100

103

X-rays/opt/radio
Internal
Dissipation
e

Jet

e

ForwardShock

Observerlog(T[sec])
105

Prompt emission is followed by afterglowsStandard afterglow: Forward shock model, power-law decay
Extended Emission, plateau emission, X-ray flares have
similar features to prompt bursts —> Late-time engine activity?
Late time activities have comparable total energy to prompt burst

th
6

th

GRB Neutrinos

FIG. 8. The multiplicity of photons and leptons produced in
one interaction of a relativistic proton with 2.7 CMBR.

p

π0

μ+

π+

νμ

•

i.
th

p

e+
ννee
Eγ’[GeV]

Cosmic rays in the GRB emission
regionFIG. 9. The total cross sections of production of $ andas a function of energy of the incident gamma ray$ mesons—> neutrino emissions throughpγ interactionin the rest frame of a proton. The experimental points are takenfrom http://wwwppds.ihep.su:8001.
Peaky cross section:
Target γ ~ 200 keV —> neutrino energy
10 TeV OF
IV. ~PRODUCTION
ELECTRON-POSITRON PAIRS
Neutrino detection: Evidence of
hadronic Cosmic-rays
At energies below the photomeson production, the main
0

•

T

ν
νμμ
n

•

Kelner 08

γ
γ
γγ

p

＜
〜
〜

target γ
HEP

σ[mb]

Photomeson production (pγ)•

Waxman 97

w
u
"
a
p
in
a

þ

channel of inelastic interactions for protons with ambient
photons proceeds through the direct production of

w

a
v
d
&

7

IceCube GRB Analysis

Figure 6. Energy PDFs and signal-to-background ratios for the northern hemisphere (left) and southern hemisphere (right) nm track analyses. Left verticalaxis:reconstructed muon energy PDFs of background off-time data (black points) and E -2 nm signal simulation (blue line); simulated background used for PDF
extrapolation is provided in the northern track analysis (green line). Right vertical axis: per-bin PDF ratios (red points) and spline ﬁt (red line).

IceCube 17•

Figure 7. Differential median sensitivity of the northern hemisphere track, allsky cascade(Aartsen et al. 2016a), and southern hemisphere track stacked
GRB analyses to a per-ﬂavor E -2 ν quasi-diffuse ﬂux in half-decadal ν energybins, with the ﬁnal combined analysis shown in the black line. Integrated
sensitivities are shown as dashed lines over the expected 90% energy central
interval in detected neutrinos for a given analysis. The IceCube measured 68%
CL astrophysical per-ﬂavor neutrino ﬂux band is given for reference from a
global ﬁt of IceCube analyses(Aartsen et al. 2015a) and a recent six-year
northern hemispheres nm track analysis (light blue, Aartsen et al. 2016c).

uncertainties in calculated model limits is determined in amodel-dependent way, as their effect is found to be much more
pronounced at low energy than at high energy.
Constraints were ﬁrst determined for a generic double
broken power-law neutrino ﬂux of the form
⎧ e-1E -1,
En - eb⎪ b n
Fn (En ) = F0 ´ ⎨ En-2,
eb < En - 10eb
⎪ -4
⎩ En (10eb )2 , 10eb < En

as a function of ﬁrst break energy eb and quasi-diffuse spectral
normalization F0 . These limits are presented in Figure 8 as
excluded regions in this parameter space. Two models ofneutrino production in GRBs where GRBs are assumed to be
the sole origin of the measured UHECR ﬂux are provided inthis parameter space: the neutron escape model of Ahlers et al.
(2011) and the proton escape model of Waxman & Bahcall(1997), which has been updated with recent measurements of
the UHECR ﬂux(Katz et al. 2009). Both models are excluded
at over 90% conﬁdence level (CL) with most of the model
assumption phase space excluded at over the 99% CL. A
thorough reconsideration of whether GRBs can be the sourcesof UHECRs from Baerwald et al. (2015) shows that the internal
shock ﬁreball model is still plausible if cosmic-ray protons can
efﬁciently escape the ﬁreball with a low pion-production
efﬁciency for a range of fp and Γ, which predict neutrino ﬂuxes
below the current limits.
Similar constraints were calculated for simple power-law
spectra consistent with IceCube’s observed astrophysicalneutrino ﬂux (Aartsen et al. 2014c, 2015b, 2015c, 2016c),
concluding that 10.4% of the astrophysical neutrino ﬂux canbe the result of a GRB prompt, quasi-diffuse ﬂux assuming no

Using the timing and position information of each GRB,
IceCube put the limit on GRB associated neutrinos—> GRB contribution < 1%published searches, these models are expected to yield 6.51,
11.02, and 0.25 neutrino events, respectively. Though anumber of events have been found temporally coincident with
GRBs, none haveappeared to beparticularly compelling
signals and they have occurred at a rate consistent withbackground.
Having found results consistent with background, limits canbe placed on neutrino production models in GRBs. These

(8 )

8

IceCube GRB Analysis

Figure 6. Energy PDFs and signal-to-background ratios for the northern hemisphere (left) and southern hemisphere (right) nm track analyses. Left verticalaxis:reconstructed muon energy PDFs of background off-time data (black points) and E -2 nm signal simulation (blue line); simulated background used for PDF
extrapolation is provided in the northern track analysis (green line). Right vertical axis: per-bin PDF ratios (red points) and spline ﬁt (red line).
The Astrophysical Journal, 843:112 (13pp), 2017 July 10
Aartsen et al.

IceCube 17

Figure 7. Differential median sensitivity of the northern hemisphere track, allFigure
Effective areas,et scaled
to all-sky,of the northern
and tracksouthernsky 2.
cascade(Aartsen
al. 2016a),and southernhemispherestackednm track
comparedto that of the ﬂux
all-skycascade analysis
hemisphereE -2 ν quasi-diffuse
in half-decadalν energy
GRB analyses
to analysesa per-ﬂavor
for thedetector conﬁguration.
bins,79-stringwith theIceCube
ﬁnal combined
analysis shown in the black line. Integrated

uncertainties in calculated model limits is determined in amodel-dependent way, as their effect is found to be much more
pronounced at low energy than at high energy.
Constraints were ﬁrst determined for a generic double
broken power-law neutrino ﬂux of the form
⎧ e-1E -1,
En - eb⎪ b n
Fn (En ) = F0 ´ ⎨ En-2,
eb < En - 10eb
⎪ -4
⎩ En (10eb )2 , 10eb < En

(8 )as a function of ﬁrst break energy eb and quasi-diffuse spectral
normalization F0 . These limits are presented in Figure 8 as
excluded regions in this parameter space. Two models ofneutrino production in GRBs where GRBs are assumed to be
the sole origin of the measured UHECR ﬂux are provided inthis parameter space: the neutron escape model of Ahlers et al.
(2011) and the proton escape model of Waxman & Bahcall
Figure 3. Signal-to-background PDF ratios for three GRB durations. The(1997), which has been updated with recent measurements ofearliest reported start time T1, and the latest reported stop time T2, deﬁne the
UHECR. al. 2009). Both models are excluded
mosttheinclusive
GRB ﬂux(Katz
duration T100et
at over 90% conﬁdence level (CL) with most of the model
assumption
excluded atatover
the 99%CL. Anˆs ) divided
by themaximizedwith phaserespectspace
to ns (maximized
thorough reconsideration
whetherGRBs
can be the0), whichbackground-only
likelihood (nofs =simpliﬁes
to sourcesof UHECRs from Baerwald et al.
(2015) shows that the internalN⎤⎡ nˆs + (xi )⎡ $ (nˆs ) ⎤shockﬁreball
model
is
still
plausible
if
cosmic-ray
protons, = ln ⎢+ 1⎥. (5) can⎥ = - nˆs + å ln ⎢⎣ $ (nsescape(xi ) pion-production= 0) ⎦ the ﬁreball⎦⎣ ánabñ low
efﬁciently
i = 1 with
efﬁciency for a range of fp and Γ, which predict neutrino ﬂuxes
The average expected numberof background events can be
below the current limits.
determined from off-time data, denoted as ánbñ.
Similar constraints were calculated for simple power-law
The time component of the signal and background PDFs,
spectra consistent with IceCube’s observed astrophysicalshown as a signal-to-background PDF ratio in Figure 3, isneutrino ﬂux (Aartsen et al. 2014c, 2015b, 2015c, 2016c),
deﬁned by the T100 of each burst. The signal time PDF is
concluding that 10.4% of the astrophysical neutrino ﬂux canconstant during T100, with Gaussian tails before and after thebe the result of a GRB prompt, quasi-diffuse ﬂux assuming no

