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Observations of high energy primary electrons and their astrophysical significance
Adv. Space Res. Vol. 19, No. 5, pp. 767-770, 1997 0 1997 COSPAR. Published by Elsevier Science Ltd. All rights reserved Printed in Great Britain 0273-1177/97 $17.00 + 0.00 PII: SO273-1177(96)00144-5
Pergamon
OBSERVATIONS OF HIGH ENERGY PRIMARY ELECTRONS AND THEIR ASTROPHYSICAL SIGNIFICANCE J.
Nishimura”, T. Kobayashi**, Y. Komori*** and K. Yoshida*
* Faculty of Engineering, Kmagaw University, Yokolzama, Japan ** Aoyama Gakuin University, Tokyo *** Kanagawa Prefectural College, Yokohnnaa
ABSTRACT Cosmic-ray electrons lose their energy by synchrotron and inverse Compton processes during their propagation through the Galaxy. It has been recognized that the spectrum of electrons will bring us unique information about the propagation and acceleration of cosmic-rays. First, we review the past observations of high energy primary electrons beyond 100 GeV, and discuss the inherent technical difficulties in the direct observations of electrons in this high energy region. Next, discussion is made on the observations and their astrophysical significance to the electron spectrum in the TeV region. By these observations, it is promising to find the major contributing sources and to determine the parameters in models of propagation of cosmic-rays. Astrophysical importance is emphasized in the observations of electrons with more statistical significance in TeV region. Q 1997 COSPAR. Published by Elsevk Science Ltd. INTRODUCTION Many theoretical and experimental efforts have been made to study cosmic-ray electrons since their discovery in 196 1. Early attempts were made to see the relation of the Galactic radio waves and the confinement time of the cosmic-rays, since the electron loses its energy by the synchrotron and inverse Compton processes. Low energy electrons below 1 GeV are now also studied in relation to the recent gamma-ray observations. In this paper, we focus on the astrophysical significance of the electrons beyond the 1 TeV region. Those electrons lose their energy within a time of order of IO’ yrs, and cannot travel far distances from the sources. Therefore, the number of sources contributing to the electrons decreases progressively with electron energy, and one would expect large fluctuations in the spectrum at high energy end and also in the anisotropy of the electrons. We can discuss the spectral shape referring to the recent data of the supernova remnants and Pulsars. The spectral shape also depends on the propagation parameters, and the observations of high energy electrons are important to resolve the problem of the propagation models. One of the difficulties in studying such high energy electrons is the detection of those electrons. The flux is low, and we need a detector of large SR,with a high rejection power against hadron showers. Here, we first review the problem of the detection of high energy electrons, and show the reason why visual detectors such as Emulsion Chambers have been successfully used in high energy region. We also discuss the possibilities of other direct observations of high energy electrons in future. OBSERVATIONS
OF HIGH
ENERGY
ELECTRONS
We illustrate the observed spectrum of electrons in Figure 1. The flux of primary electrons is about 3kn’str.day’ beyond 1 TeV. The ratio of electrons to protons decreases progressively with energy, and becomes 10.’ at around 1 TeV. We need a detector of large 32 with high rejection power of at least IO” to observe electrons in this energy region. The highest energy electrons observed with electronic detectors are by the Chicago group (Tang et nl. 19841. Their detector is one of the most sophisticated, consisting of transition radiation and shower detectors. The transition radiation detector discriminates the electrons from the protons using Lorentz factor. Beyond a few hundred GeV, however, this rejection power decreases. Electronic detectors usually have a small solid angle. This also prevents the observation of high energy electrons due to the low flux by the counter type detectors. The pioneering work using emulsions to detect high energy electrons W;IS first made by the Tata group (Daniel & Stephens 1965). They used emulsion stacks, and detected the electrons up to several hundred GeV. They mentioned that the emulsion stacks have several merits over the electronic detectors, namely;
