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Theoretical and Mathematical Physics, 170(2): 249262 (2012)

AXIONS AND COSMIC RAYS

D. Espriu and A. Renau

We investigate the propagation of a charged particle in a spatially constant but time-dependent pseudo-

scalar background. Physically, this pseudoscalar background could be provided by a relic axion density.

The background leads to an explicit breaking of Lorentz invariance; processes such as p p or e eare consequently possible under some kinematic constraints. The phenomenon is described by the QED

Lagrangian extended with a ChernSimons term that contains a four-vector characterizing the breaking

of Lorentz invariance induced by the time-dependent background. While the induced radiation (similar to

the Cherenkov eect) is too small to inuence the propagation of cosmic rays signicantly, the hypothetical

detection of the photons radiated by high-energy cosmic rays via this mechanism would provide an indirect

way to verify the cosmological relevance of axions. We discuss the order of magnitude of the eect.

Keywords: axion, high-energy cosmic ray, galactic magnetic eld

1. Axions

Cold relic axions resulting from vacuum misalignment [1], [2] in the early universe is a popular and sofar viable candidate for dark matter. If we assume that cold axions are the only contributors to the matterdensity of the universe apart from ordinary baryonic matter, then its density must be [3]

1030 g/cm3 1046 GeV4.Of course, dark matter is not uniformly distributed: its distribution traces that of visible matter (or ratherthe other way around). The galactic halo of dark matter (assumed to consist of axions) would correspondto a typical value for the density [4]

a 1024 g/cm3 1040 GeV4,extending over a distance of 30 to 100 kpc in a galaxy such as the Milky Way. The precise details of thedensity prole are not so important at this point. The axion background provides a very diuse concen-tration of pseudoscalar particles interacting very weakly with photons and therefore indirectly with cosmicrays. What are the consequences of this diuse axion background on high-energy cosmic ray propagation?Could this inuence cosmic ray propagation similarly to the GreisenZatsepinKuzmin (GZK) cuto [5]?This is the question that we answer here.

It is quite relevant that the axion is a pseudoscalar, being the pseudo-Goldstone boson of the brokenPecceiQuinn symmetry [6]. Its coupling to photons occurs through the anomaly term; the coecient ishence easily calculable if the axion model is known:

L = ga 2a

faFF, (1)

CERN, Geneva, Switzerland, e-mail: domenec.espriu@cern.ch.Department of Structure and Constituents of Matter, Institute of Cosmos Sciences, University of Barcelona,

Barcelona, Spain, e-mail: arencer@gmail.com.

Prepared from an English manuscript submitted by the authors; for the Russian version, see Teoreticheskaya i

Matematicheskaya Fizika, Vol. 170, No. 2, pp. 304320, February, 2012.

0040-5779/12/1702-0249

c

249

Two popular axion models are the model in [7] and the model in [8], in both of which ga 1. We providesome additional details of the theory in Sec. 3.

2. Cosmic rays

Cosmic rays consist of particles (such as electrons, protons, helium nuclei, and other nuclei) reachingthe Earth from outside. Primary cosmic rays are those produced by astrophysical sources (e.g., supernovas),while secondary cosmic rays are particles produced by the interaction of primaries with interstellar gas. Westudy the eect of axions on the propagation of these cosmic rays. We separately consider proton andelectron cosmic rays and ignore heavier nuclei because the eect on them is far less important, as becomesclear later (the axion-induced bremsstrahlung depends on the mass of the charged particle).

2.1. Cosmic ray energy spectrum. The number of protons in cosmic rays is an important char-acteristic. Experimental data indicate that the number of cosmic ray particles with a given energy dependson energy according to a power law J(E) = NiEi , where the spectral index i takes dierent values indierent regions of the spectrum [9]. For protons, we have

Jp(E) =

5.87 1019E2.68 eV1m2s1sr1, 109 eV E 4 1015 eV,6.57 1028E3.26 eV1m2s1sr1, 4 1015 eV E 4 1018 eV,2.23 1016E2.59 eV1m2s1sr1, 4 1018 eV E 2.9 1019 eV,4.22 1049E4.3 eV1m2s1sr1, E 2.9 1019 eV,

(2)

while the power law for electrons is [10]

Je(E) =

5.87 1017E2.68 eV1m2s1sr1, E 5 1010 eV,4.16 1021E3.04 eV1m2s1sr1, E 5 1010 eV,

(3)

with the ux typically two orders of magnitude below the proton ux, although it is known less well. Ourignorance about electron cosmic rays is quite regrettable because it has a substantial impact in our estimateof the radiation yield.

We note that the above values measured locally in the inner solar system. It is known that the intensityof cosmic rays increases with distance from the sun because the modulation due to the solar wind makes itmore dicult for them to reach us, particularly so for electrons. In addition, the hypothesis of homogeneityand isotropy holds for proton cosmic rays but not necessarily for electron cosmic rays. Indeed, becausecosmic rays are deected by magnetic elds they follow a nearly random trajectory within the Galaxy.We know that a hadronic cosmic ray spends an average of about 107 years in the galaxy before escapinginto intergalactic space. This ensures the uniformity of the ux, at least for protons of galactic origin. Incontrast, electron cosmic rays travel for an average of about 1 kpc before being slowed down. But becausel t1/2 for a random walk, 1 kpc corresponds to a typical age of an electron cosmic ray of 105 yr [11]. Inaddition, the lifetime of an electron cosmic ray depends on the energy as

t(E) 5 105(

1TeVE

)

yr =T0E

, T0 2.4 1040.

To complicate matters further, it has been argued that the local interstellar electron ux is not evenrepresentative of the Galactic ux and may reect the electron debris from a nearby supernova 104 yearsago [12].

250

2.2. The GZK cuto. It was stated in [5] that the number of cosmic rays above a certain energythreshold (we call this the GZK limit) should be very small. Cosmic rays particles interact with photons fromthe cosmic microwave background (CMB) to produce pions CMB + p p + 0 or CMB + p n + +.The energy threshold is about 1020 eV. Because of the mean free path associated with these reactions,cosmic rays with energies above the threshold and traveling over distances larger than 50Mpc should notbe observed on the Earth. This is the reason for the rapid fall o of the proton cosmic ray spectrum above1020 eV because there are very few nearby sources capable of providing such tremendous energies.

We note that the change in the slope of the spectrum at around 1018 eV is believed to be due to theappearance of extragalactic cosmic rays at that energy.

3. Solving QED in a cold axion background

In this section, in great detail, we describe the theoretical tools needed for understanding the interac-tions between the high-energy cosmic rays we have just described and the cold axion background describedin Sec. 1.

The coupling of axions to photons is described by the following part of the Lagrangian:

La = ga 2a

faF F , (4)

whereF =

12F

is the dual eld strength tensor.The axion eld is originally in a nonequilibrium state, and in the process of its relaxation to the

equilibrium conguration, coherent oscillations with q = 0 are produced if the reheating temperature afterination is below the PecceiQuinn transition scale [6]. In late times, the axion eld evolves according toa(t) = a0 cos(mat), where the amplitude a0 is related to the initial misalignment angle. Lagrangian (4)then becomes

La = ga 21fa

a0 cos(mat)F F = ga

faa0 cos(mat)AF .

Integrating by parts (dropping total derivatives) and taking into account that F = 0, we obtain

La = ga maa0fa

sin(mat)ijkAiFjk, (5)

where the Latin indices run over only the spatial components.A cosmic ray particle (which travels at almost the speed of light) sees regions with quasiconstant values

of the axion background of a size depending on the axion mass but always many orders of magnitude largerthan its wavelength. Therefore, we can approximate the sine in (5) by a constant (1/2, for example). Itcan then be written as

La = 14AF , = (, 0, 0, 0), = 4ga

maa0fa

. (6)

The constant changes sign with a period 1/ma.The oscillator has the energy density a = a2max/2 = (maa0)

2/2, and hence maa0 =

2a. Theconstant is then

= ga4

2afa

1020 eV

for a = 104 eV and fa = 107 GeV = 1016 eV.The extra term in (6) corresponds to the MaxwellChernSimons electrodynamics. Although we can

then in principle have any four-vector , the axion background provides a purely temporal vector. Weassume that is constant within a time interval 1/ma.