Using the timing and position information of each GRB,
These
analyses
focus
on
the
promptphase
sensitivities
are
shown
as
dashed
lines
over
the
expected
90%
energycentralIceCube
put
the
limit
on
GRB associated neutrinos
interval
for a givenarea
analysis.
The IceCube
measured 68%
whereis the effectiveof neutrino
interaction
A in(detected
E , W¢)neutrinos
CL astrophysical per-ﬂavor neutrino ﬂux band is given for reference from afor global
an—>
event
signalneutrino
ﬂux,
F (cannot
E , W¢) isettheGRBbeasourceobserved
neutrinosﬁt of selection,
IceCubeanalyses(Aartsenal.
2015a)
anda recent
six-yearafterglow
phase
is ofnotconstrained
andnorthern
the integral
is performedover(lightthe analysissolidΩ
n track analysis
blue, Aartsen
et al.angle
2016c).
hemispheres
•

eff

n

n

m

n

and neutrino energy En range. The effective areas, scaled to allsky, of the northern and southern hemisphere track selections
published searches, these models are expected to yield 6.51,
are11.02,
shownandin Figure2, compared
the all-sky cascade
0.25 neutrino
events,to respectively.
Though aselectionofAartsenetal.
(2016a).
Thenorthernhemisphere
number of events have been found temporally coincident withselection
demonstrated
to be mosttosensitive
to neutrinos
withGRBs,is nonehaveappearedbeparticularlycompellingenergies1
PeV,they
while
effectiveatareaof thesouthernwith
signals andhavetheoccurreda rateconsistenthemisphereselection displays the enhanced sensitivity of thisbackground.
Havingfound results
limits canchannelto neutrinos
above consistenta few PeV.with
Thebackground,resonant scattering
be placed on neutrino production models in GRBs. These

SGRBs are minority—> constraint is not strong

The Astrophysical Journal Letters, 848:L12 (59pp), 2017 October 20

9GW170817
LIGO+ 17, J-GEM 17•

The first detection of
NS-NS merger event
by GW, radio, IR/opt/UV, Xray, MeV γ-ray
Talk by Tanaka-san, Sakamoto-san, Mori-san,Hachiya-san, Kawabata-san, Kokeyama-san•

Do NS-NS mergers emit
high-energy γ-rays andneutrinos?
weeks~months

Ioka+ 17

X/Radio
afterglow
1

Jet-ISM shock~sec>yearsOff-axissGRB

X/Radioflare

1

EjectaEjecta-ISM
shock~1day opt.
macronova

Jet
~100

CocoonCv~0.2-0.4c
v~

v

Mer
Mergerejectaejec
v~0.03-0v~0.03-0.2c(dynamical/(dynamishock/wind)shock/w~10day IR
macronova

Disk

Fig. 1 Schematic figure of our unified picture.

A

10

11)

ecta

12)

fast

max

Me .
yer
sity
for
6⇥
t an
with

Choked or Successful?
Choked Jet: Weak γ-rays
Liso~1046 erg/s for on-axis observer•

Gottlieb+17

•

t al.
one
lug"
om
verdally
3D
last
o the
axis
f the
m
to
nitial
cm
ared
ngles
and

d the
s tailsimiation
s for
t is
eragere.
this
thetail’stactis
ity
jet
kout

e as
ftical1s
difore
ided
nto
o
the
oweunito
8 cm
the
8 cm.
that
up
ttedtoe exom
and
the
axes,
s and
of

Figure 1. Maps of the logarithmic energy density excluding the rest-massenergy (left) in c.g.s units and logarithmic four velocity (right). The upper figure is taken before the breakout of the forward shock from the coreejecta. Although the50
forward shock will break out, the jet material behindisothe reverse shock will remain trapped inside and will be choked with thetermination of the engine. The lower figure is taken when the shock breaks
out of the tail at q = 0.7rad at t = 6.2 s and r = 1.3 ⇥ 1011 cm. The shock
has a quasi-spherical shape, reaching most of the ejecta. (An animation isavailable in the online journal.)

Successful Jet: classical SGRBs
L >10 erg/s for on-axis observer
Gottlieb+17

constant and checked the effect of the tail by considering several
configurations (without doing an exhaustive parameter phase space
search). We examined tail parameters in the following ranges: thedensity power-law (5 15), total mass (10 4 5⇥10 2 ) M andmaximal velocity (0.5 0.85)c.
The outcome depends only on the parameters near the shockupon breakout, which are determined by these initial conditions.
The light curves we obtained showed a large range of observed values, yet almost all light curves showed the expected common feaFigureof5.low-luminosity
Same as Fig. 1 for(compared
the successfultojetthe
case.
Top:ejecta
A wellenergy),
collimatedlow
tures
total

Both models can reproduce
prompt γ-ray and afterglow
of GW170817
Choked scenario implies
new subclass of SGRBs

Kasliwal+17
Mooley+1811High-energy Neutrinos
fromshort gamma-ray bursts
1) Kimura, Murase, Meszaros, Kiuchi, 2017, ApJL, 848:L4(This work was done before the detection of GW170817)

12

Questions:
Can SGRB emit observable neutrinos?
Which component can emit observable neutrinos?
Can we expect simultaneous detection with GW?
log(L [erg/s])

Example of
SGRB Light Curve

1051
Extended Emission(for SGRBs)

1049

X-rays/opt/radio
Internal
DissipationePlateau Emission
1048

promptBurst

X-ray flare

Dashed line:
Standard classical afterglow
1

100

103

log(T[sec])
105

Jet

e

ForwardShock

Observer(Γrel − 1)Eiso
2 dN
Ep
≈
exp
dEp
ln(Ep,max /Ep,min )

13

"

−

Ep
Ep,max#,(13)

Multi-component One-zone model

where Eiso ≈ Liso tdur is the isotropic equivalent kinetic
KimuraDissipation Region @ R=Rdis
energy,Central Ep,max and Ep,min are γthe maximum and miniνμ
Lorentz factor ΓEngineenergyformulaeof the non-thermal
at theobservereau mumﬁtting
based p on
GEANT4(see
Muraseμ+ protons
νe 2respectively.
Γj Γrel mp c andrgy frame,
Nagataki
2006a). We useγ Ep,min
Kimura≈et al.
νμπObserver dep
the Ep,max Pions= Γjgenerated
εp,max is obtained
byphotomeson
the balanceproduction
betweenthroughtheﬁtting formulae based on GEANT4 (see Murase&oscillation
NeutrinoGO/theaccelerationandcooling,
i.e.,
t
≈
t
.
into
muons
and
muon
neutrinos.
Using
the
meson
producp,accp,clNagataki• 2006a).
Calculateν fluencefromeachcomponentby one-zone
model
f
º
t
t
fpg <
(whichalwayssatisﬁes
efﬁciency,
These
protons
produce
pions
that
decay
to
muons
and
p
,cool
p
g
Pions generated through ptheg photomeson production decay
5 injection:•
Power-law
protoninto muonsandneutrinos.
muon
neutrinos.
Using
the mesonproduction
thisdeﬁnition), themuon
neutrino
spectrumproduced
muon
The
muon
neutrino
spectrumby pionby p
2dN
E
/ln(E
/Ep,minfp)g < 1 infº pexpressed
tp/dE
t~pgξ(which
satisﬁes
efﬁciency,ppestimatedp Eγ,iso always
p,max
,cool
pg is
decay
is
to
be
decay
as
5
this deﬁnition), thecoolingmuon neutrino
spectrumproduced by&pion• Protonprocesses:
synchrotron
adiabatic coolings"#
dieddecay is estimated to be• μ and dN
π alsodown
bysynchrotron&
adiabaticcoolings
dN
1
1
dN
dN
π−νcool
1
n
p
p
µ
m
2
2
2
2
the Eπ−νµ
≈
f
+
f
f
E
.,
(14)
E
»
f
f
E
,
pγ
pp
π,sup
p
p
n
g
p
p
sup
m
dN m
dNp 6
1
2 dEnπ−ν
2
8
.g.,
Enm
» µ fpg fsup pdE
Epnm , 8(3)dEpdEp
dEnm
8
dEp
s of
−1 −1
−1 −1
where
f
=
t
/t
and
f
=
t
/t
are
the
neuE
»
0.05
E
f
=
1
exp
(
t
tp,dec ) is
where
and
pγ
pp
pγ
pp
ion
n
p
p
,cool
p,cl
p,cl
m
sup
p
where En » 0.05Ep and fsup p = 1 - exp (-tp,cool tp,dec ) is the+mPlateau30