J. Nishimum et al.
168
The shower profiles at the starting point are I 1 llllrl, 1 I lll111, I a 111111, 1 “llm. ourdata accurately determined, and discriminate -7 _ tl Golden er al. 1994 > clearly between hadron and electron showers. 0 Tang1984 ! 10’ r The emulsion stacks have a large acceptance Golden et al. 1984 0 angle, giving a large SSZ even with a ti P relatively small detector. . ;,,,,,,(,,,,,,,,,,,, itjhpvp;r + “~web~re~a’.‘~io)1g80~ Energy determination is made by counting “E 2 the shower particles within a small radius from g 102= $ the shower axis. The total depth for full +$t tt+ 1 - 0 development of the shower and thus the g I -P weight of detector is reduced. t , , ,,,,,,, ! , ,,,,,,, ( , (,,,! The cost of emulsion stacks is high, and this ol prevents very large size detectors. 10’ ’ ’ “““’ IO4 10’ 10’ 10’ loo The emulsion chamber is essentially a stack of ENERGY f GeV ) lead and nuclear emulsion plates, and has the following merits over the emulsion stacks Fig. 1. Observed Spectrum of Cosmic Electrons. (Nishimura et al., 1980). The cost is much reduced, since the volume of emulsions occupies less than a few % of that of the detectors. The scanning problem is resolved by placing high sensitivity screen-type X-ray films (phosphoric plates with photographic films) on nuclear emulsion plates. We can detect the showers by naked eye scanning as a dark spot on the films, and locate the shower on the adjacent emulsion plates, and inspect by a microscope. Lead plates are used to develop the shower, and the weight is reduced to almost half of pure emulsion stacks. Identification of the electrons is also made by inspecting the starting point of the shower (see Figure 2). One of the difficulties of emulsion detectors is the accumulation of the background tracks, which limits the exposure duration. Also it has usually no timing information, and cannot be used to observe the anisotropy of the electrons. To resolve this problem, a new detector is developed with scintillating fibers by replacing the emulsion plates in the emulsion chamber. However, emulsion plates are placed at several layers to check the performance of these scintillating fibers for a while. The series of test flights were performed, and the results are promising. This can be used for observations of long exposures (Torii ef al., 1995). Other possible detectors for high energy electrons are also proposed. (Stephens etnl., 1983, Nishimura, 1994). They proposed to observe the synchrotron X-rays in the geomagnetic field or to utilize upgraded future Atmospheric Cherenkov telescopes having high rejection power for the detection of the gamma-ray sources. l
l
n
l
l l
l
l
ASTROPHYSICAL
SIGNIFICANCE
OF HIGH ENERGY ELECTRONS
The energy loss by synchrotron and inverse Compton process is proportional to the square of the electron energy, and the life time, T, of electrons of energy E is known as T=2.3.1 O5yr/E(TeV)=l/bE, where we assume < B’>=6.7pG, taking account of the depression of the cross-sections of the Compton effect at high energies (Van der Walt, 199 1). The density of electrons, N, at a distance R from a point source where electrons accelerated t yrs ago is: R’ (I- btE)Y-’ N= Exp(-E) EY 7 (4$:! where B is the time average of the diffusion constant of D for the particles of energy, and can be expressed as B=Do(E/GeV)s( 1-( 1-bEt)‘-‘/( (I-S)bEt ). We see from this formula, the average distance of R traveled by an electron during time t is: R-(4B t) I” Taking the case of E=lTeV, and assuming Do -1 OZR(E/GeV)’ cm’/sec, R is given by R-300pc for 6=0.3, R- 1kpc for 6=0.6. Then it is clear that only nearby sources contribute to the high energy electrons in the Solar system. Soectrum of Electrons from nearbv Sources
Assuming the electrons are accelerated at supernovae and the Pulsars, we can estimate the possible contributions of those nearby sources (Shen, 1970, Nishimura et (II., 1979, 1980, Aharonian et&., 1995). Supernova remnants and pulsars at a distance within 1kpc with ages less than lo” yrs are shown in Table 1. Loop 1 was once supposed to be the strongest candidate contributing to TeV region. However, its estimated age changed from -1O’yr to -10’ yr, and the contribution is much reduced. Furthermore there are some arguments Loop 1 could be a star burst near the Galactic Center. (Sofue, 1994). If this is the case, we cannot expect contributions from Loop 1 for high energy electrons. Recently the estimate of the age of Vela is also revised (Lyne et ~1. 1996), and the contribution of Vela becomes more dominant in TeV region than we had suspected at the time of Rome Conference (Nishimura et al.. 1995).
Observationsof High EnergyElzctrcm
Proion tan 8 = 0.43
Fe tan 0 = 0.60
#30184 Electron 100 GeV tan 9 = 0.60
Fig. 2. Starting
Profile of Electron
showers
Observed
in Emulsion
Chamber.
D=10”s(E/GeV)0.3
10’
10’
10’
10”
cm’/sec
Yt-
Fig. 3. Contributions to the high energy Electrons from each Source. The shaded region is corresponding with the location where the contributions of the supernovae are hard.
IO’
Fig. 4. Possible
contribution
of nearby
Sources
to the high energy
electron
spectrum.
J. Nishimuraet al.
770
Table. 1 List of Nearby Supernova SNR Pulsar SN 185 s 147 G65.3+5.7 Cygnus Loop Vela Monogem Loop 1 Geminga B105552
Distance 0.95 kpc 0.8 0.8 0.77 0.5 0.3 0.17 0.3 0.92
Remnants and Pulsars
Aqe l.8.1O’yr 4.6103 2.0104 2.010” 2-3’1 o4 8.6 10” 2.0.105 3.4.10s 5.3.10s
Emax 130 TeV 50 12 12 8-12 2.7 1.2 0.7 0.3
Ref.
(Lyne, et al., 1996) (Plucinsky, et al., 1996) (Eggar & Ashenbach, 1995)
We illustrate the degree of contribution of the nearby supernova remnants referring to the data shown in Table 1 (Figure 3). Only a few supernova remnants contribute to the high energy electron flux, and one would expect large fluctuations of the spectrum in this energy region. Contributions from each source also depend on the propagation parameters. We illustrate the change of the spectrum with diffusion constant, as shown in Figure 4. If we have enough statistical data in TeV region, we can judge which supernovae give major contributions. Moreover, we can estimate the probable value of the diffusion coefficient in this energy region. If the diffusion coefficient D is about lO’“cm’/sec around at TeV region, our observation is already at the high energy edge of the electron spectrum. Then it is possible to find the effect of the nearby sources by increasing the statistics of electron data in the few TeV region. Relating to the supernova origin of cosmic-rays, recent observations of X-rays reveal evidence of acceleration of electrons of a single power low spectrum up to 100 TeV in the shock of SN1006 (Koyama et&, 1995). The total amount of energy of electrons accelerated in SN1006 is estimated as -104”erg (Reynolds, 1995, Ozaki, 1996, Willingale et al., 1996), which is enough to explain the absolute flux of the observed electrons. AnisotroDv
of high enerev
electrons
Anisotropy of high energy electrons has been proposed as due to (Earl & Lenchek, 1969 Shen & Mao, 197 1). Aligned magnetic field line along the Galactic arm; Diffusion from nearby sources. As indicated by Ptuskin & Ormes (1995), and other papers, the anisotropy from a single source is given by
l l
,_3DVN_3R c N 2 ct’ Ptsukin & Ormes indicated that Vela is the most promising candidate to show a large anisotropy of 20%. This value could be reduced by a factor of 2-3, by taking the recent values of age. We would be able to observe this anisotropy, if the diffusion constant is large to give the enough flux from Vela as shown in Figure 4. References Aharonian, F.A., Atoyan, A.M., and Volk, H., A&A. 294 L41. (1995) Daniel, R.R. & S.A.Stephens, Phys. Rev. Letters, 5615,769. (1965) Earl, J.A.& Lenchek, ApJ.157,87 (1969) Eggar, R.J. & B. Aschenbach, A&A, 294 ( 1995) L25 Koyama, K., Petre, R., Gotthelf, E.V., Hwang, V., Matsuura, M., et d., Nuture, 378, L25S (1995) Lyne, A.G., Pritchard, R.S., Graham-Smith, F., and Camilo, F., Nuture 381,497 (1996) Nishimura, J. Proc. of .3rd Air Chrre~lkov Telescope Meeting, ICR. University of Tokyo 1 ( 1994) Nishimura, J., Fujii, M., and Taira, T., 16th ICRC(KJoto), 1, 488 (1979) Nishimura, J., Fujii, M., