We study the evolution of the energy spectrum of cosmicray electrons accelerated at spherically expanding shocks with low Mach numbers and the ensuing spectral signatures imprinted in radio synchrotron emission. Timedependent simulations of diffusive shock acceleration (DSA) of electrons in the testparticle limit have been performed for spherical shocks with parameters relevant for typical shocks in the intracluster medium. The electron and radiation spectra at the shock location can be described properly by the testparticle DSA predictions with instantaneous shock parameters. However, the volume integrated spectra of both electrons and radiation deviate significantly from the testparticle powerlaws, because the shock compression ratio and the flux of injected electrons at the shock gradually decrease as the shock slows down in time. So one needs to be cautious about interpreting observed radio spectra of evolving shocks based on simple DSA models in the testparticle regime.
1. INTRODUCTION
Xray emitting baryonic gas in galaxy clusters is thought to be heated via collisioness shocks that are induced during the formation of the large scale structure in the Universe (e.g.,
Ryu et al. 2003
;
Kang et al. 2007
;
Vazza et al. 2009
;
Skillman et al. 2011
). The gas infalling toward nonlinear structures go through strong accretion shocks first, and later may encounter much weaker shocks produced by mergers and chaotic flow motions inside the hot intracluster medium (ICM) (
Ryu et al. 2003
). In fact, about a dozen of these weak ICM shocks have been detected through Xray observations as bow shocks associated with merging activities (e.g., the socalled bullet cluster 1E 065756,
Markevitch et al. 2002
;
Akamatsu & Kawahara 2013
;
Ogrean & Brüggen 2013
).
In addition, quiet a few ICM shocks have been identified as “radio relic shocks” in the outskirts of galaxy clusters through radio synchrotron radiation from cosmicray (CR) electrons that are thought to be accelerated via diffusive shock acceleration (DSA) (
van Weeren et al. 2010
,
2011
;
Nuza et al. 2012
). Radio relics are interpreted as synchrotron emitting structures that contain relativistic electrons with Lorentz factor
γ_{e}
~ 10
^{4}
gyrating in mean magnetic fields
B
~ 5−7
μ
G in the ICM (
Kang et al. 2012
;
Feretti et al. 2012
;
Pinzke et al. 2013
;
Brunetti & Jones 2014
). The electrons are expected to be accelerated at weak shocks with the sonic Mach number,
M_{s}
~ 2 − 4, inferred from the spectral index
α
_{inj}
＞ 0.6 of synchrotron emission immediately behind the shock. Note that the spectral index of the volumeintegrated synchrotron radiation is normally
A_{ν}
=
α
_{inj}
+ 0.5, since electrons cool via synchrotron and inverseCompton (iC) losses behind the shock (e.g.,
Kang 2011
).
It is well accepted that the DSA theory can explain how nonthermal particles are produced through their interactions with MHD waves in the converging flows across astrophysical shocks (
Bell 1978
;
Drury 1983
;
Malkov & Drury 2001
). In the
test particle limit
where the CR pressure is dynamically insignificant, it predicts that the CR energy spectrum at the shock position has a powerlaw energy spectrum,
N
(
E
) ∝ E
^{−s}
, where
s
= (
σ
+ 2)/(
σ
− 1) and
σ
=
ρ
_{2}
/
ρ
_{1}
is the shock compression ratio. (Hereafter, we use the subscripts ‘1’ and ‘2’ to denote the conditions upstream and downstream of shock, respectively.) Hence, the DSA model can provide a simple and natural explanation for the spatially resolved synchrotron radiation spectrum at the shock location,
j_{ν}
(
x_{s}
) ∝
ν
^{−αinj}
, where
α
_{inj}
= (
s
− 1)/2.
On the other hand, preacceleration of thermal electrons to suprathermal energies and the subsequent injection into the DSA process still remains very uncertain, especially at low Mach number shocks in high beta (
β
=
P_{g}
/
P_{B}
) plasmas (e.g.,
Kang et al. 2014
). Recently, there have been some serious efforts to understand complex plasma instabilities and waveparticle interactions operating at nonrelativistic shocks with the parameters relevant for weak cluster shocks, through ParticleinCell (PIC) and hybrid plasma simulations (e.g.,
Riquelme & Spitkovsky 2011
;
Caprioli & Spitkovsky 2014a
,
b
;
Guo et al. 2014
). They have demonstrated that selfexcitation of MHD/plasma waves via various instabilities, such as Bell’s resonant and nonresonant instabilities and firehose instability, are crucial in the injection process. Moreover, it has been shown that protons can be injected efficiently at quasiparallel shocks, while the electron injection occurs preferentially at quasiperpendicular shocks. Here, we adopt a
phenomenological
injection model in which suprathermal electrons with momentum
p
>
p
_{inj}
are allowed to cross the shock transition and participate in the Fermi 1storder acceleration, where the injection momentum,
p
_{inj}
, is several times the peak momentum of the postshock thermal protons (
Kang et al. 2002
). This subject will be discussed in detail in Section 2.2.
The morphology of observed radio relics varies rather widely from diffuse patches to welldefined, thin elongated structures. For example, the socalled Sausage relic in CIZA J2242.8+5301 is located at a distance of ~ 1.5 Mpc from the cluster center and has an arclike shape of ~ 55 kpc in width and ~ 1 − 2 Mpc in length (
van Weeren et al. 2010
). So it would be natural to assume, at least for cases similar to the Sausage relic, that radio relics consist of some portions of a spherical shell of radiating electrons projected onto the sky plane. The physical width of this shell is mainly determined by the cooling length of electrons, ∆
l
~
u
_{2}
t
_{rad}
(
γ
_{e}
), where
u
_{2}
is the postshock flow speed and
t
_{rad}
is the electron cooling time. But the projection of a partial shell could involve some geometrical complexities such as the extent of the partial shell and the viewing angle (e.g.,
van Weeren et al. 2010
,
2011
;
Kang et al. 2012
).
For several observed radio relics the
projected
spatial profiles of the radio flux,
S_{ν}
, and the spectral index,
α_{ν}
, have been explained by the electron energy spectrum
N_{e}
(
r
,
γ_{e}
) cooling radiatively behind onedimensional planar shocks (e.g.,
van Weeren et al. 2010
,
2011
;
Kang et al. 2012
). This may be justified, since the width (~ 50 − 100 kpc) of those relics is much smaller than their length (~ 1 − 2 Mpc). Note that a simple projection of a spherical shell on the sky plane was attempted in these previous papers, although the DSA simulation results for a planar shock was adopted for
N_{e}
(
r
,
γ_{e}
).