103×103×100.1

0.3–10Quantities
14

B (G)

L g ,iso (erg s−1)

Eg,iso (erg)Ep, M (EeV)

En , m (EeV)

En , p (EeV

EE-modEE-optPromptFlare
Plateau

2.9×103
5.0×104
6.7×103
5.3×102
1.8×102

1.2×1049
3.4×1049
6.1×1051
3.5×1048
3.8×1047

3.8×1051
1.1×1052
6.1×1051
1.0×1051
1.1×1051

21
6.0
60
25
13

0.020
3.9×10−4
0.29
0.11
0.33

0.28
5.4×10−
4.0
1.5
4.6

Neutrino Fluence+horizontal muon neutrino events (d > -5 ), the atmosphmuons are shielded by the Earth. For IceCube-Gen2, we
ModelPlateauareasPrompt
10 2 3 timesEElarger effective
than those ofFlareboth upgo+horizontal and downgoing events for IceCube. The effec
Γarea of downgoing
10–30 muon30neutrino events
1000in IceCube-Gen230
not be simply scaled, but the simple scaling is sufﬁcient for
demonstrativeof 14this work.
We set3x10
the 14thresh
14
Rdis [cm]1013–10purpose
3x10
3x1013energy for neutrino detection to 100GeV for IceCube
for IceCube-Gen2.
Eγ,1TeV
pk
1—10
0.1
500
0.3
[keV]The probability of detecting k neutrino events, pk
described by the Poisson distribution. The detection probab
Eγiso.kof more thanas p (&501051 k neutrinos
3x1050 is represented
1051
3x10 m1 - åi < k pi . We ﬁnd that for EE-mod (G = 30 ), the p[erg]ability for upgoing+horizontal events, p (&m . 1), is 0.04
0.16 with IceCube and IceCube-Gen2, respectively. For EE
Figure 1. Neutrino ﬂuences from the EE-mod, EE-opt, prompt emission, ﬂare,(G = 10 ), &m  1.7 and 7.9 with IceCube and IceCube-G
Set
Mpc.dL = 300 Mpc (GW horizon for design sensitivity)and plateau for d•L=300
respectively. It is possible for IceCube to detect neutrinos f
EEs, while detections with IceCube-Gen2 are more promis•
Extendedemission(EE)
canproduceneutrinos efficiently
number density. This makes EEs more luminous than the
However, for dL=300 Mpc, the neutrino detection forothers. The magnetic ﬁelds are so strong that spectral breaks
prompt,ﬂare, fluenceand plateau φ↑neutrino emissions may still
•
Γ
↓
or
R
↓
—>
photondensity↑—>dis
due to both the muon and pion cooling supressions are seen inchallenging even with IceCube-Gen2, since p (&m . 1)
Figure 1. The proton maximum energy is determined by thethem is less than 0.01.
ν to the Lor
photomeson production, leading to relatively lower values of
The neutrino ﬂuence ofGRBs
issensitiveμ
Ep, M . For the other three models, fpg < 1 is satisﬁed and theνfactor. To take this effect into account in a reasonablemanν the detec
lower ﬂuences are obtained. The magnetic ﬁelds are so weakπ
we consider the distributionof Γ to calculate

en
can
be =
asneutrinos
brightas
to0.04high.wepion
wheresGF0=
is 2the(EEnormalization
In eachcase,
examine30(G0 0.42=isEEs10represented). owing
P (& .than
2) G0 k0.25
0.69
0.24 fp (&k) =of they
more
asthe
m
p (&for1)ﬂuencesfor(G = 30
eCube-Gen2,
since
m .
roduction
efﬁciency.
Also,
neutrino
fromprompt
115
p
-å
.
We
ﬁnd
that
EE-mod),theprob)
EE-opt-dist-B
IC(up+hor)IC(down)IC (all)(all) Lorentz factor,
IC ( A
is(EE-mod-dist-B
the Gen2
means=4
mod-dist-A
and
EE-opt-dist-A)
and
i
<
k
i
G
where
A
is
the
eﬀective
area.
Since
downgoing
eventseﬀ than the plateau and ﬂares if G 1 300 is
missioncanbe higher
6. 1), is 0.04 andabilityforupgoing+horizontalC)
P (& . 1)
0.60 events, p (&m 0.19
0.41
0.73 space. Here, 0.47
logarithmic
we in
and
EE-opt-dist-B).
suﬀer
from
the
atmosphericbackground,
we
focus
on
alized.
GRBs
isIceCube
to 0.31the Lorentz
P (& .
2) sensitive
0.02
0.18
0.55
0.17
0.16
with
and IceCube-Gen2,
respectively.
For EE-opt
G»2factor
,
below
which
w
min we use
The
resultant
P
are
shown
in
Table
2,
where
the
upgoing+horizontal
eventsthat
havedeclination
δ
>
k
&

1.7
10
),
and
7.9
with
IceCubeand IceCube-Gen2,(G =ntoaccountinareasonable
manner,m
Note. IC:
Gen2: IceCube-Gen2, up+hor: upgoing+horizontal events, down: downgoing events, all: does
covering-factor-weighted
averageover the up+hor&
andm
notexist.higher
We calculate◦3.IceCube,
ProbabilityofNeutrino
Detection
−5
.
We
use
the
eﬀective
area
shown
in
Ref.
[11]
d
=300
Mpc.
The
upgoing+horizontal
eventshave
respectively.
is thepossiblefor
IceCubeto detect neutrinos from
:Itusing
declination-averaged
effective area.
down, A
Lcalculate
on
of
Γ
to
the
detection
2/3 estimate the detection prob
2)EEs, whilefor
IceCube.
For
IceCube-Gen2,
we
use
10
times
detections
with
IceCube-Gen2
aremorepromising.
probabilitythanthedowngoing
eventsowing
to a higher AeffnThe
expected
numberof-induced
eventsisestimated
to
be
m
nt
and
futureneutrino
experiments.
detection
Since Mpc,
severalparameters
are uncertain,• Expected
numberofν
events:=300theneutrino
detection
fortheHowever,
forprobability.
dlarger
A
than
that
for
IceCube,althoughitlowerdependsL
P (s&Gm model. k ) = 1(EE- åi < k Pi . No
eﬀ
for
low
E
.
In
EE-mod-dist
cases,
the4
weconsidermoderate(EE-mod-dist)andoptimistic(EE-optν
tionisdist)toplateaubelognormal:
prompt,ﬂare,
andbasic
neutrino
emissions
may still beonassumed
themod-dist-A)
specific
configurations.
The
thresholdenergyforδ throughandfn and Aeff , resmodels.
Theparametersfor
EE-mod-dist(EE-opt&
=
f
A
(
d
,
E
)
dE
,
(
7
)
has
slightlylower
detection
probabilities,
becausemeffnnnIceCube-Gen2,
p (&mTeV. 1) forfor IceCube
challenging
even
with
since
dist)
are
the
same
as
those
for EE-mod(EE-opt)
withupgoing+horizontal and downg
theneutrino
detection
is
set
to
0.1
2
⎛
⎞
s
=
2
.
In
each
case,
we
examine
(EEG
=
30
(
G
=
10
)
they
have
a
smaller
fractionoflower-Γ
EEs.
On
the other hand,them isthanG
0 less(ln0 (G0.01.
G))0
consider a covering-factor-we
and
1 TeV
for IceCube-Gen2.
The
detection
probabil4 sensitive
mod-dist-A
andﬂuence
EE-opt-dist-A)
and sG =(EE-mod-dist-B
F
exp
,
(
8
)
⎜
⎟
0
The
neutrino
of
GRBs
is
to
the
Lorentz
Aandis theeffectiveThe
effectiveareasof upgoing• Detection
EE-opt-dist-A
hasdescribedhigherdetection
probabilities
than EE-opt)here
2area.
eff EE-opt-dist-B).
probability
ispoisson:
⎝
⎠
s
2
(
ln
(
))
ity
of
k
neutrinosisbythe
PoissondistriG shown
factor.
To
take
this
effectinto
accountina reasonable
manner,The
resultant
P
are
in
Table
2,
where
we
use
6
horizontal anddowngoing
tracksfor
IceCube
is
shown
in
k
kto a smaller fraction of higher-Γ
dist-B
due
EEs.
We
also
Althoughthe exact
shape of F (G) is
dbution,=300
Mpc.
The
upgoing+horizontal
eventshavehigherwe
considerthedistributionofΓ
tocalculatethedetectionLp=N
exp(−N
)/k!,
and
the
probabilityof¥k2)
look lognormal,
ratherfor
than Gaussian (artsen probability
et al.estimate(2017)
asPa eventsfunctionofa higher
Eν. AForupgoing"thanthedowngoing
owing
to
using
declination-averaged
effectivearea
eff
Fbyk(Gcurrent) dkcases,
ln andGthe=future1),s Gonfactor(ofthanprobability
EEs
neutrino
experiments.
0
more
neutrino
detectionsis
1 − i<k pk . We
forlow Eò
.
In
EE-mod-dist
lower
model
(EEG
ν Gmin
The Lorentz
factordistribution
is assumed
toA
belognormal:
A
IceCube,
,
shown
as
IC
(
)
in
A
=
d
W
mod-dist-A)
has slightly
lower
detection
probabilities,
becauseeff,aveeff,aveeffòcalculatethedetection
probability
formodelsA
andB,
3•
distribution
tor, Assume
andis fraction
the ofdispersion
they havesaGsmaller
lower-ΓofEEs.ΓOn theinother hand,2 ⎞ higher
which
istabulated
the
upper
part
of Table⎛P (&II.
1) for EE-optTable
2,
which
shows
slightly
EE-opt-dist-A
has
higherdetectionin
probabilities
thanEE-optG
G
dN
(
ln
(
))
m .Ice0G
we introducetheminimum
LorentzG) is= aunlikely=F
F (due(8 )⎜ - any
dist-Btosmaller fractionofhigher-Γ
EEs. Wealso⎟ ,
0 exp
Cube
to
detectcoincident
neutrino
sig2
A
dist.
Althoughthedeclination
dependence
of
does
not
⎝
⎠
G
s
d
ln
2
(
ln
(
))
eff
estimateP
using
declination-averaged
effective
area
for
G
k
bch we assume that such a slow jetmeff,avemmeff,ave

Detection Probabilityònals, whilecanneutrino
IceCube,shown as aIC( Aeff,ave ) inevent with IceCubeAeff,ave =we
dW
Aeff , expectò¥edwhereP (&m(.
1) for
EE-opt2, whichshowsslightly• m Table
Fneutron(G) d ln Gstar= 1),merger
F0 is
the
factorG0
Gen2
fornormalization
model
A.higher
Using
the
rate
G
min
dist. Although the declination dependence of Aeff does3 notc−3−1
obtained
by
LIGO,
R
∼
1.5
×
10
Gpc
yr
s
is thechange
mean
Lorentz
factor,andisthedispersion
in , -around2
our conclusion much, the declination-dependent
G
k
G
k
m Mpc
6 for more quantitative evaluations.
L year.
analysisisspace.
important
nm every170
mergereventshappenwithin300
logarithmic
Here,
we
introducetheminimum
Lorentz-2
Using the relation &m µ fnm µ dL , we estimate2
kbelowG»
2
factorPThe
,
which
we
assume
that
such
a
slow
jet
fractionofon-axiseventsisf∼
0.045θ
m
0
L , leadmin(&m . 1) = 1 - P0 as a function of dL, which is shownbin
j,0.3
−1 we
Figure2.
Here,weon-axis
ignore the effectsof cosmologicalredshift,&m forEEs
variousΓ,andexist.
We
calculate5)does not
k withing
toanmerger
rate
R
≃4.1
yr
within the
0 The
since we focus on the local universe at dL 1 2 Gpc.
Parea.
dG
FG pk and that
estimatetheprobabilities
upgoing+horizontal
coverage
Supposingverticaldotted detectionlines show dL=300
Mpc and dL=600
k =Mpc,L allwhich corresponds to the sensitivity ranges of face-on NS–NS

change our conclusion much, the declination-dependent& Estimate
for EEs withvariousΓ, andwethedetection
probabilities
ò
analysis is important for more quantitative evaluations.
P = ò dthe
GF prelation
probabilities Usingand & µ f µ d , we estimate. Note thatP (p& is.a1)function= 1 - Pof asΓ a function of d , which is shown in, respectively.
We2.calculate
for the effects of cosmological redshift,
FigureHere, we Pignore
owngoing events
andthewelocalsince separately,
we focus onò universe at d 1 2 Gpc. The

16

and NS–BH mergers by aLIGO, respectively. For NS–BH
dL.
mergers, since the distance is longer, P (&m . 1) is lower thanThe
those for NS–NS mergers. The detection probability of nearbysG , while
eventsis affectedby Letters,that(6pp),
of distanteventsThe
AstrophysicalJournal
848:L42017 October
10 is not.
Ac
We estimate the detection probability within a given timeN
Table 3window
of
E
acc
DT
(
=
1
P
,
which
is
estimated
to
be
,
where
N
interval,T
Detection
probability
in aΔT(DT 0
The Detection
Probabilities within
Given TimeDInterval,
ideally beins
is the number of EEs for the time interval within the covering
localization
NS–NS (DT = 10 years)
IC (all)
Gen2 (all)Hearea of neutrino detectors. The local SGRB rate Fermi
is GBM-1
EE-mod-dist-A
0.11–0.25~4 Gpc-3 yr-1 - 10 Gpc-3 yr(e.g., Nakar et 0.37–0.69
al. 2006; Wanfor the to
GW
EE-mod-dist-B
0.16–0.35
0.44–0.77
NSderman&Piran2015),
so
the
event
rate
withinthesensitivity
without
elec
EE-opt-dist-A
0.76–0.97
0.98–1.00-1 - 1.1 yr 1. atmosph
the
~
0.46
yr
~3
range
of
aLIGO
(300
Mpc)
is
EE-opt-dist-B
0.65–0.93
0.93–1.00

Detection Probability

the signal in
NS–BH (DT = 5 years)
IC (all)
Gen2 (all)atmospheric4
In the 2
EE-mod-dist-A
0.12–0.28
0.45–0.88
Einstein Te
EE-mod-dist-B
0.18–0.39
0.57–0.88CE), might
EE-opt-dist-A
0.85–0.99
1.00–1.00
EE-opt-dist-B
0.77–0.97
0.99–1.00
mergers eve
prakash et a
Wanderman & Piran 15, Nakar + 06
Note. The SGRB rateis assumed
to be 4 Gpc-3 yr-1 - 10 Gpc-3 yr-1.
gamma-ray-3-3• RSGRB~ 4 —10 Gpc yr & half of SGRBs have EE
also being p—> N ~ 2-5 for NS-NS (10 yr), N~9-22 for NS-BH (5 yr)z 2 1 with
operation. The estimated values of (DT are tabulated in Table 3.
• For optimistic case, simultaneous detection with GW
Since GW d
We ﬁnd that the simultaneous detection of gamma-rays,
cosmologica
is highly
probablewith
IceCubeneutrinos,and even
GWs ispossiblein the era of IceCube-Gen2
mergers andandaLIGO/aVirgo/KAGRA,
assumingacosmic-ray
loading• Fore moderate case, IceCube-Gen2 is likely to detect neutrinos
stacking ana
factor, x ~ 10 . This will allow us to probe the physical17

Implications for GW170817

neutrino conpical GRB paons, we use the
IceCube+ 17
ese to our conFor extended
rresponding tole 1 of Kimura
timescales, wee window with, namely emism 3 to 30 days
fiducial emisTeV
PeV
EeV
n in Fig. 2. We• There is no extended emission from this event
istic extendedMpc, viewed —>atneutrinos from EE should not observed

The jet may be off-axis —> the flux is considerably lower• This event is in southern sky—> atmospheric noise is strong for lower energy
m the first bi•

Ws, GW170817,

18

Sub-photospheric
Neutrinos from
Neutron Star Mergers
2) Kimura., Murase, Meszaros, Bartos, in prep

Preliminary

19

Questions:
Can the choked jet produces detectable neutrinos?
choked jet —>EMarehiddenbyejecta7—> neutrinos are only an available signal(11)

ejecta(12)

e fastvc,maxs Me .
layer
ensity8 s for2.6 ⇥at an
d withKasliwal+17
Choked Jet: Weak γ-rays
Liso~1046 erg/s for on-axis observer
Gottlieb+17
Mooley+18

et al.
st onefrom
icallyhe lastz-axiscm to
12 cm
s, andFigure 1. Maps of the logarithmic energy density excluding the rest-massenergy (left) in c.g.s units and logarithmic four velocity (right). The upper figure is taken before the breakout of the forward shock from the core

Figure 4. Radio light curves arising from quasi-spherical ejecta with a velocity gradient

20

Schematic Picture••

swept-up ejecta forms cocoon surrounding the jet—> push the jet inward —> form collimation shock
Velocity fluctuations induce the internal shocks at
the pre-collimated jet
cf. Murase & Ioka 13

Jet Head
Collimated
γ Jet

γ

Ejectap

Cocoon1010 cm

Collimation
shock
(CS)

p

internal shock
(IS)

21

Particle Acceleration

4

86

0.5

upstream
0.4

Transition
regionΓβ

BUDNIK ET

downstream

A
Se
of
de
R
sh
ex
R

kTe

Γβ
2
T/(mec )0.3

10 X P
0.2

γ
γ

0.1

Treumann+ 08

0Weaveranalytic

0

10

20

τ*

30

40

50

Budnik+10

Figure 45. Shock structure for ε = 50 MeV. The dotted black line is theanalytic solution for Γβ obtained by Weaver (1976, Equation (5.10)), withaverage Compton cross section σ̄C = 0.56σT .
mfp(A color version of this figure is available in the online journal.)

Particle acceleration requires sharp velocity jump in λ
FIG. 3: Cartoon of the di↵usive shock acceleration (left) and shock heating mechanisms [after 30, after an original sketch by M.
•Scholer].
High
upstream
densitynoparticle
acceleration
In di↵usive
shock acceleration
the particle is —>scatteredaroundthe shock being
much faster than the shock. The requirement isconclude,
the thepreliminary
foundofusing
our numerthe presence of upstream waves and downstream turbulence or waves. InToshockheating
particle solution
is a member
the mainparticle(highdensity—>pressure
dominant@Weaver’s
downstreamdistribution,is trapped
for aradiation
while at the shockand therebyicalthermalised
accelerated
until
leaving
the Inshock.
schemeis andconsistentwithresults.
addition, thedetailed spectra support the validity of Weaver’s approximations—> photons diffuse to upstream —>decelerate
theTheupstream
fluid
regarding
the radiation spectrum.
fact that the resultsforNR shocks are in agreement with previous work supports the
The —>
belief ingradualCosmic Rayvelocity
acceleration byshocksislargefuelledby theMediatedspatial isotropyShock]of Cosmic Rays
as well fromchange[Radiation)
validity of thenumerical scheme.
its approximate power law shape over wide ranges of the spectrum even though the spectrum exhibits several breaks
•

de
th
pr
w
fo
va
ha
sp
of

22

Particle Acceleration
2

ξB
10−3

del

Rcs > RisΓcj
3.3

IS Rad Med.

dels

s

Preliminary
CS Rad Med.

α22.0

• Cosmic-ray production requires high Lorentz factor jets
εγ,pk [keV]
1. Theallowedparameter
j -Liso plane forΓ ~ FIG.
200 forinternal
shocks,Γ ~range
500 onforΓcollimation
shocks1.7
internal shocks (blue-solid lines) and collimation shocks (redΓ forlines)
internal
shock
leads tomediation
larger dissipation
radius3.3 • Highdashedthroughthe radiation
condition, τu <(thick lines) orwith
τu <ourτcr (thinlines). The dissipation radius—> 1inconsistentassumption
condition, Ris < Rcs (purple line), and the jet head condition,

23

Critical Energies
•

•
•
•

High-photon density—> photomeson production limits acceleration
Ep,max ~ 1 PeV for CS, Ep,max ~ 10 PeV for IS
Strong magnetic field —> synchrotron is effective
Eπ,syn ~ 0.2 TeV for CS, Eπ,syn > 100 PeV for IS
High baryon density —> Hadronic collisions is important
Eπ,had ~ 1.5 TeV for CS, Eπ,had ~ 5 PeV for IS
Small dissipation radius:—> adiabatic cooling is effective for IS: Eπ,ad ~ 1 PeVCollimation shock cannot produce
high-energy neutrinos of > 10 TeV

Jet Head

CollimatedJetγγ

Collimation
shock

p

p

EjectaCocooninternal
shock

y
.
,
y

24

Neutrino Emission from IS
4

dL=300 Mpc• 1-100 TeV neutrinos for IS
Preliminaryegood for IceCube detection
IS model Anh• Merger rate:
IS
model
B
R~1500 Gpc-3 yr-1
1
Beaming factor: fb~ 0.045
s
f
where
is the eﬀective—> Aon
event area.
rateSince downg
eﬀ axis- TABLE II. Detection probability of neutrinos by IceCube (IC)suﬀer from the-1 atmospheric background, w
R~4yr- and IceCube-Gen2 (Gen2)
theon
upgoing+horizontal events that have decl.
Detection Probability for a single event◦ northern hemispherein−5. We use the eﬀective area shown in
FIG. 3. The muon neutrino fluences from the internal shockcmodel
p1 (IC)(model A: solid
p1 (Gen2)
p2 (Gen2)
for IceCube. For IceCube-Gen2, we use 1
models for optimisticline) and conservative≃•
IceCubecandetectneutrinoA(model B: dashed0.11
0.40
0.093
line)−3cases.
larger
A
than
that
for
IceCube,althougheﬀ−4= B
6.2×100.026
3.5×10on
the specificconfigurations.
The threshold
with
a
few
years
operation,
the neutrino detection is set to 0.1 TeV f
Detection probability−2for a given interval
is επ,dyn ≃ 5.0tvar,−4 Γj,300 Γrel,4 TeV. For low Γj case,
• and
Gen2
can
detect with The
10 detectio
1 TeV
for IceCube-Gen2.
y modelthe hadronic
coolingP3yr
can(IC)be importantdue to theirvery
P1yr (IC)P1yr (Gen2)
P3yr (Gen2)
ity
of koperation
neutrinos is evendescribed
by the Poi
−4
2
6
year
with
o Astrong Γj 0.38
dependence:0.76
εpπ ≃ 16L−1
t
Γ
Γ
0.88
0.998
k
iso,51 var,−4 j,300 rel,4n modelbution, pk = N exp(−N )/k!, and the "pro
TeV. Sincethe(Gen2)
emission PregionforP1yr the(IC)Lorentz
P10yrfactor(IC) atP1yr
10yr (Gen2)
conservativecase
more than k neutrino detections is 1 −shock case0.23
is high, we can0.10
expect much higher
Bthe internal
0.025
0.67
•

i

neutrino fluence at Eν > 10 TeV.

calculate the detection probability for model

25

Quantities

B (G)L g ,iso (erg s−1)

Eg,isoEE-modEE-optPromptFlare
4Plateau

2.9×103
5.0×104
6.7×103
5.3×102
1.8×102

1.2×1049
3.4×1049
6.1×1051
3.5×1048
3.8×1047

3.8×
1.1×
6.1×
1.0×
1.1×

Choked or Successful?

Figure 1. Neutrino ﬂuences from the EE-mod, EE-opt, prompt emission, ﬂare
FIG. 3. The muon neutrino fluences from the internal shockand plateau for dL=300 Mpc.

E < 300 TeV
models for optimistic (model A: solid line) and conservative••

E > PeVnumber density. This makes EEs more luminous than theothers. The magnetic ﬁelds are so strong that spectral breaks
Jet
due to both the muon and pion cooling supressions are seen in−2
Collimatedis επ,dyn ≃ 5.0tvar,−4ΓΓTeV. For low Γj case,Figure 1. The proton maximum energy is determined by the
Jet j,300 rel,4Collimation
the hadronic cooling can be importantdue to their veryphotomeson production, leading to relatively lower values ofshockν and the−426strong Γj dependence: εpπ ≃ 16L−1
t
Γ
Γ
E
f
<
1
.
For
the
other
three
models,issatisﬁedp,Mμvar,−4j,300
iso,51pg
rel,4
ν so weak
Cocoo
TeV. Since
Lorentz factor at the emission region forlower ﬂuences are obtained. The magnetic ﬁelds areν
Eject theinternalπin these models. The
the internal shock case is high, weshock
can expect much higherthat pion cooling is not important(model B: dashed line) cases.

γ

γ

26

Long-duration
High-energy Photons
from NS mergers
3) Murase, Toomey, Fang, Oikonomou, Kimura et al. 2018, ApJ, 854, 60(Most part of this work was done before the report of GW170817 )27Central Remnant?

Rep. Prog. Phys. 80 (2017) 096901

Review

Baiotti & Rezzolla 17

How
long
is
lifetimeofHMNS?
of the initial mass of the binary. Depending on the initial total mass of the binary M, and on how it relates to the maximum mass of anonrotating neutron star M, the binary can either collapse promptly to a black hole surrounded by a torus (top row), or give rise to an•
Final
remnant
BH(middle
or row),NS?
HMNS (or an SMNS) that ultimately
collapses
to a black hole isand torus
or even lead to an SMNS (first differentially and•

Figure 1. A schematic diagram illustrating the various stages in the evolution of an equal-mass binary system of neutron stars as a function
TOV

subsequently uniformly rotating) neutron star that eventually yields a black hole or a nonrotating neutron star (bottom row). Also indicatedin red are the typical frequencies at which gravitational waves are expected to be emitted (adapted with permission from [59]).

is uncertain,
theofphenomenologi
Gruzinov (2the formula
suchthea long-lived
or ma
injection pulsar
luminosity
beentime
investigatedin byvarious
cont
are
given
the
spin
w
SGRBs,
isotropic
opticaland X2spin-downtimeas
nd mass of Mw = L
0.02
M
.
While
this
is
suﬃcientas(11)
w = ηw⊙Ṁd Vw ,GWsources, andradio bur
The Astrophysical Journal, 805:82 (13pp), 2015 May
20
cf. Talkfast
by Mori-san
illustration, the value is highly uncertain.
4 2 6
4π
B∗2001;
R∗dRP
Zhang
&
Mészáros
1998;
and we
set
η
=
0.3
and
V
=
0.3
c.
In
this
work,
we
w fallback accretion
w
•
•
NS:
spindown
energyBH:≈
L≈
We
have(6pp),
found
that
detecting X-ray and gamma-ray Fan & XuLKisaka
inj
sd
ers,
804:L16
2015
May
1
&
Ioka
3 dt
2006; Metzger etcal.
only consider the acceleration of primary electrons with—>Wind
nebulaformation—>
Remnant
accretion
disk
missionfromahidden
BH
disk
wind
is
diﬃcultfor48
Rowlinsonetal.≃2013;
Totani2011;
a simple power-law injection. For simplicity, the spectral
The internalenergytrapped
1.4
×10High-Energy
Signaturesfrom
Neutron Star
rLate-Time
fiducial
simplybecausetheCompact
mass ac-Remnants
Murase+15
Metzgerof& Double
Piro
2014;
Gaodetint al. 2
index
is setparameters,
to s = 2.2,Kisaka+15
theelectroninjection
Lorentz= L emetion
rate
quicklydeclineswithtime.
However,asdt
Murase
et
al.
2017;
Piro
et a
2016;
factoris fixed
to γe,i that
= 1,a fraction
and the ofenergy
fractions of nant BH and
It is natural
to expect
the dynamicaland dedicated radiative magnet
own
in Figuresradio
emissionmayical simulations
fo
the formulawillof Gruzinov
non-thermal
is assumed
to beϵe Fallback= 0.1.
where
tdyn 2» (2005)
Rej3Vej is the
ecta
fall
backelectrons
onto4,a high-frequency
centralcompact
object,and Wethe alsobenecessary.
Pi emission
Ic is exp
gamma-ray
the weinjection
and≃mthassumethat ϵB may= 0.003of late-prompt
thediskoutflow
kinetic
energydetectable.
For
example,ALMAemission
has sensitivity
tconsider≈ disk2 emission1
inj ≈ tsdluminosity
ll-backmaterial
powerofFirst,
the2
6
we only
considerenergy
2π
B
R
2
∗
∗
time
are
given
by
the
spin-downM∼w V0.01/2 isstoredthemagnetic
and parameters,
weuse a that X-rays are produced
mJy
νas=
100
GHz.etfield,For
ourw(Kisaka
&at
Ioka
2015;
Kisakaal. 2016).
The
GRBs
of
phase, byas ain coronanormal SNe,
56
asthe initial
wind
mass
ofrateM= 0.02M⊙time. Whilesuﬃcientsuch
Nislimcanfall-backbepeakwrittenas (atthis
diskwhere
itself. time
Analogously
to asthePi isisotopes
rotat
eass
radio
flux
atwcanthe∼ is30− 100 asd) isor a spin-down
may assume that late inter−2
an∼illustration,
the
value
is
highlyuncertain.
4at
2 the
6 (where−4
super-Eddington
accretion
flows
p
!
"
magnetic
field
surfaca few µJy (d/40
Mpc)
,
and
then
it
decreases−5/3
Figure1. The schematic
pictureofpulsar-aided SNe. We consider
the left case,
4π
B
R
P
are
negligible,
andinjecti∗ ∗2M
t whereX-rayda pulsar wind
is quasi-sphericaland the windbubbleis embedded
in
theL
L
≈
≈
We
have
found
that
detectingandgamma-ray
−5/3
inj
sd
is
eﬀective),one
may
expectthedisklumṀfollowing,(7)
irrelevant
after
their
lifetime
3d (t) =
inertia.
In
particular,thelFν ∝ t
Ṁ
.
Note
that
the
synchrotronSNejecta.
dc
3teje BH
tejedisk wind is diﬃcult for
Visible photons
leave theSemission from a hidden
48−1
been very≃populartoµergexpla
bular spectrum in this BH disk wind model is approx1.4
× 10rem
sH c
4πGM
m
e
max
our M
fiducialparameters,
simplybecausemass
ac- two-dimensional
down
power
isisthehighenough,somesimulahereisthefall-back
disk
mass,
and
t
the
charLdisk = ηradLηrad the observation
−s/2
−1.1
(
d
eje
Edd =andej
SGRBs,
ately
given
by
F
∝
ν
∝
ν
.
Specific
implicaνtions suggestthat However,
the equatorial windredirected by theσT tesc »cretionrate quickly
with
time.
as can be andteristic massejection declines
time scale.
The
fall-backmass!
16 $ #
anisotropicpressure,
andhoop
stressesleadtobipolar$#•∼ 40
a optical
few ×
10depthG and
P
ns of GW+EM
170817are describedinSection2.3.(Bucciantini
Theejectaabsorbshigh-energy
photons
untiletBal.∗their
µ
η
10
−1
e
rad
−1
shown
in
Figures4,high-frequency
emission
maythat −3
could
GRBs
2007,
te has been
estimated
to be Ṁ0 ∼outﬂows
10−4radio−10
Mexplains
where⊙≃ 1.0 × 10 ergP
si2 Ic3the Thomson optica−ej
2008;
Komissarov& sensitivity
Barkov 2007). If not, we expect
a quasiMuraseetal.
2017).
2013;
15
2s BW
becomes
lower
than
unity
,
which
t
»
K
r
(
R
)
R
For
example,
ALMA
has
t
≈
t
≈
≃
130
T
ej
ej
T
ej
inj
sd
rbeMdetectable.
∼
M
∼
0.03
M
and
t
∼
0.01
−
0.1
s
(e.g.,∗
d⊙ejesphericalexpandingstar
ﬂow embedded in the expanding stellar 2π 2 B 2 R6
2.2. Long-lived
spinning
neutron∗ ∗typicallonger(3 - d ) val
KT M
Fν ∼ 0.01
mJy
at &ν Begelman= 100 after
GHz.
For
our
parameters,material(see
Figure1).
Assuminga SN explosion
with than theosswog
2007;
Rossi
2009;
Kisaka&Ioka—>
observablearound
10
30
days
ej
»
where
ηrad Vej∼and10 − 30 istT possible
51 erg, the SN ejecta expands with2inits
velocity

~
10
sn
The
birth
of ata1. therapidly-rotating,massiveNS
uesPiseem
to4pexpl
Rperis thenecessary
initial rotation
the
flux
peak
time
(atBH∼
30−short100d) SeeisjustFigureSchematic
picturesof our
model
for
GRBs.
Section 2where
for details.
15). radioThe
post-merger
ejectamass
may
belarger
than
ej
model
with
Ṁ
≫
Ṁ
,
dependingradius
.
The
early
PWN
radiusR
also
increases
nonR
ej
w
d
Edd
−2
magnetic
field at inthethesurface,
F
∼
a
few
µJy
(d/40
Mpc)
, toandthenitasisdecreases
eerdynamical
ejectamass,
leadinga largerremnant
thecoalescenceseems
ubiquitous,
hasbybeenplateaus
long-liveν
relativistically,
which
given
(e.g.,ofMetzgeret as
al. 2014)
5000´ ( Vej and
the
disk
outer
radius
and
inner
radi
In particular,
the-1long-liv
ass,
a series
disk-driven
outflowor itsmay sim- inertia.
as Fνwhile
followingṀd . relativistic
Notethatradiation
thenumericalsynchrotronundin∝ ta−5/3
of fullyThus,
allowingforwa2014).
2008;
Ohsugaet al.
2005).
In this3-d

used
for particle
acceleration,
where
one
can
expectsyn-alson-thermal
electronsisassumed
to
be
ϵ
=
0.1.
We
e
28
chrotronemission
consideredpulsar wind
nebulae.
sume
that
ϵB = as0.003
of theinoutflow
kinetic
energyThe2 kinetic energy of the outflow is written as
w V /2 is stored as the magnetic field, and we use a

Emission by Remnant

The Astrophysical Journal Letters, 804:L16 (6pp), 2015 May 1

29 39

5

10

15
t [d]

20

25

10

30

-18

10

5

10

6

10

7

10

8

t [s]

Remnant Disk Emission

Figure 3. X-ray light curves from a BH with a remnant disk

Figure 1.

Thermal bolometric luminosities from a NSNS merger. We show the case where thermal radiation
is mainly powered by radioactive decay of r-process elements and modified by X-ray emission from the disk with
Ldisk = 1040 erg s−1 for demonstration purposes.

Murase et al.

-13

10

-14

10

-15

10

-16

10

-17

10

5

sun

106 sec

10

-14

10

-15

10

-16

10

-17

10

10

-2

10

-3

10

-4

10

-5

10

-6

30 keV

3 keV

10

5

10

6

10

7

10

8

t [s]

-18

10

dissipative outflow (M=0.02Msun, V=0.2c, high K)
dissipative outflow (M=0.02Msun, V=0.2c, modest K)

disk emission,baseline (M=0.02Msun, V=0.3c, modest K)-1
10
disk emission, baseline (M=0.02Msun, V=0.2c, high K)
disk emission, optimistic (M=0.02Msun, V=0.3c, modest K)-13
disk emission, optimistic (M=0.02Msun, V=0.2c, high K)
10

disk emission, optimistic (M=0.02Msun, V=0.2c, high K)

107 sec

Light Curve

0-12Fν [mJy]

10

10radiation luminosity
disk emission, baseline(M=0.02M
sun, V=0.3c, modest K)
radioactive
decay
disk emission, baseline (M=0.02Msun, V=0.2c, high K)
fall-back disk
disk emission, optimistic (M=0.02M, V=0.3c, modest K)

-2 -1

-12

10

E FE [erg cm s ]-2 -1

E FE [erg cm s ]Spectrum10

in the disk emission model, for E = 3 keV (thick curves) andE = 30 keV (thin curves).

3

10

15
t [d]

4

20

10
E [eV] 25

5

10

6

30

Figure 2. X-ray spectra from a BH with a remnant disk

-18

10 Figure 4. High-frequency radio light curves from a BH with
5
6
7
8
10
10
10
10
a remnant disk in the dissipative
outflow model, for ν =
t [s]
10
11

10 Hz (thin curves) and ν = 10 Hz (thick curves).
in the disk
emission model,
at t = 10
curves) and
Figure 3. X-ray light curves from a BH with a remnant disk
Thermalbolometricluminosities
froms (thina NS7•s (thickt = 10curves).
The distanceset to d = 40 Mpc.
in the disk emission model, for E = 3 keV (thick curves) andWeshow
the case
where
thermal isradiation
Note thatthe baseline decay
and optimistic
cases areE = 30 keV (thin curves).
oweredby6 radioactiveof r-process
ele-degenerateor its coronaFigure
may1. Schematic
be reasonablemass
at t = 10 s.
pictures of ourwhen
BH modelthe
for short
GRBs.accretion
See Section 2
6

Opacity for hard X-ray is lower than that for soft X-ray—> emission
Hard from
X raysbecomebright earlier
modified by X-ray
the disk
withrate is not far from that of ultra-luminous X-ray sources.
erg s−1 for demonstration purposes.
• Observable by NuSTAR, but not@
At detected
earlier times, thediskGW170817
radiation luminosity should not0
neutrino–neutrinoannihilation
isthe
notEddington
effective at late
time
where
, the ang
κ the
≈ 0.0510 largely
sion, Athena+ (Nandra et al. 2013), is planned to reachexceedluminosity,
but
outflow
dissipative
outflow(M=0.02M,
V=0.2c,highK)(> 1−10 s). Our model is based on general topologicalsunoutflowcan(M=0.02M
K)
a limiting sensitivity of EFE ∼ 10−17 erg cm−2 s−1 inkineticdissipative
luminosity
be larger.
Amodestsignificant
fractionsun, V=0.2c,

30

8

10

5

10

6

10

7

10

8

10

-2

10

-3

10

5

10

6

10

t [s]

7

10

8

t [s]

Wind Nebulae Emission

X-ray light curves from a long-lived pulsar as aemnant, for E = 3 keV (thick curves) and E =hin curves).

Figure 8.

High-frequency radio light curves from a longlived pulsar as a merger remnant, for ν = 100 GHz (thickcurves) and ν = 10 GHz (thin curves).

Late-Time High-Energy Signatures from Compact Remnants of Double Neutron Star M
Figure 1. The schematic picture of pulsar-aided SNe. We consider the left case,where a pulsar wind is quasi-spherical and the wind bubble is embedded in the
SN ejecta.

X-ray Light Curve

γ-ray Light Curve
6

10
down
power16isG,
highenough, sometwo-dimensional
simula- hig
16
B*=5x10
Pi=10ms(M=0.02M, V=0.2c,sun
B*=5x10G,P=10ms(M=0.02M,
V=0.2c,
high
K)
15
tions
suggestthattheequatorial
wind
can
be
redirectedby theisun (M=0.02M
15 5
B(M=0.02M=1013G,
Pi=3ms, V=0.2c,
mode
*
sun
B*=1013G,
P=3ms,
V=0.2c,modestK)
anisotropicpressure,
andhoop
stressesleadtobipolarsun
10 i
10 G, PB
=10
=1ms
(M=0.02M,
V=0.2c,modethat could
GRBs(Bucciantini
outﬂows*
i , explain
sunK) et al. 2007,
B*=10 G, Pi=1ms(M=0.02M
V=0.2c,modestsun

16
B*=5x10
G, P =10ms (M=0.02Msun, V=0.2c, high K)
15 −2 i −1
B*=1013G, Pi=3ms (M=0.02Msun, V=0.2c, modest K)
B*=10 G, Pi=1ms (M=0.02Msun, V=0.2c,
15 modest K)

51

sn

ej

ej

Fν [mJy]-2 -1

E FE [erg cm s ]-2 -1

E FE [erg cm s ]−14-8
EFE ∼ 10erg cm s , hard X-ray emis-8
10
10
m the embryonic nebula with B∗ = 10 G and2008; Komissarov & Barkov 2007). If not, we expect a quasi4
spherical expanding ﬂow embedded in the expanding stellar
13G
10
B=10
ms is detectable
up
to
z
∼
0.2
−
0.3.
The
dematerial (see Figure 1). Assuming a SN explosion with
-10
-10
10
10

~ 10 erg, the SN ejecta expands with its velocity V and
30 keV3
prospects are quitesensitive to the spin-down
10
radius R . The B=10
early PWN13
radiusG R also increases nonrelativistically, which is given by (e.g., Metzger et al. 2014)ers. We -12find that radio and sub-mm3 obser2
-12
10
keV
10
10
dR
 æç R ö÷
R
7
15
÷÷
çç
,
(7)
=
+
are more promising, whichisconsistentwith
B=10 G
1 GeV 1 B=1015G
÷
dt
6(3 - d ) M çè R ÷øt10
et al. (2016),
who
proposed
synchrotron
nebufor R < R , otherwise R » R is used. Note that we have
-14
-14
10
10
0
used the ejecta density
10
ion as a probe of the connection between fast
100 GeV(3 - d ) M æç R ö÷
-1
rsts and pulsar-driven supernovae, including
÷÷ ,
çç
(8)
r
=
16
10G
÷÷ø
B=10
R
ç
4
R
p
-16
-16
è
10
10
16G
minous supernovae.
High-frequency radioB=10
emis-2where δ ∼ 0–1 is a typical value used in the literature (Kasen &1067
Bildsten 2010; Metzger et al. 2014). The mixture of material
escape around ∼ 10 − 10 s thanks to theallows us to approximate the inner density proﬁle to be
-18
-3
-18
10
10
reasonably smooth and ﬂat (Chevalier 1977; Chevalier &10
cta mass, the5 fast velocity, 6and the expectation
7
8
5
6
7
5
6
Fransson 1992). For demonstration,
10
10
10
10
10
10 107 we adopt d = 1 throughout
10 10810
10
this work (Kasen & Bildsten 2010; Metzger et al. 2014),
ejecta are largely neutral. In thet [s] case where
t [s]
t [s]
and that the radiation pressure is given by (3 ) » (6 7) r V . Here  is the PWN volume and
015FigureG andms,curvesthe sub-mm
emission pulsar Figure
i = 3
7. • P
X-ray
light
from a long-lived
as a 9. Gamma-ray
Figure 8. 16
High-frequency
radio
light
curveslight
fromvelocity
a long-lived
pulsarthe PWN expansionthat canbe different
from
V is curves
VerybrightX-rays&γ-raysunless
B
>10
G
for
P~10
ms
In
general,R
issmaller
than
,
and
both
of
and
R
V
.
R
R
detected
to z ∼for1.5
with curves)sensi- and asa mergerlived
remnant,
Ea merger= 1 GeV(thick curves)andmerger up
remnant,
E by= 3ALMAkeV (thickE =pulsarforasremnant,for ν =are100
numerically determined in this work. Roughly speaking,approximation
for small values of
R becomes
E16= 100 GeV
(thin curves).
keV mJy.
(thin
curves)and ν R=» 10
GHza good(thin
curves).
∼300.01
Notethat the beradio
synchrotron• curves).
Bshould
very
high
(~10
G)
or
lifetimeisshort
P such that ). The ejecta
2  (implying P 1 5 ms 
−0.5
andradiuscanbedeterminedbyvelocityVRm (which cannot be harder than Fν ∝ ν) isto casesavoid
0.8
2 éê ò dt (  t ) +  ùú
Murase et al. 2016, for lower P− ν −0.7 in our(seenon-detection
ë
û
V =
(9)
M
−2
2014). More generally, we conclude
that gamma-ray deailed discussion), andit declines
as F ∝ t .
16−14−2 −1ν
B
=5x10
G,
P
=10ms
(M=0.02Msun, V=0.2c, hig*isitivity EFE ∼ 10
erg cm s , hard X-ray emis15
In this case,wind radiuswhen
is R » ct . the spintection of a pulsar -8remnant
isthe (collimated)possible
w

3-d

w

ej

w

rot

w

ej

ej

w

ej

w

-d

ej

ej

rot

3
ej

2
ej nb

nb

ej

nb

nb

ej

w

ej

w

ej

ej

rot, i

ej

sn

-1 2
sn,51

i

ej

int

dyn

sn

ej

ej

10

w

w

31

Inverse Compton Emission
from Delayed Jets

cf. Talk by Inoue-san•

Photons of EE/Plateau are up-scattered @ forward shock—> TeV γ emission from SGRBs
12Internal Forward
Dissipation Shock

10

-9

-10

10

-11

10

-12

10

-13

10

-14

Detection might be possible
for on-axis observer
Very weak for off-axis observer

10

-15

10

-16

10

-17

γ-raysoff-axis
Observer-2 -1e

X-rays•-8

E FE [erg cm s ]

e

10

10

Jet

•

on-axisObserver10

Figure 10.

Muextended EIC (θ=0 deg)plateau EIC (θ=0 deg)plateau EIC (θ=15 deg)plateau EIC (θ=30 deg)θj=11°
n=0.001

2

10

3

10

4

5

10
t [s]

10

6

10

7

10

8

Light curves of high-energy gamma-rays gen-

32

Summary33

Summary•

Neutrino detection is a smoking-gun for hadronic CRs
Neutrinos associated with SGRBs are detectable with
IceCube-Gen2 if SGRBs accompanies extended
emissionsSub-photospheric neutrinos are detectable with IceCube
if internal shocks are formed in the pre-collimated jets
BNS mergers might emit high-energy photons

The Astrophysical Journal Letters, 848:L4 (6pp), 2017 October 10•

Parameters
EE-modEE-optPromptFlare
Plateau•Quantities
EE-modEE-optPromptFlare
Plateau•Kimura et al.

Table 1
Used Parameters (Top Section) and Resultant Quantities (Bottom Section)Γ

Lg*,iso ( erg s-1)

E*g,iso (erg)rdiss (cm)Eg ,pk (keV)

Energy Band (keV)

30
10
103
30
30

3×10483×1048
1051
1048
1047

1051
1051
1051
3×1050
3×1050

1014
3×1013
3×1013
3×1014
3×1014

1
10
500
0.3
0.1

0.3–100.3–10
10–103
0.3–100.3–10B (G)L g ,iso (erg s−1)

Eg,iso (erg)Ep, M (EeV)

En , m (EeV)

En , p (EeV)

2.9×103
5.0×104
6.7×103
5.3×102
1.8×102

1.2×1049
3.4×1049
6.1×1051
3.5×1048
3.8×1047

3.8×1051
1.1×1052
6.1×1051
1.0×1051
1.1×1051

21
6.0
60
25
13

0.020
3.9×10−4
0.29
0.11
0.33

0.28
5.4×10−3
4.0
1.5
4.6

4

Late-Time High-Energy Signatures from Comp16

B*=5x10
G, Pi=10ms (M=0.02Msun, V=0.2c, high K)
15
B*=1013G, Pi=3ms (M=0.02Msun, V=0.2c, modest K)
B*=10 G, Pi=1ms (M=0.02Msun, V=0.2c, modest K)-2 -1

E FE [erg cm s ]Note that the temperature and radiation energy density
in the collimation jet is independent of both !
Liso and Γj .
In the collimation jet, np ≈ ncj and B ≈ 8πξB aT 4 ,where ξB is the ratio of the magnetic field energy density
to the radiation energy density.
We plot the timescales for the collimation shock in theupper panel of Figure 2, and tabulate the parameters in
Table I. We do not show other relevant timescales, suchas the advection time tadv = Rh /(cΓcj ) and tp,syn because they are much longer. We can see that the BetheHeigler process suppresses the pion production for 0.01
TeV ! εp ! 1 TeV, while the pion production eﬃciency isalmost unity above εp "1 TeV. The maximum energy ofthe protons εp ≃ 3.1 × 102 TeV. However, the pion coolingsare 1.significant
for εfrom
0.1EE-mod,
TeV dueto prompt
the high
den-ﬂare,π " theFigureNeutrino ﬂuences
EE-opt,emission,
Mpc.
plateaudL=300
sityandandtheforstrong
magnetic
field in the collimation jet.

+horizontal muon neutrino events (d > -5 ), the atmosphericmuons are shielded by the Earth. For IceCube-Gen2, we use
10 2 3 times larger effective areas than those of both upgoing-8
10+horizontal and downgoing events for IceCube. The effective
area of downgoing muon neutrino events in IceCube-Gen2 may-10
not be simply scaled, but the simple scaling is sufﬁcient for 10
thedemonstrative purpose of this work. We set the threshold-12
10
energy for neutrino detection to 100GeV for IceCube and
1TeV for IceCube-Gen2.
The probability of detecting k neutrino events, pk,10is-14
described by the Poisson distribution. The detection probabilityof more than k neutrinos is represented as p (&m . k ) =-16
10
1 - åi < k pi . We ﬁnd that for EE-mod (G = 30 ), the probability for upgoing+horizontal events, p (&m . 1), is 0.04 and0.16 with IceCube and IceCube-Gen2, respectively. For EE-opt-18
10(G = 10 ), &m  1.7 and 7.9 with IceCube and IceCube-Gen2,105
respectively. It is possible for IceCube to detect neutrinos from

10

6

10
t [s]

7

10

8