251

3.1. EulerLagrange equations. In the presence of an axion background, the QED Lagrangian is

L = 14FF + (i e Ame) + 12m

2AA

+14A F

, (7)

where A = A and similarly for other slashed symbols. Here, we also introduce the eective photon withthe mass m (equivalent to a refractive index [2]). We can write the approximation m2 4ne/me for it.The electron density in the Universe is expected to be at most ne 107 cm3 1021 eV3. This densitycorresponds to m 1015 eV, but we here use the more conservative limit m = 1018 eV (compatiblewith [13]).

The second term in (7) gives the kinetic and mass term for the fermions and also their interaction withphotons. Dropping it, we obtain the Lagrangian for (free) photons in the axion background (see [14] forfurther details):

L = 14FF +

12mAA

+14A F

=

= 12A(A A) + 12m

2AA

+14AA .

The EulerLagrange (EL) equations are

L

(A) L

A= 0,

where

L

(A)=

(A)

[

12

Agg(A A) + 14

AA

]

=

=

{

12[gg(A A) +

+ A(gg gg)] + 14 A

}

=

=

[

(A A) + 14A

]

=

= A + A + 14 A ,

LA

=

A

(

12m2g

AA +14AA

)

=

=12m2(g

A + gA) +14A = m2A

+14A .

Rearranging the indices, we bring the equations to the form

A + A m2A 12 A = 0.

If we choose the Lorenz gauge A = 0, then the second term vanishes. The equations can also be writtenas

gA gm2A 12 A = 0.

252

We are interested in writing these equations in the momentum space. For this, we dene the Fouriertransform of the eld:

A(x) =

d4k

(2)4eikxA(k).

The relevant derivatives are written as

A =

d4k

(2)4(ik)eikxA(k), A =

d4k

(2)4(k2)eikxA(k).

The EL equations are then

d4k

(2)4

[

g(k2 m2) +i

2k

]

eikxA(k) = 0.

Therefore,[

g(k2 m2) +i

2k

]

A(k) = 0

or, equivalently,

KA(k) = 0, K = g(k2 m2) +i

2k . (8)

3.2. Polarization vectors and dispersion relation. We now dene

S = k

k.

This can be more conveniently expressed using the contraction of two Levi-Civita symbols =3![ ] (the minus sign is here because 0123 = 0123 in Minkowski space):

S = [( k)2 2k2]g ( k)(k + k) + k2 + 2kk .

It satisesS

= S k = 0, S = S = 2[( k)2 2k2], SS =

S

2S .

If = (, 0, 0, 0), then we have S = 22k 2 > 0.We now introduce two projectors:

P =S

S i

2Sk .

These projectors have the properties

P = P k = 0, gP

= 1, (P

)

= P = P ,

P P = P , P

P = 0, P

+ + P

= 2

S

S.

With these projectors, we can build a pair of polarization vectors to solve (8). We start from a spacelikeunit vector, for example, = (0, 1, 1, 1)/

3. We then project it: = P . To obtain a normalized

vector, we need

() = P

P

= P =

SS

=

=S/2 + 2( k)2

S= 1

2+

( k)22k2

.

253

(this of course is negative because is spacelike). The polarization vectors are then

=

=

[k2 ( k)22k2

]1/2P .

These polarization vectors satisfy

g

= 1, g = 0, (9)

+

= 2

S

S

= S

2k2. (10)

Using the projectors, we can write the tensor in (8) as

K = g(k2 m2) +

S

2(P P+ ).

For k = (, k), we then have

K =

[

(k2 m2)

S

2

]

= (k2 m2 |k|) =

(

2 k2 m2 |k|)

.

Therefore, A = is a solution of (8) i

(k) =

m2 |k|+ k2.

This is the new dispersion relation of photons in the cold axion background in the approximation where is assumed to be piecewise constant.

4. The process p p4.1. Kinematic constraints. We now consider p(p) p(q)(k) or e(p) e(q)(k). This process

is forbidden in normal QED by the energy conservation, but it is possible in this background (the cold axionbackground even allows the process e+e [15]). Momentum conservation implies that q = p k. Ifm is interpreted as the mass of the charged particle (proton or electron), then conservation of energy leadsto

E(q) + (k) = E(p),

m2 + (p k)2 +

m2 |k|+ k2 =

m2 + p 2,

E2 + k2 2pk cos +

m2 k + k2 E = 0,(11)

where we use the expressions

E = E(p) =

m2 + p 2, p = |p |, k = |k|, pk cos = p k

in the last line. As becomes clear in what follows, if is positive or negative, then the process is onlypossible for the respective negative or positive polarization. Therefore, = || in these cases. To takeboth of them into account, we use the minus sign and write instead of ||.

We square the last equality in (11) twice:

(4E2 4p2 cos2 + 4p cos 2)k2 2(2E2 + 2m2p cos m2)k + 4E2m2 m4 = 0.

254

Neglecting 2k2, m2k, and m4 in this equation, we obtain

(E2 p2 cos2 + p cos )k2 (E2 + m2p cos )k + E2m2 = 0.

This equation has two solutions

k =E2 + pm2 cos E

E22 4E2m2 + 4p2m2 cos2 2pm2 cos 2(E2 p2 cos2 + p cos ) , (12)

which make sense only if the discriminant is positive. In the approximation cos 1 sin2 /2, thecondition 0 becomes

sin2 p22 2pm2 + m2(2 4m2)

4p2m2(1 /4p),

which can be rewritten as

sin2 2

4p2m2

11 /4p(p p+)(p p), (13)

where

p =m2 2mm

1 2

4m2 2mm

1 2

4m2.

It is clear that p+ > 0 and p < 0. For sin2 to be positive, the condition

p > p+ = pth =2mm

1 2

4m2

must be satised. This is the threshold below which the process cannot occur kinematically. The energythreshold E2th = m

2 + p2th is written as

Eth =2mm

.

As 0, the threshold value goes to innity (as expected, the process becomes impossible if vanishes).There is another relevant scale in the problem: m2/. It is many orders of magnitude above the GZK

cuto. Therefore, we always assume the limit p m2/. The maximum angle of emission for a givenmomentum is given by (13). Its largest value is obtained when p is large (p pth):

sin2 max =2

4m2.

Because this is a small number, photons are emitted in a narrow cone max = /2m . This justies theapproximation for cos .

At max(p), the square root in (12) vanishes, and we have

k+[max(p)] = k[max(p)]pthpm2/ 2m

2

.

The minimum value for the angle is = 0:

k(0) E2 + pm2 (E2 pm2 2m2m2/)

2(m2 + p).

255

In the framework of the assumption pth p m2/, we hence obtain the maximum and minimum valuesof the photon momentum:

kmax = k+(0) =E2

m2, (14)

kmin = k(0) =m2

. (15)

These two values coincide at the energy threshold. Here, we can see that the process is possible for negativeor positive polarization only if respectively > 0 or < 0. Otherwise, the modulus of the photon momentumwould be negative.

We note that the incoming cosmic ray wavelength ts perfectly within the 1/ma size, and the quantity therefore turns out to be almost perfectly constant. Whether is positive or negative, there is alwaysa state with slightly less energy into which it is possible to decay with a loss of part of its energy (of theorder O()), emitting a soft photon. Therefore, even if the process is rare, it does not average to zero. Wewill present an exact analysis of this situation in a future publication.

4.2. Amplitude. The next thing we need is to determine is the matrix element for the process. Usingthe standard Feynman rules, we obtain

iM = u(q)ieu(p)(k).

Its square is

|M|2 = u(q)ieu(p)(k)[u(q)ieu(p)(k)] =

= e2(k)(k) tr[u(q)u(q)u(p)u(p) ].

We now must sum and average correspondingly over the initial and nal proton helicities. We do notaverage over photon polarizations, because the process is possible only for one polarization. Taking thetrace, we obtain

|M|2 = 12e2(k)(k) tr[(q + m)(p + m) ] =

=12e2(k)(k) tr[q p + m2] =

=12e2(k)(k)[q

p qpg + qp + m2g ].

Using four-momentum conservation, (9), and the fact that pp = m2, we obtain

|M|2 = 2e2(pk + 2pp) = 2e2[pk + ( + )pp ].

Now using (10), we obtain

( + )p

p = Spp2k2

= p2 sin2 .

The average squared amplitude is then

|M|2 = 2e2(pk + p2 sin2 ). (16)

256

The rst term in the right-hand side is positive,

pk = E + pk cos = E pkm2 k 2E

2pk=

=12

(

k m2

)

=12(k kmin) > 0,

and |M|2 is therefore also positive.4.3. Dierential decay width. The dierential decay width is

d = (2)4(4)(q + k p) 12E|M|2 dQ,

where the phase space element is

dQ =d3q

(2)32E(q)d3k

(2)32(k).

We can use (3)(q +k p) to eliminate d3q. The remaining factor corresponds to the energy conservationlaw: E(q) = E . We next use a property of the Dirac delta function,

[f(x)] =

i

(x xi)|f (xi)| ,

where xi are the zeros of the function f(x). In our case, we consider E(q) a function of cos :

[E(q) + E] = (

E2 + k2 2pk cos + E) =

=

2pk2

E2 + k2 2pk cos

1

(

cos m2 k 2E2pk

)

.

Further, we write d3k = k2 dk d(cos ) d, integrate over the angle (which gives the factor 2), and usethe delta function to eliminate d(cos ). This xes the value of cos :

cos =m2 k 2E

2pk . (17)

Finally, the dierential decay width is

d =

2k

Ep(pk + p2 sin2 ) dk,

where = e2/4 and sin is given by (17). This decay width can be written more conveniently for futurecomputations:

ddk

=

81k

[A(k) + B(k)E1 + C(k)E2](

E2

m2 k

)

,

where is the step function,

A(k) = 4(k m2), B(k) = 4(m2 k),

C(k) = 2m2k2 + 2k3 m4 2k2 + 2m2k.

257

4.4. Eects on cosmic rays. We now compute the energy loss of protons in this background:

dE

dx=

dt

dx

dE

dt=

1v

(

d)

.

Using the previous results and v = p/E, we obtain the energy loss (with the integration limits given by (14)and (15))

dE

dx=

21p2

kmax

kmin

dk k

[

12(k m2) + p2(1 cos2 )

]

=

= 8p2

kmax

kmin

dk

[

2k2 (4m2 + 2m2 + 2)k + 2(2E2 + m2)

m2(4E2 + m2)1k

+ 4E

m2 k + k2 + 4Em2

m2 k + k2k

]

=

= 8p2

[

23

(

3E6

m6 m

6

3

)

12(4m2 + 2m2 +

2)(

2E4

m4 m

4

2

)

+

+ 2(2E2 + m2)(

E2

m2 m

2

)

m2(4E2 + m2) log(

2E2

m2m2

)

+ . . .]

.

The leading term isdE

dx=

8p222E4

m2=

2E2

4m2v2

2E2

4m2.

The energy as a function of the traveled distance is then

E(x) =E(0)

1 + (2/4m2)E(0)x.

The fractional energy loss for a cosmic ray with the initial energy E(0) traveling a distance x is

E(0) E(x)E(0)

=(2/4m2)E(0)x

1 + (2/4m2)E(0)x.

This loss becomes more important as the cosmic ray energy increases. But 2/4m2 is a very smallnumber. If we take E(0) = 1020 eV (the energy of the most energetic cosmic rays) and x = 1026 cm (aboutthe distance to Andromeda, the nearest galaxy, and therefore larger than the galactic halo), then the energyloss is smaller than 1 eV. For less energetic cosmic rays, the eect is even weaker.

As we have shown, the eect of the axion background on cosmic rays is quite negligible. Nevertheless,the emitted photons may be detectable. Using m = 1018 eV and = 1020 eV as characteristic valuesand having in mind the GZK cuto for protons (and a similar one for electrons1), we nd that the emittedphoton momenta are in the range 1016 eV < k < 100 eV for primary protons and 1016 eV < k < 400MeVfor primary electrons.

The number of cosmic rays with a given energy crossing a surface element per unit time is d3N =J(E) dE dS dt0, where J(E) is the cosmic ray ux. These cosmic rays radiate at a time t. The number ofphotons is given by

d5N = d3Nd(E, k)

dkdk dt = J(E)

d(E, k)dk

dE dk dt0 dS dt.

1It is very doubtful that electrons could be accelerated to such energies, but this is unimportant for our discussion becausethe intensity is extremely small at these energies.

258

Assuming that the cosmic ray ux is independent of time, we integrate this equation over t0 and obtaina factor t(E): the age of the average cosmic ray with the energy E. Because we do not care about theenergy of the primary cosmic ray (only that of the photon matters), we also integrate over E, starting fromEmin(k), the minimum energy that the cosmic ray can have in order to produce a photon with momentumk given by (14). We obtain the photon ux

d3Ndk dS dt

=

Emin(k)

dE t(E)J(E)d(E, k)

dk, Eth = 2

mm

. (18)

Further, we assume that t(E) is approximately constant and take t(E) Tp = 107 yr for protons andt(E) Te = 5 105 yr for electrons. We know that this last approximation is incorrect because t(E) 1/E,but we are now interested only in obtaining an order of magnitude estimate of the eect.

The photon energy ux is obtained by multiplying photon ux (18) by the energy of a photon withmomentum k:

I(k) = (k)

Emin(k)>Eth

dE t(E)J(E)ddk

T8k

Emin(k)

dE Ni[A(k)Ei + B(k)E(i+1) + C(k)E(i+2)],

where Emin(k) = m

k/ (see (14). A numerical analysis shows that the only relevant term in the decayrate is 4k from A(k). The integral can then be approximated by

I(k) T2

J [Emin(k)]Emin(k)min 1 k

(1)/2.

The value min should be taken from (2) or (3) depending on the range where Emin(k) falls. Substitutingthe numerical values, we obtain the approximate expressions for Ip(k) and Ie(k):

Ip(k) = 6(

Tp107 yr

)(

1020 eV

)1.84(k

107 eV

)0.84m2s1sr1,

Ie(k) = 200(

Te5 105 yr

)(

1020 eV

)2.02(k

107 eV

)1.02m2s1sr1.

As mentioned above, these expressions are only an estimate of the order of magnitude and assume constantaverage values for the age of a cosmic ray (either proton or electron). The interested reader can nd a moredetailed discussion in our recent paper [16], from which we take Fig. 1 describing the radiation yield.

5. Conclusions and outlook

We have investigated the eect on charged particles of a pseudoscalar background that is mildly timedependent (compared with the particle momentum). We considered both proton and electron cosmic rays.This eect is calculable because the axion background induces a modication of QED that turns out to beexactly solvable. This modication has several interesting features, such as the possibility of the photonemission process p p and e e (which we called the axion-induced bremsstrahlung processes). Weobtained kinematic constraints on the process; in particular, we showed that the process is possible onlyfor proton energies higher than a certain threshold. We computed the energy loss of protons in such abackground. For protons with energies below the GZK cuto, this loss is negligibly small.

259

I(k), m2s1sr1

EGZKEGZK

k, eV

Fig. 1. Radiation yield using the exact formulas and a more appropriate parameterization of the

electron cosmic ray average lifetime as a function of the energy (from [16]): We note that electrons in

general dominate the eect at low energies.

Nevertheless, the radiated photons could still be detected. We have computed their ux and energyspectrum in some detail. Because the energy threshold depends on the mass of the charged particle, it islower for lighter particles. Also, the energy loss is proportional to the inverse squared mass of the chargedparticle, and the eect is therefore more important for electrons. The value of kmin is independent ofthe charged particle mass, and the radiation spectra for electrons or protons dier very little (but theaverage lifetimes of electron and proton cosmic rays dier signicantly, and this has observable eects onthe radiation power spectrum).

The interested reader can refer to [16], where we describe this phenomenon in more detail and discussthe possibility of measuring this diuse radiation. Below, we summarize the main conclusions in that work.

The dominant contribution to the radiation via the considered mechanism is from electron (andpositron) cosmic rays. If we assume that the cosmic ray power spectrum is characterized by an expo-nent , then the produced radiation has an spectrum k(1)/2 for proton primaries, which becomes k/2

for electron primaries. The dependence on the key parameter /fa comes with the respective ex-ponents (1+)/2 and (2+)/2 for protons and electrons. But for the regions where the radiation yield islargest, the electron contribution is determinative. We assumed that the ux of electron cosmic rays isuniform throughout the Galaxy and hence identical to the ux observed in our neighborhood, but relax-ing this hypothesis could provide an enhancement of the eect by a relatively large factor. The eect forthe lowest wavelengths, for which the atmosphere is transparent, and for values of corresponding to thecurrent experimental limit is of the order O(101)mJy. This is at the sensitivity limit of antenna arraysthat are already being deployed, and it there seems possible to conduct experimental observations of thediscussed eect.

In the case of radiation originating from our galaxy, there is no answer to the main question relevant toour discussion: Is the ux of electron cosmic rays measured in our neighborhood representative of the wholegalaxy? Because it seems possible to relate this ux to the galactic synchrotron radiation, this radiationcould be deduced from measuring the ux. It appears that either the total number of electron cosmic raysis substantially larger than the number of rays measured in the Solar System or the galactic magnetic eldsmust be stronger than expected [17]. The quantitative analysis of this phenomenon requires further studies.There has also been no attempt to quantify the signal from possible extragalactic sources.

We note that the eect discussed here is a collective eect. This contrasts with the GZK eect discussedin Sec. 1: the CMB radiation is not coherent over large scales. For instance, no similar eect exists for

260

hot axions. One more observation is that some of the scales that play a role in the present discussion aresomewhat nonintuitive (e.g., the crossover scale m2/ or the threshold scale mm/). This is relatedto the fact that this eect is not Lorentz invariant. It may seem surprising at rst glance that an eectthat has such a low probability can give a small, but not negligibly small, contribution. The reason is thatthe number of cosmic rays is huge. It is known that their contribution to the galactic energy density iscomparable to the contribution of the galactic magnetic eld [18].

There are several aspects of the present analysis that could be improved to make it more precise. Inparticular, the problem with a piecewise-constant oscillating axion background or even with a backgroundwith a serrated time prole can be easily solved exactly without using special functions (the sine proleinvolves Mathieu functions). This problem will be considered in subsequent publications, but the presentanalysis already suces to obtain the order of magnitude of the eect. We hope that the consideredmechanism will help to understand the possible relevance of cold axions as a candidate for the role of darkmatter.

Acknowledgments. The authors thank their collaborator F. Mescia, who participated in the researchreported here. The authors also thank A. De Rujula for some comments on cosmic rays and especiallyJ. M. Paredes and P. Planesas for the discussions about the galactic synchrotron radiation. One of theauthors (D. E.) thanks the organizers of the XVI International Seminar on High Energy Physics Quarks-2010 for the warm hospitality.

This work was supported by the Ministerio de Educacion y Ciencia (Grant Nos. FPA-2007-66665 and2009SGR502), the Spanish Consolider-Ingenio 2010 program (Grant No. CPAN CSD2007-00042), and theMarie Curie Research Training Network FLAVIAnet.

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Axions and cosmic raysAbstract1. Axions2. Cosmic rays2.1. Cosmic ray energy spectrum2.2. The GZK cutoff

3. Solving QED in a cold axion background3.1. Euler-Lagrange equations3.2. Polarization vectors and dispersion relation

4. The process $p\longrightarrow p\gamma$4.1. Kinematic constraints4.2. Amplitude4.3. Differential decay width4.4. Effects on cosmic rays

5. Conclusions and outlookAcknowledgmentsReferences