Multi-Wavelength Transients with CTA
Kazumi Kashiyama(UC Berkeley, NASA Einstein fellow)

A, the long-lived, low level emission could easily have been missed, while GRB 111209A was seen and.

Diversity of GRB
LGRB

GRB 121027A

SGRBs
GRB 111209A
TDEs?
LLGRBsGRB 101225A
SGRs

Galac-c sources

Levan+2014

ter space for transients in the -ray sky, showing the duration of the burst, and the approximate
At low luminosity there are numerous Galactic sources that we do not include in further detail; a
ft-gamma repeaters (SGRs) in our own Galaxy are shown, as well as extragalactic transients such
GRBs and SGRBs), and the likely population of low luminosity GRBs (LLGRBs). Two recently dThere are however some diﬀerences between the two
events. First, the X-ray data of XRF 060218 demand asoft thermal emission component with an evolving temperature, which has been interpreted as due to shock
breakout from the progenitor star (Campana et al.
2006). For XRF 100316D, we find that the data
do not demand such a component. We noticed that
Starling et al. (2010) claimed a thermal component withkT ∼ 0.14 keV and an energy that is ∼ 3% of the entire
X-ray emission.
time
iso To check
50the consistency, we3fit theisodependentLL
BAT+XRT joint spectra by a Cutoﬀ HLPowerlaw model (with absorption from both Milk Way andthe host galaxy, wabs*zwabs*cutoﬀpl in Xspec 12) and a
Blackbody+Cutoﬀ Power-law model. We find that bothmodels can give equally acceptable fits to the data. For
example, for the slice 2 (from Ttrig +240 s to Ttrig +734 s),
the former model gives χ2 /dof = 1305/1085 with Ngal =2-N
3
121 cm−2 , 3
20
7.05 ×⇢
10LL
cm
and
=3
4.6
× 10
H,host⇠−210
Gpc
yr& 10 while⇢HL
2
the latter gives χ /dof = 1264/1085 with a larger hostgalaxy absorption NH,host = 1.3 × 1022 cm−2 and a thermal temperature that is consistent with Starling et al.
(2010). Note that the peak of the proposed black body2
component
is near ⇠the1-10
low endof theband"peak,LL
keV⇠XRT
10 energy"peak,HL
and the neutral hydrogen absorption around the thermal
peak is large. This makes great uncertainty on identification of such a thermal emission component fromthe data. Our spectral analysis cannot robustly identify
this component, although the χ2 is slightly improved byadding it. Therefore, we do not claim a thermal emission
in the observed spectrum and only stick to the Cutoﬀ
power-law model to discuss possible theoretical implications. The second diﬀerence between the two events liesin the supernova data. The modeling of SN 2006aj suggests a kinetic energy ∼ 2.5 × 1051 erg (Mazzali et al.

Cutoﬀ Power-law model fit (line) for the time-integrated sfrom Ttrig + 138 s to Ttrig + 736 s. The model paramΓ = 1.32 ± 0.03 and Ep = 19.6+3.3−2.8 keV.

Low-Luminosity GRBs

• Much dimmerE⇠ 10

erg ⇠ 10E

• Nearby (ex. 060218@140Mpc )• More frequent• Quasi-thermal soft spectrum• Associate broad line type Ic SN→ Walf-Rayet star

Zhang+2010

Two Competing Scenarios
Trans-relativistic
shock breakout fromoptically-thick wind

Low-power
relativistic jetVSmagnetar not BH?

Toma+2007
Fan+2010

Waxman+2010
Nakar & Sari 2012

Shock Breakoutü The shock downstream is radiation-dominated.

Prad & Pgas

Stellar envelope
or CMS⇢ü The shock is initially inside optically-thick media.

⌧ ⇡ ⇢T r

1

Radiation-mediated shockr
SN shock

Shock acc is suppressed. e.g., Weaver 1976

vsh

c/
⌧

Shock breakout @ r = rsb
where c/⌧ ⇡ vsh
The downstream photons begin to escape.
ü Shock breakout emission à LL GRB?
ü No longer radiation-mediatedà Shock acc starts à CR, γ, ν?

2008 Ja

ratio
count

0.1 0.5 1 1.50.5 0.02
1 1.5
0.

The temporal evolution is characterized by a fast rise and exponential decay, often observed for a variety of X-ray flare phenomena(Fig. 1). We determine the onset of the X-ray emission to be 9z20
{8 s
before the beginning of the observation, implying an outburst startNNtime (t0) of January 9.5644 UT. The X-ray spectrum is best fitted by a
30"
30"
E
0.5
1
2
5
E
power law (N(E) / E2C, where N and E are the photon number andchannel energy (keV)energy,22
Soderberg etal. respectively) with a photon index of C 5 2.3 6 0.3, and a
c 10
21
{2
, in excess ofhydrogen column density of NH ~6:9z1:8{1:5 |10 cm
Soderberg+2008
theabsorption
within
the
Milky
Way
(see
SupplementaryBlack Bodydyn
1 unabsorbed
2
Information).0.5The inferredpeak flux 5is FX,p <
210
22
21
8
(keV)
6.9 3 10erg cm s channel(0.3–10energykeV). Wealso measure significant
spectral softening during the outburst.
The XRO was in the field of view of the Swift Burst Alert TelescopeBlack Body
6
(BAT; 15–150 keV) beginning 30 min before and continuing throughout
the outburst, but no c-ray counterpart was detected. Thus, the outburst
was not a GRB (see also Supplementary Information). Integrating overthe duration of the outburst, we place a limit on the c-ray fluence of4
iso dyn1028 erg cm22 (3s), a factor of three times higher than an
fc = 8 3 Power−law
extrapolation of the X-ray spectrum to the BAT energy band.
The total energy of the outburst is thus EX < 2 3 1046 erg, at least2
three orders of magnitude lower10 than GRBs. The peak luminosity is
43
1 orders of2 magnitude larger
5 than theerg s21, several
LX,p < 6.1 3 100.5
channelenergy(keV)
Eddington luminosity (the maximum luminosity for a spherically0accreting source) of a solar mass object, outbursts from ultra-luminous
X-rayspectrum
sources andtypeX-ray
I X-ray
bursts. In
the properties
ofFigure 1. Swift/XRTof theoutburst
fitsummary,with power-law(top) an
400
500
600
0
100
200
300
XRO
080109aredistinct
from
those
of
all
known
X-ray
transients.
2 panels) reveals
5 that th(bottom) models. A comparison 0.5of the model 1residuals (lower
Time since tobs = 2008 January 9.5645 UT (s)channelenergy(keV)model provides a better fit to the data.
The birth of a supernova
Figure 1 | Discovery image and X-ray light curve of XRO 080109/
dyn
sb images sh
Simultaneous
observationsof the fit
field
the co-aligned
1. Swift/XRT
spectrum of
the X-ray outburst
withwith
power-law(top) anand ultraviolet (right)
ofFigure
the field obtained
SN 2008D. a, X-ray (left)
Ultraviolet/Optical
Telescope(UVOT)
on
board
Swift
showednothe type(bottom)
Ibc
supernova
on 2008 January 7 UT during Swift observations of3comparison0.5of the model1residuals (lower
revealsth
2 2 models. Aevidence
2 panels)
5 thatfor a contemporaneous
counterpart.
However,
UVOT
2007uy. No source is detected at the position of SN 2008D to a limit of
iso
sb sh
model provides a better fit to the data. channel energy (keV)=1023 counts s21 in the X-ray band and U > 20.3 mag. b, Repeatedobservations just 1.4 h after the outburst revealed11 a brightening
SN 2007uy

SN 2007uy

ratio

0.01
0.1

20.2 30.01
10.022 0.05
3 10.1⇠ E /t0.5 1 1.5

X-ray count rate (s–1)

⇠t

−1
−1
count s−1 keV
counts−1 keV−1
ratio
count
s−1 keV
ratio

ratio

What Shock Breakouts Tell Us⇠ r /cSpectrum: quasi-thermal
E ⇠ 2⇡⇢rc
The temperature and the power low tailultraviolet and X-ray observations of the field from January 9 during which
ultraviolet/optical counterpart. Subsequent ground-based optica
⇢
⇠
1/
r
we serendipitously discovered
XRO
080109andits ultravioletTsb
sh counterpart. observations
a coincident
source.
also Body
uncovereddepends
on the detailsof plasma.
Black
The position of XRO 080109 is right ascension a 5 09 h 09 min 30.70 s,t