In this work we consider electron acceleration at spherical shocks with
M_{s}
~ 2.5−4.5, which expand into a hot uniform ICM. A selfsimilar solution for SedovTaylor blast wave is adopted as the initial states of the
postshock
gasdynamic quantities (i.e.,
ρ
,
u
,
P
) at the beginning of the simulations. Later on the shock structure deviates from the SedovTaylor similarity solution, because the sonic Mach numbers of our model shocks are not large enough. Nevertheless, the shock slows down approximately as
u_{s}
(
t
) ∝
t
^{−3/5}
. This model may represent a spherically decelerating shock, but it is not meant to mimic realistic cluster shocks. Note that the flow speed,
u
(
r
,
t
), is the only dynamical information that is fed into the diffusion convection equation for the evolution of the electron distribution function,
f_{e}
(
r
,
p
,
t
). As the shock slows down and weakens, the shock compression ratio and the injection flux of CR electrons should decrease in time. Consequently, the amplitude of electron spectrum at the injection momentum decreases and the electron energy spectrum steepens gradually.
Considering that the synchrotron/iC cooling time scale (
t
_{rad}
~ 10
^{8}
yr) for the electrons of our interest (
γ_{e}
~ 10
^{4}
) is shorter than the dynamical time scale of typical cluster shocks (
t
_{dyn}
~ 10
^{9}
yr), the electron energy spectrum at the shock location is expected to follow the testparticle DSA predictions with the
instantaneous
shock parameters. But the spatial profile of the electron distribution function,
f_{e}
(
r
,
p
), and the volumeintegrated distribution function,
F_{e}
(
p
) =
4
πf_{e}
(
r
,
p
)
r
^{2}
dr
, would be affected by the timedependent evolution of the spherically decelerating shock. Our primary objective here is to study any signatures of such timedependence through DSA simulations of a heuristic example (i.e., a spherical blast wave).
In the next section we describe the numerical method and some basic physics of the DSA theory. The simulation results of a planar shock and a spherical shock will be compared, and spherical shocks with different magnetic field models will be discussed in Section 3. A brief summary will be given in Section 4.
2. DSA NUMERICAL SIMULATIONS
 2.1. 1D Spherical CRASH Code
We consider DSA of CR electrons at nonrelativistic, gasdynamical shocks in onedimensional (1D) spherical geometry. So we solve the following timedependent diffusionconvection equation for the pitchangleaveraged phase space distribution function for CR electrons,
g_{e}
(
r
,
p
,
t
) =
f_{e}
(
r
,
p
,
t
)
p
^{4}
:
where
u
(
r
,
t
) is the flow velocity,
y
= ln(
p
/
m_{e}c
),
m_{e}
is the electron mass,
c
is the speed of light, and
D
(
r
,
p
) is the spatial diffusion coefficient (
Skilling 1975
). For
D
(
r
,
p
), we adopt a Bohmlike diffusion coefficient with a weaker nonrelativistic momentum dependence
The cooling term
b
(
p
) = −
dp
/
dt
accounts for electron synchrotron and iC losses:
in cgs units, where
e
is the electron charge. The ‘effective’ magnetic field strength
takes account for radiative losses due to both synchrotron and iC processes, where
B
_{rad}
= 3.24
μ
G(1+
z
)
^{2}
corresponds to the cosmic background radiation (CBR) at redshift
z
(
Schlickeiser 2002
). In this study, we set
z
= 0.2 as a reference epoch and so
B
_{rad}
= 4.7
μ
G. Then the cooling time scale for electrons is given as
For typical ICM conditions relevant for radio relics (i.e.,
n_{H}
~ 10
^{−4}
cm
^{−3}
and
B
~ 1
μ
G), the Coulomb loss dominates over the other losses for relativistic electrons with
γ_{e}
100, while the synchrotron and iC losses are the main cooling processes for
γ_{e}
100 (
Sarazin 1999
). Nonthermal bremsstrahlung loss is much less important and can be ignored for relativistic electrons, since its cooling time scale is longer than the Hubble time (
Petrosian 2001
).
Since we focus on the synchrotron emission from CR electrons here, we do not need to consider the acceleration of CR protons. Moreover, the CR proton pressure is expected to be dynamically insignificant at weak shocks with
M_{s}
5, so we can just follow the hydrodynamic evolution of the shock without following DSA of CR protons. This significantly alleviates the requirements for computational resources, since much wider ranges of particle energy and diffusion length/time scales should be included, if both proton and electron populations were to be calculated. For instance, the maximum energy of electrons are limited to ~ 100 TeV for the problem considered here, while that of protons can go up much higher energies by several orders of magnitude.
At weak shocks in the testparticle limit, the CR feedback becomes negligible and the background flow,
u
(
r
,
t
), is governed by the usual gasdynamic conservation equations in 1D spherical coordinates (
Kang & Jones 2006
):
where
P
is the gas pressure,
e
_{g}
=
P
/[
ρ
(
γ
_{g}
− 1)] +
u
^{2}
/2 is the total energy of the gas per unit mass and the rest of variables have their usual meanings. The gas adiabatic index is assumed to be
γ
_{g}
= 5/3, since the background gas flow is nonrelativistic.
In order to optimize the shock tracking scheme, a comoving frame that expands with the instantaneous shock speed is adopted. In such a frame, a spherically expanding shock can be made to be stationary by adopting comoving variables which factor out a uniform expansion. The details of the CRASH (CosmicRay Amr SHock) code in a comoving spherical grid can be found in
Kang & Jones (2006)
.
 2.2. Electron Preheating and Injection
Injection of suprathermal particles into the Fermi process remains one of the outstanding problems in the DSA theory. In the socalled thermal leakage injection model, suprathermal particles with
p
p
_{inj}
~ 3
p
_{th,p}
are expected to recross the shock upstream and participate in the DSA process, because the shock thickness is of the order of the gyroradius of postshock thermal protons (e.g.,
Kang et al. 2002
). Here,
is the most probable momentum of thermal protons with postshock temperature
T
_{2}
and
k_{B}
is the Boltzmann constant. On the other hand, recent hybrid and PIC simulations of collisioness shocks have revealed a somewhat different picture, in which incoming protons and electrons are specularly reflected at the shock surface and scattered by plasma waves excited in the foreshock region, leading to suprathermal populations (e.g.,
Guo & Giacalone 2013
;
Guo et al. 2014
;
Caprioli et al. 2015
). The trajectories of these suprathermal particles are nondiffusive and confined to the region mostly upstream of the shock transition.
If one assumes that downstream electrons and protons have Maxwellian distributions defined with the same kinetic temperature,
T_{e}
≈
T_{p}
, then thermal electrons have much higher speeds but much smaller momenta compared to thermal protons (i.e.,
p
_{th,e}
= (
m_{e}
/
m_{p}
)
^{1/2}
p
_{th,p}
). So electrons must be preaccelerated from their thermal momenta to
p
_{inj}
~ 3
p
_{th,p}
~ 130
p
_{th,e}
, before they can begin to take part in the DSA process. In fact, preheating of electrons above the thermal distribution has been widely observed in space and laboratory plasmas (e.g.,
Vasyliunas 1968
). Such suprathermal nonMaxwellian distributions can be described by the combination of a Maxwellianlike core and a powerlaw tail, which is known as the
κ
distribution:
where
is the thermal peak momentum, the mass of the particle is
m
=
m_{e}
for electrons and
m
=
m_{p}
for protons, and Γ(
x
) is the Gamma function (e.g.,
Pierrard & Lazar 2010
). In the limit of large
κ
, it asymptotes to the Maxwellian distribution.
The powerlaw index of the
κ
distributions for electrons and protons,
κ_{e}
and
κ_{p}
, respectively, should depend on plasma and shock parameters, i.e.,
β
,
M_{s}
,
M_{A}
=
u_{s}
/
v_{A}
(where
is the Alfvén speed), and the shock obliquity angle, Θ
_{Bn}
. For example, the electron distributions measured near interplanetary shocks can be fitted with the
κ
distributions with
κ_{e}
~ 2 − 5, while the proton distributions prefer a somewhat larger
κ_{p}
10 (e.g.,
Pierrard & Lazar 2010
).
As an illustration,
Figure 1
compares the electron
κ
 distributions,
f_{e,κ}
(
p
) with
κ_{e}
= 1.6 − 3.0 for
p
≤
p
_{inj}
and the electron powerlaw distribution,
f_{e,κ}
(
p
_{inj}
) · (
p
/
p
_{inj}
) −q for p > pinj. Here the shock parameters adopted are
u_{s}
= 3.2 × 10
^{3}
km s
^{−1}
,
M_{s}
= 3, and the DSA powerlaw slope,
q
= 4.5. The Maxwellian distributions for electrons and protons are also plotted for comparison. These
κ
distributions may represent the preheated electron populations, while the powerlaw distributions may represent the shock accelerated electron populations. We note that preheated electrons are assumed to be injected at
p
=
p
_{inj}
with the amplitude,
f_{e,κ}
(
p
_{inj}
), which is much larger than the corresponding value for the electron Maxwellian distribution,
f_{e,M}
(
p
_{inj}
). Here we adopt
p
_{inj}
≈ 5.34
m_{p}u
_{2}
, which becomes
p
_{inj}
≈ 3.3
p
_{th,p}
for
M_{s}
= 3. The filled circle on the proton Maxwellian distribution marks the amplitude of proton distribution,
f_{p,M}
(
p
_{inj}
), at the injection momentum, which is larger than
f_{e,κ}
(
p
_{inj}
) by several orders of magnitude.
Schematic model demonstrating the preheated suprathermal distributions and the powerlaw distributions of electrons accelerated at a weak shock of u_{s} = 3.2 × 10^{3} km s^{−1} and M_{s} = 3. The Maxwellian distribution (black solid line) and κdistributions with κ_{e} = 1.6 (red dotted), 2.0 (blue dashed), 2.5 (green dotdashed), and 3.0 (magenta dotlong dashed) are shown for electrons. For protons only the Maxwellian distribution (black dotlong dashed line) is shown. The vertical line demarcates the injection momentum, so f_{e}(p) = f_{e,κ}(p) for p < p_{inj} and f_{e}(p) = f_{e,κ}(p_{inj}) · (p/p_{inj})^{−q} for p ≥ p_{inj}. The filled circle marks the amplitude of f_{p,M} (p_{inj}) at the injection boundary for the proton Maxwellian distribution.
For the case shown in
Figure 1
, the ratio of CR electron to proton numbers can be approximated by
K_{e/p}
≈
f_{e,κ}
(
p
_{inj}
)/
f_{p,M}
(
p
_{inj}
), which becomes
K_{e/p}
~ 1/300 for
κ_{e}
= 1.6. So
Figure 1
demonstrates that the ratio becomes
K_{e/p}
~ 10
^{−3}
− 10
^{−2}
, depending on the value of
p
_{inj}
, if the electrons are preaccelerated to
κ
distributions with
κ_{e}
2.
Since the main goal of this study is to explore how the electron energy spectrum and the synchrotron emission are affected by the dynamical evolution of spherically expanding shocks (i.e.,
u_{s}
(
t
) and
r_{s}
(
t
)), the specific value of the amplitude
f_{e,κ}
(
p
_{inj}
) for a given set of models parameters (i.e.,
κ_{e}
and
p
_{inj}
/
p
_{th,p}
) is not important. So we simply assume that the injection momentum decreases with the shock speed as
p
_{inj}
(
t
) ≈ 5.34
m_{p}u_{s}
(
t
)/
σ
, and the electrons are injected with the value
K_{e/p}
= 1 in the simulations described below. Note that we do not attempt to compare the theoretically estimated radio flux,
S_{ν}
, with the observed radio flux of any specific radio relics.
 2.3. Shock Accelerated CR Electron Spectrum
In the testparticle regime of DSA, the distribution of CR distribution function at the shock position can be approximated, once it reaches equilibrium, by a powerlaw spectrum with superexponential cutoff,
where the powerlaw slope is
q
= 3
σ
/(
σ
− 1) and
f
_{inj}
is the amplitude at
p
=
p
_{inj}
(
Kang 2011
). The cutoff momentum can be derived from the equilibrium condition that the DSA momentum gains per cycle are equal to the synchrotron/iC losses per cycle (
Kang 2011
):
The corresponding Lorentz factor for typical cluster shock parameters is then
Postshock electrons cool radiatively while advecting downstream, so the cutoff of the electron momentum spectrum decreases as one moves away from the shock. As a result, the volume integrated electron energy spectrum,
F_{e}
(
p
), becomes steeper than the powerlaw in Equation (9) by one power of the momentum above the ‘break momentum’ (
p
>
p
_{e,br}
). At the shock age
t
, the ‘break Lorentz factor’ can be estimated from the condition
t
_{age}
=
t
_{rad}
=
p
/
b
(
p
):
The synchrotron emission from monoenergetic electrons with
γ_{e}
peaks around the frequency
ν
_{peak}
≈ 0.3(3
eB
/4
πm_{e}c
)
, so it will be useful to have the following relation:
Then the synchrotron radiation spectrum emitted by the electron population with the powerlaw given in Equation (9) also has a powerlaw form of
j_{ν}
∝
ν
^{−αinj}
with
.
For
diffuse
radio structures such as radio halos and radio relics with
B
~ 5
μ
G and the radio flux density,
S_{ν}
~ 5 mJy, the optical depth for synchrotron emission is small for the frequencies considered here (ν > 10 MHz) (
Jones et al. 1974
;
Lang 1999
). So the synchrotron selfabsorption can be ignored in the calculation of synchrotron radiation spectrum of the model shocks.
3. DSA SIMULATION RESULTS
 3.1. Planar versus Spherical Shocks
In this section we compare the electron acceleration at a spherical shock with that at a planar shock with similar properties. First, the planar shock is characterized by the following parameters: the shock speed,
u_{s}
= 3.5 × 10
^{3}
km s
^{−1}
, the sonic Mach number,
M_{s}
= 4.5, the preshock density,
n
_{H,1}
= 10
^{−3}
cm
^{−3}
, and the preshock and postshock magnetic field strength,
B
_{1}
= 2
μ
G (for
x
>
x_{s}
) and
B
_{2}
= 7
μ
G (for
x
< x
_{s}
).
In the upper panel of
Figure 2
, we show the spatial profiles of the electron distribution function,
g_{e}
(
x
,
γ_{e}
) for different energies, log
γ_{e}
= 3.3, 3.7, 4.0, 4.3, 4.5, and 4.8. Note that the shock location is shifted slightly from left to right at later time epochs,
t
_{age}
= (1.8, 4.4, 8.0) × 10
^{7}
yr. The downstream width of
g_{e}
,
L
(
γ_{e}
), for low energy electrons with log
γ_{e}
≲ 3.3 increases simply as
l
_{adv}
= (
u_{s}
/
σ
)
t
_{age}
≈ 100kpc(
t
_{age}
/10
^{8}
yr). For high energy electrons, on the other hand,
L
(
γ_{e}
) decreases with
γ_{e}
, since the radiative cooling time scales inversely with electron energy as in Equation (4). As can be seen in the figure, for
t
_{age}
> 4.4 × 10
^{7}
yr, for instance, the width
L
(
γ_{e}
) for
γ_{e}
≥ 104 becomes constant in time (tage > trad) as L(
γ_{e}
) ≈ (
u_{s}
/
σ
)
t
_{rad}
(
γ_{e}
) (see blue dashed, green dotdashed, magenta dotlong dashed, and cyan dashlong dashed lines).
Time evolution of the planar shock of u_{s} = 3.5 × 10^{3} km s^{−1}, M_{s} = 4.5, and B_{2} = 7 μG is shown for t_{age} = (1.8, 4.4, 8.0)×10^{7} yr (from left to right). For clarity the shock location is shifted to the right by arbitrary amount at different time epochs. Top: Spatial distributions of the electron distribution function, g_{e}(x, γ_{e}), for log(γ_{e}) = 3.3 (black solid lines), 3.7 (red dotted), 4.0 (blue dashed), 4.3 (green dotdashed), 4.5 (magenta dotlong dashed), and 4.8 (cyan dashlong dashed). Middle: Spatial distributions of the synchrotron emissivity, j_{ν}(x, ν) for log(ν) = 7.6 (black solid lines), 8.5 (red dotted), 9.0 (blue dashed), 9.6 (green dotdashed), 10. (magenta dotlong dashed), and 10.7 (cyan dashlong dashed). Bottom: Spectral index α_{log ν1−log ν2}, where (log ν_{1},log ν_{2}) is (7.6,8.5) (black solid line), (8.5,9.0) (red dotted), (9.0,9.6) (blue dashed), (9.6,10.) (green dotdashed), and (10.,10.7) (magenta dotlong dashed).
From Equation (13), the electrons with
γ_{e}
are expected to have the peak of synchrotron emission at
ν
_{peak}
≈ 1GHz(
γ_{e}
/10
^{4}
)
^{2}
. The middle panel of
Figure 2
shows the synchrotron emissivity,
j_{ν}
(
x
,
ν
), for the peak frequency that corresponds to the Lorentz factor chosen in the upper panel (i.e., log
ν
= 7.6, 8.5, 9.0, 9.6, 10., and 10.7), at the same epochs as in the upper panel. Note that the spatial distribution of synchrotron emission at
ν_{peak}
is much broader than that of electrons with the corresponding
γ_{e}
. This is because more abundant lower energy electrons also contribute significantly to the synchrotron emissivity at
ν_{peak}
. For example, the spatial profile of
j_{ν}
(1GHz) (blue dashed lines) is much broader and decreases much gradually, compared to that of
g_{e}
of the corresponding
γ_{e}
= 10
^{4}
(also blue dashed lines) with
L
(
γ_{e}
= 10
^{4}
) ≈ 40 kpc. So the downstream width of synchrotron emitting volume would be greater than that of the spatial distribution of
g_{e}
(
γ_{e}
) for the value of
γ_{e}
related by Equation (13), that is,
L
(
ν_{peak}
) >
L
(
γ_{e}
), if the width is constrained by the cooling time
t
_{rad}
. Note that
g_{e}
and
j_{ν}
are plotted in arbitrary units in
Figure 2
.
The bottom panel of
Figure 2
shows the spectral index which is defined as
estimated between
ν
_{1}
and
ν
_{2}
. The peak frequencies chosen in the middle panel are adopted for
ν
_{1}
and
ν
_{2}
, so
α
_{7.6−8.5}
(black solid line),
α
_{8.5−9.0}
(red dotted),
α
_{9.0−9.6}
(blue dashed),
α
_{9.6−10}
(green dotdashed), and
α
_{10−10.7}
(magenta dotlong dashed) are plotted. The characteristic scale over which
α
_{log ν1−log ν2}
increases is governed mainly by the length scale of
j_{ν}
at the higher frequency
ν
_{2}
. For example, the spectral index
α
_{9.0−9.6}
(blue dashed line in the bottom panel) behaves similarly to the emission
j_{ν}
at log
ν
= 9.6 (green long dashed line in the middle panel).
For the spherical shock, we adopt a SedovTaylor similarity solution with the shock speed,
u_{s,i}
=
u_{o}
(
t_{i}
/
t_{o}
)
^{−3/5}
, as the initial set up. Here
t_{i}
is the initial time from which electrons are injected at the shock and
u_{o}
= 2.3 × 10
^{3}
km s
^{−1}
and
t_{o}
= 1.73 × 10
^{8}
yr. For
t_{i}
= 0.5
t_{o}
, the initial shock speed,
u_{s,i}
≈ 3.5 × 10
^{3}
km s
^{−1}
and
M_{s}
≈ 4.5, which are similar to the parameters for the planar shock shown in
Figure 2
. Since the sonic Mach number is not large and the shock compression ratio is
σ
=
ρ
_{2}
/
ρ
_{1}
≈ 3.5, the shock evolution behaves slightly differently from the true SedovTaylor solution. So at later time the shock speed evolves only approximately as
u_{s}
(
t
) ≈
u_{o}
(
t
/
t_{o}
)
^{−3/5}
. Again note that only the structure and evolution of the flow velocity
u
(
r
,
t
), is fed into the diffusion convection Equation (1) and the other gasdynamical variables do not influence the DSA process. For the comparison with the planar shock, the magnetic field strengths are assumed to be constant in both space and time with
B
_{1}
= 2
µ
G and
B
_{2}
= 7
µ
G. In the next section, we will consider several models with variable magnetic field strength,
B
(
r
,
t
), behind the shock.
The characteristics of the spherical shock model shown in
Figure 3
can be compared directly with those of the planar shock shown in
Figure 2
. Here the results are shown at the time elapsed from the initial injection,
t
_{age}
≡ (
t
−
t_{i}
) = (1.7, 4.3, 8.7)×10
^{7}
yr, during which the shock speed decreases from 3.5×10
^{3}
km s
^{−1}
(
M_{s}
= 4.5) to 1.5 × 10
^{3}
km s
^{−1}
(
M_{s}
= 3.1). The obvious feature, which is different from the planar shock case, is that the injected particle flux decreases as the spherical shock slows down in time. As a result, the amplitude of
g_{e}
(
r
,
γ_{e}
) increases downstream away from the shock at a given time, especially in the case of low energy electrons, which simply advect downstream without much cooling. So low frequency (
ν
360 MHz) radio emission may increase towards downstream as well. For high energy electrons, on the other hand, the decline of
g_{e}
(
r
,
γ_{e}
) due to cooling dominates over the effect of decreasing injection flux. Thus radio emission at higher frequencies (
ν
1 GHz) should decrease downstream behind the shock. Again note that
g_{e}
and
j_{ν}
are plotted in arbitrary units in
Figure 3
.
Time evolution of the spherical shock of u_{s} = 2.3 × 10^{3} km s^{−1} (t/t_{o})^{−3/5}, M_{s} = 3.0(t/t_{o}) ^{−3/5}, and B_{2} = 7 µG (where t_{o} = 1.73 × 10^{8} yr) is shown for t_{age} = (1.7, 4.3, 8.7) × 10^{7} yr (from left to right). Electrons are injected at the shock from t_{i} = 0.5t_{o} when u_{s,i} = 3.5 × 10^{3} km s^{−1}. Top: Spatial distributions of g_{e}(r, γ_{e}). Middle: Spatial distributions of j_{ν}(r, ν). Bottom: Spectral index α_{log ν1−log ν2}. The line types are the same as in Figure 2.
Similar to the planar shock model, the spatial distribution of synchrotron emission at
ν_{peak}
is much broader than that of
g_{e}
(
γ_{e}
) for the corresponding
γ_{e}
. For example, again the width of
g_{e}
,
L
(
γ_{e}
= 10
^{4}
) ≈ 40 kpc (blue dashed line in the top panel), while the width of
j_{ν}
,
L
(
ν
≈ 1GHz) ≈ 100 kpc (blue dashed line in the middle panel). But the spectral index
α
_{9.0−9.6}
between 1 GHz and 4 GHz (blue dashed line in the bottom panel) follows the spatial profile of the emission
j_{ν}
at 4 GHz (green dotdashed line in the middle panel). Thus the main difference between a planar shock and a decelerating spherical shock is the decline of the injected particle flux in time, leading to possible increase of low frequency radio emission downstream away from the shock.
Next we attempt to find any spectral signatures imprinted in the synchrotron emission spectrum in the case of the spherical shock case. In the left panels of
Figure 4
, we show, at the three epochs, the electron spectrum at the shock,
g_{e}
(
x_{s}
,
p
), and the volume integrated electron spectrum
G_{e}
(
p
) =
p
^{4}
F_{e}
(
p
), where
F_{e}
(
p
) =
f_{e}
(
p
,
x
)
dx
, and their slopes,
q
= −
d
ln
f_{e}
(
x_{s}
)/
d
ln
p
, and Q = −d ln Fe/d ln p. The slope from the simulation results is consistent with the testparticle slope
q
= 4.2 for
σ
= 3.5, and the slope for the volume integrated spectrum,
Q
=
q
+1, above the break momentum
p
_{e,br}
. Note that
γ_{e}
=
p
/
m_{e}c
for relativistic energies. As can be seen in
Figure 4
, the electron energy spectrum at the shock has reached steady state at these ages. The volume integrated spectrum shows the gradual decline of the break Lorentz factor,
γ
_{e,br}
with the shock age.
Electron (left) and synchrotron (right) spectral properties for the planar shock shown in Figure 2 are plotted for t_{age} = (1.8, 4.4, 8.0) × 10^{7} yr (solid, dotted, and dashed lines, respectively). Upper left: electron distribution function at the shock position, g_{e}(x_{s}, p) = p^{4}f_{e}(x_{s}, p), and integrated over the simulation volume, G_{e}(p) = p^{4}F_{e}(p). Lower left: electron distribution slopes, q = −d ln f_{e}(x_{s})/d ln p and Q = −d ln F_{e}/d ln p. Upper right: synchrotron radiation spectrum at the shock position, νj_{ν}(x_{s}), and integrated over the simulation volume, νJ_{ν}. Lower right: synchrotron spectral indexes, α_{ν} = −d ln j_{ν}(x_{s})/d ln ν and A_{ν} = −d ln J_{ν}/d ln ν.
The same kinds of quantities for the spherical shock are shown in the left panels of
Figure 5
. We can see the time evolution of the electron spectrum as the shock slows down in time. Since the shock Mach number decreases, the slope
q
for
g_{e}
(
r_{s}
,
p
) increases in time, leading to a nonlinear change of the volume integrated spectrum,
G_{e}
(
p
). So the slope
Q
no longer follows the simple
q
+1 steepening above
p
_{e,br}
. In addition to usual radiative cooling, both the reduced injected particle flux and steepening of
g_{e}
(
r_{s}
,
p
) due to the slowingdown and weakening of the shock affect the spectral shape of
G_{e}
(
p
).
Same as Figure 4 except that the results for the spherical shock shown in Figure 3 are plotted for t_{age} = (1.7, 4.3, 8.7) × 10^{7} yr (solid, dotted, and dashed lines, respectively).
In addition, we show the synchrotron emission spectrum at the shock,
j_{ν}
(
x_{s}
) in
Figure 4
and
j_{ν}
(
r_{s}
) in
Figure 5
, and the volume integrated emission spectrum,
J_{ν}
=
j_{ν}
(
x
)
dx
in
Figure 4
and
J_{ν}
=
j_{ν}
(
r
)4
πr
^{2}
dr
in
Figure 5
. Their respective spectral indexes,
α_{ν}
= −
d
ln
j_{ν}
/
d
ln
ν
and
A_{ν}
= −
d
ln
J_{ν}
/
d
ln
ν
are shown in the lower right panels of
Figures 4
and
5
. For the planar shock with
σ
= 3.5, the testparticle spectral indexes are
α_{ν}
= 0.6 and
A_{ν}
= 1.1 for
ν
>
ν
_{br}
, as can be seen in
Figure 4
. We note that the transition of the volumeintegrated spectral index from
A_{ν}
=
α_{ν}
to
A_{ν}
=
α_{ν}
+ 0.5 is more gradual, compared to the transition of the particle slope from
Q
=
q
to
Q
=
q
+ 1. For the shock parameters chosen here, at the shock age of ~ 10
^{8}
yr this transition occurs over the broad frequency range, 0.01 GHz
ν
10 GHz.
For the spherical shock shown in
Figure 5
, on the other hand, the volume integrated spectrum,
J_{ν}
, develops a concave curvature and its spectral index
A_{ν}
deviates from the simple relation of
A_{ν}
=
α_{ν}
+ 0.5. Again, at the shock age of ∼ 10
^{8}
yr the nonlinear curvature occurs over the broad frequency range, 0.01 GHz
ν
10 GHz. Thus the simple relation,
A
_{integ}
≈
α
_{inj}
+ 0.5, which is often invoked for the DSA origin of synchrotron emission spectra, should be applied with some caution in the case of evolving shocks. Again in
Figures 4
and
5
,
g_{e}
,
G_{e}
,
j_{ν}
and
J_{ν}
all are plotted in arbitrary units.
 3.2. Spherical Shocks with Different Magnetic Field Models
Selfexcitation of MHD waves and amplification of magnetic fields via plasma instabilities are the integral parts of DSA at
strong
collisionless shocks. CRs streaming upstream in the shock precursor excite resonant Alfvén waves with wavelengths comparable with the CR proton gyroradii, and turbulent magnetic fields can be amplified into the nonlinear regime (
Bell 1978
;
Lucek & Bell 2000
). In addition, the nonresonant fastgrowing instability driven by the CR current can amplify the magnetic field by orders of magnitude (
Bell 2004
;
Riquelme & Spitkovsky 2009
). It was shown that, in the case of strong shocks (
M_{s}
> 10), magnetic ampli fication (MFA) via streaming instabilities and Alfvénic drift can lead to significant nonlinear feedback on the shock structure and the predicted CR spectrum (e.g.,
Caprioli 2012
;
Kang 2012
).
In a recent study of hybrid plasma simulations,
Caprioli & Spitkovsky (2014b)
has shown that the MFA factor increases with the Alfvenic Mach number as 〈
δB
/
B
〉
^{2}
~
, where
δB
is the turbulent magnetic fields perpendicular to the mean background magnetic fields. For the sonic Mach number in the range of 3
Ms
5, the CR proton acceleration efficiency varies as
~ 0.01 − 0.05 (
Kang & Ryu 2013
;
Caprioli & Spitkovsky 2014a
). For typical cluster shocks with
M_{A}
~ 10, the MFA factor due to the streaming stabilities is expected be rather small, 〈
δB
/
B
〉
^{2}
~ 0.3 − 1.5. How magnetic fields may be amplified at weak shocks in high beta ICM plasmas has not yet been fully understood. As in the previous study of
Kang et al. (2012)
, we adopt the postshock magnetic field strength
B
_{2}
~ 7
µ
G in order to produce the observed width of the Sausage relic (∆
l
~ 50 kpc), while the preshock magnetic field strength is chosen to be
B
_{1}
~ 2 − 3
µ
G. Note that
B
_{1}
cannot be constrained directly from observations, since we do not detect any radio emission upstream of radio relic shocks. Thus
B
_{1}
can be adjusted so that
B
_{2}
becomes about 7
µ
G after considering MFA via plasma instabilities or the compression of the perpendicular components of magnetic fields across the shock. In any case, the postshock electron energy spectrum and the synchrotron radiation spectrum do not depend sensitively on the value of
B
_{1}
as long as it is a few microgauss.
Here we consider the following three magnetic field models:

MF1:B1= 2µG &B2= 7µG.

MF2:B1= 3µG,B2=&B(r) =B2· (P(r)/P2)1/2forr

MF3:B1= 3µG,B2=&B(r) =B2· (ρ(r)/ρ2) forr
In
MF2
and
MF3
models, the postshock magnetic field energy is assumed to be dissipated, and so
B
(
r
) decreases with the gas pressure or density behind the shock. As shown in the bottom panels of
Figure 6
, the decay length scale is about 100150 kpc in
MF2
and
MF3
models.
Spherical shock models with three different magnetic field profiles: MF1, MF2, and MF3. The results are shown for t_{age} = 1.3×10^{8} yr, when u_{s} = 2.5×10^{3} km s^{−1} and M_{s} = 2.4. Spatial distributions of the electron distribution function, g_{e}(r, γ_{e}), the synchrotron emissivity, j_{ν}(r, ν), the spectral index, α_{log ν1−log ν2}, and magnetic field strength, B(r) (from top to bottom panels). The line types for g_{e}, j_{ν}, and α_{log ν1−log ν2} are the same as in Figures 2 and 3. Note that g_{e} and j_{ν} are plotted in arbitrary units.
As in the previous section, we consider a spherical blastwave,
u_{s,i}
≈
u_{o}
(
t_{i}
/
t_{o}
)
^{−3/5}
, where
u_{o}
= 3.0 × 10
^{3}
km s
^{−1}
and
t_{o}
= 1.3 × 10
^{8}
yr (slightly faster than the shock adopted in
Figure 3
). The calculations were started at
t_{i}
= 0.5
t_{o}
with
u_{s,i}
≈ 4.5 × 10
^{3}
km s
^{−1}
and
M_{s}
≈ 4.3 and ended at
t_{f}
= 1.5
t_{o}
with
u_{s,f}
≈ 2.5 × 10
^{3}
km s
^{−1}
and
M_{s}
≈ 2.4. The results of these three models at the final time (
t
_{age}
=
t
−
t_{i}
≈ 1.3 × 10
^{8}
yr) are compared in
Figure 6
. As in the spherical shock model shown in
Figure 3
, the spatial profile of
g_{e}
for low energy electrons increases downstream away from the shock. The distributions of
g_{e}
(
γ_{e}
≈ 5 × 10
^{3}
) (red dotted lines) show some differences among the three models due to different synchrotron cooling rates. The spatial profiles of
j_{ν}
at low frequencies (e.g., 360 MHz in red dotted lines) show rather weak dependences on the
B
(
r
) model, while those at higher frequencies (e.g., 1 GHz in blue dashed lines) behave similarly in all three cases. Moreover, the steepening patterns of
α_{ν}
behind the shock show no significant variations among the three models.
Thus it may be difficult to differentiate the postshock magnetic field profile from the spatial variation of spatially resolved radio flux,
S_{ν}
, or spectral index,
α_{ν}
. In other words, signatures imprinted on synchrotron emission due to different
B
(
r
) behind the shock would be rather subtle to detect. There are two reasons that may explain this prediction. Because the iC cooling with
B
_{rad}
= 4.7
µ
G provides the baseline for the cooling rate,
b
(
p
), the difference in the synchrotron cooling due to different
B
(
r
) has relatively minor effects. Also any sharp features in the distribution of
f_{e}
(
r
,
p
) would be smeared out in the distribution
j_{ν}
(
r
), since the synchrotron emission at a given frequency comes from electrons with a somewhat broad range of
γ_{e}
.
Finally, in
Figure 7
we compare the synchrotron radiation spectrum at the shock,
j_{ν}
(
r_{s}
), and the volume integrated radiation spectrum,
J_{ν}
, for the three magnetic field models at
t
_{age}
= (0.33, 0.67, 1.3) × 10
^{8}
yr. Departures from the predictions for the testparticle planar shock are most severe in
MF1
model, while it becomes relatively milder in the two models with decaying
B
(
r
) behind the shock. As expected, any variations in the spatial profiles of
f_{e}
(
r
,
p
) and
B
(
r
) are averaged out in the volume integrated quantities such as
F_{e}
(
p
) and
J_{ν}
.
Spherical shock models with three different magnetic field profiles shown in Figure 6. Top: synchrotron spectral distribution at the shock position, νj_{ν}(r_{s}), and integrated over the simulation volume, νJ_{ν}. Bottom: synchrotron spectral index, α_{ν} for j_{ν}(r_{s}) and A_{ν} for J_{ν} . The results are shown for t_{age} = (0.33, 0.67, 1.3) × 10^{8} yr (solid, dotted, dashed lines, respectively). Note that j_{ν} and J_{ν} are plotted in arbitrary units.
In conclusion, any features in the spectral curvature,
dA_{ν}
/
d_{ν}
, or the steepening of spectral index,
dα_{ν}
/
dr
in the postshock region, should be affected in principle by both the timedependent evolution of the shock and the spatial variation of postshock magnetic fields. But the effect of decaying magnetic fields may be somewhat subtle to detect due to the dominance of iC scattering off the CBR.
4. SUMMARY
We have performed timedependent DSA simulations for cosmicray (CR) electrons at a planar shock with a constant speed and decelerating spherical shocks similar to SedovTaylor blast waves. These shocks were chosen to have parameters relevant for weak shocks typically found in the outskirts of galaxy clusters:
u_{s}
≈ (1.5−4.5)×10
^{3}
km s
^{−1}
,
M_{s}
≈ 2.5−4.5,
B
_{1}
≈ 2−3
µ
G, and
B
_{2}
≈ 7
µ
G. For spherical shocks, we used a onedimensional spherical version of the CRASH code, in which a coexpanding spherical grid was adopted in order to optimize the shock tracking and adaptive mesh refinement features (
Kang & Jones 2006
). The diffusion convection equation for the pitchangleaveraged phase space distribution function for CR electrons in the testparticle regime was solved along with the usual gasdynamic conservation equations without including the dynamical feedback of the CR proton pressure. The synchrotron emission from CR electrons accelerated at these model shocks was estimated by adopting several models for magnetic field profile,
B
(
r
), and the electron energy spectrum,
N_{e}
(
r
,
γ_{e}
), from the DSA simulations.
Since the time scales for electron acceleration and cooling are much shorter than the dynamical time scale, the electron energy spectrum at the shock,
f_{e}
(
r_{s}
,
p
,
t
), has reached the steady state in the planar shock model, as can be seen in
Figure 4
. The slope of the volumeintegrated momentum distribution,
F_{e}
(
p
), is simply
Q
=
q
+ 1 above the break momentum (
p
>
p
_{e,br}
). However, the spectral index of the volumeintegrated synchrotron emission spectrum,
J_{ν}
, steepens gradually to
A_{ν}
=
α
_{inj}
+ 0.5 over the broad frequency range, ~ (0.1 − 10)
ν
_{br}
. They are consistent with the predictions of the DSA theory in the testparticle limit.
In the spherical shock cases, the electron energy spectrum at the shock has reached the steady state defined by the
instantaneous
shock parameters (i.e.,
u_{s}
(
t
),
r_{s}
(
t
) and
M_{s}
(
t
)), as shown in
Figure 5
. So the synchrotron radiation spectra at the shock location could be described properly by the testparticle DSA predictions. However, the volume integrated spectra of
F_{e}
(
p
,
t
) and
J_{ν}
(
ν
,
t
) both evolve differently from those of the planar shock, depending on the timedependent evolution of the shock parameters. As a result, the slopes,
Q
and
A_{ν}
, exhibit some nonlinear signatures that come from decreasing particle flux and declining shock Mach number (see
Figure 5
). This suggests that, in the case of evolving shocks, one needs to be cautious about interpreting observed radio spectra by adopting simple DSA models in the testparticle regime.
In addition, we have considered spherical shock models with decaying postshock magnetic fields, in which
B
(
r
) decreases behind the shock with a scale height of 100150 kpc, as shown in
Figure 6
. Because inverseCompton scattering off the cosmic background photons with
B
_{rad}
≈ 4.7
µ
G (for
z
= 0.2 adopted here as a reference epoch) provides the baseline cooling rate, the electron energy spectra do not depend sensitively on the B(r) profile. We also find that the impacts of the different
B
(
r
) profile on the spatial profile of synchrotron emission,
j_{ν}
(
r
), and its spectral index,
α_{ν}
(
r
), is rather minimal, since electrons in a somewhat broad range of
γ_{e}
contribute emission to
j_{ν}
at a given frequency. Of course, any nonlinear features due to the spatial variations in
f_{e}
(
r
,
p
) and
B
(
r
) are mostly averaged out and may leave only subtle signatures in the volume integrated spectrum,
J_{ν}
, and its spectral index,
A_{ν}
.