UTLight curve

rsb ,

sh

We promptly obtained optical spectroscopy of the counterpart± Is relevantwith the (Gemini North
8-m telescopefor beginning.) thesh 1.74 d after
outburst (Fig. 2). The spectrum is characterized by a smooth continuum with narrow absorption lines of Na I (wavelengths 5,890
0.1

declination d 5 33u 089 19.10 (J2000) (63.50), about 9 kpc from the centre of
NGC 2770. c, X-ray light curve of XRO 080109 in the 0.3–10 keV band. The
nnn%eXkli\%Zfd&eXkli\
data were accumulated in the photon counting mode and were processed using
version 2.8 of the Swift software package, including the most recent calibration

nt s−1 keV−1

11–13

e

& 0.1

dec ∼
8−1−1> 10−4 (vṀ ∼wind /10 cm s ) M⊙ yr , consistent with the mass-loss rate inferred from the t
X–ray component.
Lower panel: the UVOT light curve. Filled circles of diﬀerent colors represent diﬀerent UVOT
red – V (centered at 544 nm); green – B (439 nm), blue – U (345 nm), light blue – UVWnm); magenta – UVM1 (217 nm) and yellow – UVW2 (188 nm). Specific fluxes have been mul
by their FWHM widths (75, 98, 88, 70, 51 and 76 nm, respectively). Data have been rebinincrease the signal to noise ratio. The UV band light curve peaks at about 30 ks due to thebreak-out from the outer stellar surface and the surrounding dense stellar wind, while theband peaks at about 800 ks due to radioactive heating in the SN ejecta.

Shock Breakout Scenario for LL GRB
E ,iso ⇠ 1050 erg t ⇠ 3000 s
sh ⇠ 1rsb ⇠ 9 ⇥ 1013 cm

Campana+2006

Too large for a C-O WR ⇠ 1011 cm⇢(rsb ) ⇠ 10 14 g cm 3Ṁ ⇠ 0.1 M

yr 1

A trans-relativistic shock breakout froman optically-thick envelope formed by a strong wind.

TeV Gamma Rays Counterpart
Kashiyama et al.
proton acc. @ collisionless sh. + s.b.o ph. → pγ int.
-6
-7
-8
-9
-10

8

9

0

1

10

10

-13

0M

pc

-12

pc

-11
M

within the source andduring propagation2-2-1
log(Eγ Φγ [erg cm sec ])!e±attenuatedinjected
CTA (0.5hr)

KK+2013

Including the effect of

2

3

4
5
6
log(Eγ [GeV])

7

8

9

Detectable
even
away
by CTA. torelativistic
shock
Fig. 3.—from
Energy100
fluxes Mpc
of gamma
rays corresponding
We set ϵCR = 0.2
The emission duration is equal to that of X rays; tγ = 2.0×10

Nakar 2015

If the rela s.b.o is driven by a chocked jet, …
Precursor Neutrinos

3

Prompt Neutrinos

ν

γ

ν

Senno+2015

γ

Emerging Jet
Shock Breakout
Stall Radius
Choked Jet

ExtendedMaterial

ExtendedMaterial!

!
CE

Progenitor
Core

CE

ProgenitorCoreodel for jet-driven SNe. Orphan neutrinos are expected since electromagnetic emission
FIG.panel:
4:fluxesfrom chokeddiffuseneutrino flux
from
LLAll-flavor
GRBs
may
be
compatible
s mayübeTheobserved
as hypernovae.
MiddleThediﬀuse
shockneutrino
breakoutmodel
for LLj
Neutrino
emissionneutrinos
from LL GRBs
shown
for the CJdriven bychokedChokedjets(~10-100
produceprecursor
sinceis the
gamma-raywiththejets.
IceCubedata
TeV)
without
contradicting
model (this work) and EJ model [28]. In addition, orphakout later
thantheneutrinos(e.g.,[25]).
Right
panel:
emerging
jet model
forthe limits on classical GRBs.neutrino
emission
fromThechokedjets is included(thick curv
nd gamma rays are produced by the successfuljet,both messengers can be observed
See thetextandfor details.

ü Such events do not contribute significantly to the Fermi b.g..

those without accompanying GRBs, contain a chokedsuch as in the CJ model.

sizes in a large cosmic volume. The first gtion of surveys covering large volumes w
signed to find numerous distant type Ia SN
cosmological use. These observed relatively
fields of view to a great depth, placing mosteffective survey volume aredshift (5).
upernova explosions play-23
An alternative method fimportant roles in many
SLSN−I
veying a large volume of
aspects of astrophysics.
-22
SLSN−II
to use wide-field instrume
hey are sources of heavy eleSLSN−R
SLSN
threshold-21
cover a large sky area wiments, ionizing radiation, andSN IInatively shallow imaging.
nergetic particles; they drive
SN Ia
-20
most of the survey volu
as outflows and shock waves
SN Ib/c
low redshift, one can cond
hat shape star and galaxy for-19
SN IIbefficient untargeted surve
mation; and they leave behindSN
II−P
nearby SNe. Such survey
ompact neutron star and black-18
vided the first well-observ
ole remnants. The study of super-17
amples of SLSNe, suchovae has thus been actively
1999as (6), which turned
ursued for many decades.
-16
be the first example of th
The past decade has seen thetremely 56Ni-rich SLSN-R
iscovery of numerous superlu-15
minous supernova events (SLSNe;(7), and SN 1999bd (8) (Fig. 1). Their study is motivatedwhich
is probably the firs
-14
Gal-Yam12y their likely association with
documented example of the S-13
he deaths of the most massive
II class (9).
-100
0
100
200
300
400
500
600
ars, their potential contribuFurther important deteDays from peakon to the chemical evolution of
resulted from the Texas S
he universe and (at early times) Fig. 1. The luminosity evolution (light curve) of supernovae. Common SN explosions nova Survey (TSS) (10) (teo its reionization, and the possi- reach peak luminosities of ~1043 ergs s−1 (absolute magnitude > −19.5). Super- On 3 March 2005, TSS deütheySomeclasses of luminous
SNe cannot
be explainedby thatclassical
56Ni
ility thatare manifestations
SNe (SLSNe)reach luminosities
are greaterby adecay
factor ofmodel.
~10. The SN 2005ap, a hostless traf physicalexplosionmecha- prototypicalat 18.13(e.g.,mag. Its
redshift w
events of àthe three
SLSN classes—SLSN-I(4)], orSLSN-II[SN windü Needadditionalenergy source
circum-stellar
shock[PTF09cnd(e.g., IIn)pulsarIbc)
isms that differ from those of 2006gy (12, 13, 77)], and SLSN-R [SN 2007bi (7)]—are compared with a normal 0.2832, which indicated aheir more common and less lu- type Ia SN (Nugent template), the type IIn SN 2005cl (56), the average type Ib/c solute magnitude at peak aminous cousins.
light curve from (65), the type IIb SN 2011dh (78), and the prototypical type II-P SN −22.7 mag, marking it as thWith extreme luminosities ex- 1999em (79). All data are in the observed R band (80).
luminous SN detected unt

s radioactively powered (SLSN-R), hydrogen-rich (SLSN-II), and hydrogen-poor (SLSN-I, the mostuminous class). The SLSN-II and SLSN-I classes are more common, whereas the SLSN-R class isetter understood. The physical origins of the extreme luminosity emitted by SLSNe are a focus ofurrent research.

Diversity of Supernova
Absolute magnitude (mag)SHadronic signals fro

where V ≫ V i
Non-Ni-decay-powered Supernovae
suggests that a sig
ejcs

CSM mass is larg
et al. 2013b). De
Obs. tellimportant
us … for deta
profile of the ejecta
by lower-velocityü pre-collapse
4
can be describ
massthatejectionInteraction-poweredfrom(Chevaliermassive 1982a).
stars
comparable to Me
can beü earlythat
shock
acc.descri(see Truelove & M;80./(8<#)25$8/+4(#>$'708.4$/#)25$8/+4
to push the basic=23)(875+-$8$0#)25$8/+4(##@(AA(78(<#128)'?
and many other co62().7)5*$8.&(3#!"#-./0#
8$0.8$&'$0#!"#-./0#79#1.5+3(8Figure 1. The schematic picture of the interaction-powered SN scenario.
cuss non-thermal
treatment still prov
collision between the SN ejecta with a CSM shell and found that CR
non-thermal signa
Obs.
tell
us
…accompanying
protons may be accelerated, and furthermore thattheprotonsmay$/4"!#$%$&'(#
anexperience strong pionic losses via inelastic pp collisions, producing
Hereafter, we a)*)*+&,#
ü
Physicalpara.
gamma-rays and neutrinos. Thus, interaction-powered
SNe
may
be
12113$#
densityof
profileand12113$#
interesting CR accelerators and high-energy/multi-messenger emit-newborn
pulsars
expectthat this is
-./0#
-./0#
Fast-spinning-newborn-pulsar-powered!"#!"#of the possibility of“Howters. In this work, we continue our study
the magnetars
born?”
details are
uncertainon-thermal emission from the shock interaction of a SN embedded
mechanism. Then,in a dense CSM. In particular, we focus on the secondary electrons! "ü Particle acc. inRand positrons expected from the same pp collisions that give rise toyoung PWNe
ϱcs = DR0−2
R0
neutrinos and gamma rays. Importantly, we show that these secondaries can emit detectable synchrotron radiation at high-frequency
where R should bradio wavelengths including mm/submm and FIR bands.
1Constraints from SN light curvee.g., in the case of the pulsar-powered scenario“Width”“Height” Multi-Messenger Tests for Fast-Spinning
Newborn
NSs in SE SNe
Kashiyama et al.

-20

MR [mag]

MR [mag]-21 -20
-19

-19
-18

-18
-17

-17
-16

-16
-15

-15
-20
-40

P0 = 1 ms, Bp = 1×1013 G, Mej = 10 Msun
P0 = 1.5 ms, Bp = 3.5×1013 G, Mej = 7.5 Msun
P0 = 2 ms, Bp = 1×1014 G, Mej = 5 Msun
P0 = 2 ms, Bp = 3×1014 G, Mej = 5 Msun

-10
-20

0
10
20
30
20
40
60
80
100
Days since Maximum
Days since Maximum
Fig. 3.— Comparison of observed SN Ibc and BL-SN Ibc (greyFig.
2.—light curvesof pulsar-driven
SN andmodelswith mslines)and SN
theoreticallight curves.
The red solid
dotted-dashedrotationanddiﬀerent
magnetic-fieldstrength.
The
upper
and
lines show wind powered cases, whereas the blue dashed lineslowershow
0
grey56lines
show theobserved SL-SN-Ic
PTF 09cnd and BL-SNe Ic
Ni-powered
cases.
1998bw, respectively.56 Ni
ej 20

For a given P ,SNasbecomes
if
estimated
Lsn ≈ L Nithe× (tbrightest
Ni /tesc ) , or
56

t

⇠t

56

"!
!
Actually, e.g., Maeda et al. (2007)
proposeda"
newborn
−1
56
56 Ni
M
M
ej
Ni
42
−1
sd
dif
magnetarasarelevantenergysourceofatype
Ibc SNLsn ∼ 4 × 10 erg s
0.1 M⊙ to5showM⊙ in what
2005bf. In this regard, it is important

In this work, we numerically calculate light curves of0
SNe driven
by fast-spinning strongly-magnetized
KK +15newborn NSs embedded in SE progenitors. Details of themodel description are given in Appendix. We assume
that the energy injection is caused by spherical winds
rather 0.5
than jets, and both 56 Ni decay and magnetized
wind are taken into account as energy sources. The
thermalization of the non-thermal emission is approximately taken into account, and the optical SN emission and early non-thermal nebular emission are ob1
tained consistently.
The eﬀect of GW spin-down is incorporated in a simple parametric form. The present
14
Esn = 1×10
erg, Mej = 2 Msun
ms, Bp = 5×10but
G
model is basedon 51Murase
et, Pal.
with sev0 = 10 (2015)51
13
Esn = 1×10 erg, Mej = 2 Msun, P0 = 10 ms, Bp = 5×10 G 56
eral refinements,
e.g.,
including
the
eﬀect
of
Ni decay.
51
14
Esn = 1×10 erg, Mej = 5 Msun, P0 = 10 ms, Bp = 5×10 G
The simple
model allows us to explore a wide parame1.5
ter range;
1 − 30 ms,-20 the initial-10 spin 0of Pi = 10
20 poloidal30
13−15magnetic field of Bdip =10G,SNejectamassofDays since Maximum56
Mej = 1−10 M⊙ , Ni mass of M56 Ni = 0.05−1.0, SN exFig. 4.— SN light curve of pulsar-driven
SN models around the51−52
plosionenergyof Esnwith= 10erg, andbody
opacpeak. The
grey lines
arrows indicate
the greyobserved
range
of
−1
−2
SNeKIbc
and
BL-SN
Ibc.
ity
0.05
− 0.2
g
cm
.
Note
that
K
∼
0.1
and
T =
T
For a given P0 ,0.2 g−1 cm−2 corresponds to electron scattering for singly
SN helium,
becomesslowerand can be
ionized and fully ionizedrespectively,
smaller for e.g.,
partially
C-1 ×or
Mej = 2aMsun, MNi = 0.05ionized
Msun, Εsn =1051O-dominated
erg
for
a
smaller
B15
p
23
ejecta.10
1.0=
In Fig.
2, we 15show some light curve 22exSN Ibc
R,

nitude-22 -2152

Esn = 3×10 erg, Mej = 8 Msun, MNi = 0.8 Msun
51
Esn = 6×10 erg, Mej = 4 Msun, MNi = 0.4 Msun
51
Esn = 2×10 erg, Mej = 2 Msun, MNi = 0.2 Msun
51
Esn = 1.5×10 erg, Mej = 2 Msun, MNi = 0.05 Msun
51
14
Esn = 1×10 erg, Mej = 5 Msun, P0 = 2 ms, Bp = 3×10 G
51
14
Esn = 1×10 erg, Mej = 2 Msun, P0 = 20 ms, Bp = 7.5×10 G

052 erg-23 -223

MR-MR,max [mag]4Multi-Messenger Tests for Fast-Spinning N

KK +151015

Mej = 5 Msun, MNi = 0.1 Msun, Esn = 3 × 1051 erg
23=

BL-SN Ibc
1.0

SN Ibc22
21- Peak R-Band AB Magnitude

Constraints from SN light curve

declineter thelate-tim

EK = 10 52

Bp [G]

MR

erg

Beca
SN mo
20
disting1014
19
els. Toapproa18
nario fSL-SN Ic
17
and SLNi poweredtion to16
1013
sar eng1
10
in x raP0 [ms]Murase13
P0 ⇠ ms, Bp & 10 G ! SL-SNe Ic Murase
6.— Same scenario
as Fig. 5, but with Mej = 5 M14
UnifiedFig.
pulsar-powered⊙ , M56 Ni =nism o
P
⇠
ms,
B
&
5
⇥
10
G
!
BL-SNeIbc
51
0p
0.1 M⊙ , and Esn =3 × 10 erg. Parameter regions broadly confor stripped-envelopeSNelikely osistent with the observed SN Ibc,
BL-SN
Ibc,
and
SL-SN
14 Ic areP0 ⇠ 10 ms, Bp & 5 ⇥ 10 G ! SNe Ibctrons, w
indicated.
5
,1

.3

M R, 15 = 0

log(E FE [GeV cm

-2

-8
-9

Gamma-ray counterparts
-10
-11
-12
-13

1. detect supernovae with high cadence SN surveys (e.g., PTF, Pan-STARRS, Hyper-Supreme Cam)
-5
-4
-3
-2
-1
0
1
2
3
4
5
2. theoretically constraint on the model parameters, select good candidates log(Efor gamma-ray
followup[GeV])

Interaction-powered (SNe IIn, SLSN-II)

!$ !"

-7

Pulsar-powered (SNe Ibc, SLSN-I)!)

Model A
Model B (RS)Model B (FS)

Murase+14

Fig. 6.— The same as Figure 5, but at diﬀerent observation
times.

,-./0 1 0 2345 67 8 9:log(Eγ2 Fγ [GeV cm-2 s-1])

-6

4

-8
-9

Murase,KK +15!(!"#$%&'&()*

!'
!&

14;7<

!"#$%&'&+)*

=>?

!"%

-10

!""

-11

!">?#

!"#$%&'&,)*

!"$

-12

!"#

-13
-1

0

1

2
3
log(Eγ [GeV])

4

5

6

!*

!+

!#

!$

!"
%
"
,-./0 23459:

$

#

+

*

Fig. 7.— High-energy photon spectra of the early PWN embeddedin the SN ejecta for Pi = 2 ms at t = 107.5 s ≃ 316 d. Diﬀerent
FIG. 2. Energy fluxes of pionic gamma rays, corresponding
magnetic
fieldtostrengths
are considered.
Detections with CTA are
3. Fig.Follow-up
observationof inside
nearby10 Mpc)~ months~ yr after
the explosion
to
1. Gamma-ray
attenuation
and(~outside
the SNe possiblefor Bdip = 1013 G.

source is included; the double-dotted curve shows the intrinsic
spectrum without attenuation in Model A. Left dot-dashedcurves show Fermi /LAT diﬀerential sensitivities at t = 106.5 s(∼ tr in Model A) and t = 107.5 s (∼ tr in Model B). The
100 hr diﬀerential sensitivity of CTA is also overlaid (right).

Fast Radio Bursts⇠ 1.4 GHz

⇠ ms
⇠ Jy

⇠ 103 4 day 1 sky 1⇡ ⌫ 2 arrival time pulse sweep

Lorimer+07

radio dispersion in cold ionized medium

DM ⇠ 500-1100 cm 3 pc

DMmw

Figure 2: Frequency evolution and integrated pulse shape of the radio burst. The survey data,collected on 2001 August 24, are shown here as a two-dimensional ‘waterfall plot’ of intensity
as a function of radio frequency versus time. The dispersion is clearly seen as a quadratic sweep
across the frequency band, with broadening towards lower frequencies. From a measurement of
the pulse delay across the receiver band using standard pulsar timing techniques, we determine
the DM to be 375±1 cm−3 pc. The two white lines separated by 15 ms that bound the pulse show
the expected behavior for the cold-plasma dispersion law assuming a DM of 375 cm−3 pc. The
horizontal line at ∼ 1.34 GHz is an artifact in the data caused by a malfunctioning frequency
channel. This plot is for one of the offset beams in which the digitizers were not saturated.
By splitting the data into four frequency sub-bands we have measured both the half-power
pulse width and flux density spectrum over the observing bandwidth. Accounting for pulse
broadening due to known instrumental effects, we determine a frequency scaling relationshipfor the observed width W = 4.6 ms (f /1.4 GHz)−4.8±0.4 , where f is the observing frequency.

ü Cosmological distance (z ~ 1) !? à very high effective temperature à coherent emission
ü Can be an unique probe of cosmology and new physics??

Models• Galactic• Rotating radio transients
Loeb+13• Flaring stars• Extragalactic• Giant pulse from young pulsars
Popov&Postnov 08, Thornton+13, Lyubarsky+14• Magnetar giant flares
Calgate+71,75, Egorov&Postnov 08• Supernovae into nearby stars
Falcke&Rezzolla 13, Zhang 13• Core-collapses of hypermassive NSs
Hansen&Lyutikov+01, Totani 13• Binary NS mergers
KK, Ioka, Meszaros 13• Binary WD mergers
Rees 77, Blandford 77, Kavic+08, Keane+12• Evaporations of BHs
Cai+12• Superconducting cosmic stringsand more…
No counterpart detection so far.

re,
ber
s in
ith

Recent advances

LETTERRESEARCHü at lower
frequencies
(GBT)

2.1
1.0

U

800

Fluence (K ms)
Q

I

10

700
0.0

5
0
5
0
−5
5

0.0

0

Masui+15
−5

–0.8
750
800
850
900
1.2
1.8
Frequency2.4(MHz)Time (s)Figure 2 | FRB 110523 spectra in total intensity and polarization.
Plottedtemperatureis the pulse fluence(time-integrated
total
intensity
Figure 1 | Brightness
spectraversus timeflux)
forfor
FRB
110523.
(StokesI),
andlinearpolarization(Stokes
Q
and
U).
Solid
curvesThe diagonal black curve shows the pulse of radio brightness sweeping are modelfits. In addition to noise, scatter in the measurement around the models
over time. The arrival
time is differentially delayed (dispersed) by plasma
is due to the scintillation visible in Fig. 1. The decline of intensity withalong the line of sight.
A pairof curvesthebeam
FRBacross the sky
frequencyis primarily
duein
towhite,motionbracketingof the telescope
pulse, show that thefunctionand delay
is not intrinsic
to matches
the source.the one expected from cold
700

0.6

ü large scattering and rotation measure

Antenna temperature (K)15ü liner and circular polarization
I

0.8

P+

0.6

P×

0.4

V

0.2
0.0
−0.2

\ (deg)900

Frequency (MHz)

olity
Hz
ded
val
irst
ent

Table 1.
Our detection in a distinct band and with independent instrumentation compared to the 1.4-GHz detections at the Parkes and Arecibo
observatories greatly strengthens the argument that FRBs are astrophysical phenomena. In addition, as described in the Methods, theclose fit to astronomical expectations of FRB 110523 for dispersion
spectral index, Faraday rotation spectral index, and scattering spectral
index all further support an astronomical origin.