Acknowledgements
This work was supported by a 2Year Research Grant of Pusan National University.
Akamatsu H.
,
Kawahara H.
2013
Systematic XRay Analysis of Radio Relic Clusters with Suzaku
PASJ
65
16 
Brunetti G.
,
Jones T. W.
2014
Cosmic Rays in Galaxy Clusters and Their Nonthermal Emission
Int. J. of Modern Physics D.
23
30007 
Caprioli D.
2012
CosmicRay Acceleration in Supernova Remnants: NonLinear Theory Revised
JCAP
7
38 
Caprioli D.
,
Sptikovsky A.
2014
Simulations of Ion Acceleration at Nonrelativistic Shocks. I. Acceleration Efficiency
ApJ
783
91 
DOI : 10.1088/0004637X/783/2/91
Caprioli D.
,
Sptikovsky A.
2014
Simulations of Ion Acceleration at Nonrelativistic Shocks. II. Magnetic Field Amplification
ApJ
794
46 
DOI : 10.1088/0004637X/794/1/46
Caprioli D.
,
Pop A. R.
,
Sptikovsky A.
2015
Simulations and Theory of Ion Injection at Nonrelativistic Collisionless Shocks
ApJ
798
28 
Drury L. O’C.
1983
An Introduction to the Theory of Diffusive Shock Acceleration of Energetic Particles in Tenuous Plasmas
Rep. Prog. Phys.
46
973 
DOI : 10.1088/00344885/46/8/002
Feretti L.
,
Giovannini G.
,
Govoni F.
,
Murgia M.
2012
Clusters of Galaxies: Observational Properties of the Diffuse Radio Emission
A&A Rev.
20
54 
Guo F.
,
Giacalone J.
2013
The Acceleration of Thermal Protons at Parallel Collisionless Shocks: ThreeDimensional Hybrid Simulations
ApJ
773
158 
DOI : 10.1088/0004637X/773/2/158
Guo X.
,
Sironi L.
,
Narayan R.
2014
Nonthermal Electron Acceleration in Low Mach Number Collisionless Shocks. I. Particle Energy Spectra and Acceleration Mechanism
ApJ
793
153 
Jones T. W.
,
O’dell S. L.
,
Stein W. A.
1974
Physics of Compact Nonthermal Sources. I. Theory of Radiation Processes
ApJ
188
353 
DOI : 10.1086/152724
Kang H.
2011
Energy Spectrum of Nonthermal Electrons Accelerated at a Plane Shock
JKAS
44
39 
Kang H.
2012
Diffusive Shock Acceleration with Magnetic Field Amplification and Alfvénic Drift
JKAS
45
127 
Kang H.
,
Jones T. W.
,
Gieseler U. D. J.
2002
Numerical Studies of CosmicRay Injection and Acceleration
ApJ
579
337 
DOI : 10.1086/342724
Kang H.
,
Ryu D.
2013
Diffusive Shock Acceleration at Cosmological Shock Waves
ApJ
756
97 
Kang H.
,
Ryu D.
,
Cen R.
,
Ostriker J. P.
2007
Cosmological Shock Waves in the LargeScale Structure of the Universe: Nongravitational Effects
ApJ
669
729 
DOI : 10.1086/521717
Kang H.
,
Vahe P.
,
Ryu D.
,
Jones T. W.
2014
Injection of κlike Suprathermal Particles into Diffusive Shock Acceleration
ApJ
788
141 
DOI : 10.1088/0004637X/788/2/141
Lang K. R.
1999
Astrophysical Formulae, Vol. I
Springer
Radiation, Gas Processes and High Energy Astrophysics
Malkov M. A.
,
Drury L.O’C.
2001
Nonlinear Theory of Diffusive Acceleration of Particles by Shock Waves
Rep. Progr. Phys.
64
429 
DOI : 10.1088/00344885/64/4/201
Markevitch M.
,
Gonzalez A. H.
,
David L.
,
Vikhlinin A.
,
Murray S.
,
Forman W.
,
Jones C.
,
Tucker W.
2002
A Textbook Example of a Bow Shock in the Merging Galaxy Cluster 1E 065756
ApJ
567
27 
DOI : 10.1086/339619
Petrosian V.
2001
On the Nonthermal Emission and Acceleration of Electrons in Coma and Other Clusters of Galaxies
ApJ
557
560 
DOI : 10.1086/321557
Pierrard V.
,
Lazar M.
2010
Kappa Distributions: Theory and Applications in Space Plasmas
Sol. Phys.
265
153 
Pinzke A.
,
Oh S. P.
,
Pfrommer C.
2013
Giant Radio Relics in Galaxy Clusters: Reacceleration of Fossil Relativistic Electrons?
MNRAS
435
1061 
DOI : 10.1093/mnras/stt1308
Ryu D.
,
Kang H.
,
Hallman E.
,
Jones T. W.
2003
Cosmological Shock Waves and Their Role in the LargeScale Structure of the Universe
ApJ
593
599 
DOI : 10.1086/376723
Sarazin C.
1999
The Energy Spectrum of Primary CosmicRay Electrons in Clusters of Galaxies and Inverse Compton Emission
ApJ
520
529 
DOI : 10.1086/307501
Schlickeiser R.
2002
Cosmic Ray Astrophysics
Springer
Berlin
Skillman S. W.
,
Hallman E. J.
,
O’Shea W.
,
Burns J. O.
,
Smith B. D.
,
Turk M. J.
2011
Galaxy Cluster Radio Relics in Adaptive Mesh Refinement Cosmological Simulations: Relic Properties and Scaling Relationships
ApJ
735
96 
DOI : 10.1088/0004637X/735/2/96
van Weeren R.
,
Röttgering H. J. A.
,
Brüggen M.
,
Hoeft M.
2010
Particle Acceleration on Megaparsec Scales in a Merging Galaxy Cluster
Science
330
347 
DOI : 10.1126/science.1194293
van Weeren R.
,
Hoeft M.
,
Röttgering H. J. A.
,
Brüggen M.
,
Intema H. T.
,
van Velzen S.
2011
A Double Radio Relic in the Merging Galaxy Cluster ZwCl 0008.8+5215
A&A
528
A38 
38
Vasyliunas V. M.
1968
A Survey of LowEnergy Electrons in the Evening Sector of the Magnetosphere with OGO 1 and OGO 3
JGR
73
2839 
DOI : 10.1029/JA073i009p02839
Vazza F.
,
Brunetti G.
,
Gheller C.
2009
Shock Waves in Eulerian Cosmological Simulations: Main Properties and Acceleration of Cosmic Rays
MNRAS
395
1333 
DOI : 10.1111/j.13652966.2009.14691.x