Antenna temperature (K)

ure
val
23.
ues
tic
tic
0.5.
his
m,
or
ma
ula
n of
nd
ng
ng
be

0
−20

Masui+15

−40
−10

−5

0

5
Time (ms)

10

15

20

Figure 3 | Polarized pulse profiles averaged over spectral frequency.
Plotted is total intensity (I), linear polarization (P+ and P×), and circularpolarization (V, which may be instrumental). Before taking thenoise-weighted mean over frequency, the data are scaled to 800 MHz
using the best-fit spectral index and the linear polarization is rotated tocompensate for the best-fit Faraday rotation. The linear polarization basis
coordinates are aligned with (+), and rotated with respect to (×), themean polarization over time. The bottom panel shows the polarization
angle (where measurable) in these coordinates. The error bars show thestandard deviation of 10,000 simulated measurements with independent
noise realizations.Kulkani+15, Masu+15

à sufficient plasma density and magnetic field surrounding the source?
à star forming region?
à Magnetar giant flare is the most promising?

plasma. The grey horizontal bars show where data has been omitted owing
to resonances within
GBT
receiver.
Theburst
insetdata,shows
fluctuations
inBythe
fittinga modelto thewe foundthe detection
signifibrightness causedcanceby scintillation.
to be over 40σ, with fluence 3.79(15) K ms at our centre spectral
frequency of 800 MHz. The burst has a steep spectral index of −7.8(4),oad, Vancouver, British Columbia V6T 1Z1, Canada. 2Canadian Institute for Advanced Research,which we attribute to telescope motion. The peak antenna temperature
Center for Cosmology, Carnegie Mellon University, Department of Physics, 5000 Forbes Avenue,at 800 MHz is 1.16(5) K. We do not know the location of the sourceof Chemistry and Physics, University of KwaZulu-Natal, Durban 4001, South Africa. 5National Institutewithin the GBT beam profile, but if the source location were at beamysics, University of Wisconsin, Madison, Wisconsin 53706-1390, USA. 7Academia
Sinica Institute ofcentre where the antenna gain is 2 K Jy−1 the measured antenna tem8d, Taipei 10617, Taiwan. National Astronomical Observatories, Chinese Academy of Science,
perature would translate to 0.6 Jy. Off-centre the antenna has lower gain
Beijing 100871, China.410Astrophysics and Cosmology Research Unit, School of Mathematics, Statistics,
so, as University
in previously
reportedis aToronto,
lower limit
toM5S
the FRB flux.
ent of Astronomy and Astrophysics,of Toronto,
50 Stbursts,
Georgethis
Street,Ontariointrinsic
at half-maximum(FWHM)
13d, Gushan District, Kaohsiung The
City 804,
Taiwan.full-width
Departmentof Physics and Astronomy,
Westduration
Virginia of the(withscatteringremoved)1.74(17)Institute,ms, similar
to the widths ofophysics, 60 St George Street, burst
Toronto,Ontario
M5S 3H8,Canada. 15isPerimeter
31 Caroline

detect so we conclude that the linear polarization is of astronomical
Lyubarsky+14
origin rather than due to leakage.
The rotation measure and dispersion measure we detected imply anelectron-weighted average line-of-sight component of the magnetic
field of 0.38 µ G, compared to typical large-scale fields of ∼10 µ G inspiral galaxies26. This field strength is a lower limit for the magnetizedregion owing to cancellations along the line of sight. Also, the magnetized region may only weakly overlap the dispersing region and so

In a magnetar model,TeV gamma rays are detectable up to ~ 100 Mpc with CTA.

If ~10 per day per sky up to ~ Gpc, ~3 per 5deg x 5deg per yr up to ~ 100 Mpc…

They repeat!

More hopes for just waiting!?

Summary and Discussion• Three approaches to multi-wavelength transients
1. ”Need to harry!” (20 sec < tγ < day)• e.g., low-luminosity GRBs
ü TeV gamma ray is the smoking gun signal of the shock breakout model.
ü also may give an important hint on the origin of IceCube neutrino.

2. ”Need not too much harry” (day < tγ < year)• e.g., non-Ni-powered SNeü Follow up obs. from months to year after the explosion is enough.
ü Gamma rays can test supernova models, probe newborn pulsars, etc

3. ”Just wait” (tγ < 20 sec)
• e.g., FRBsü In a magnetar model, a few times per CTA FoV per yrü Some of them do repeat! à targeted observation is also feasible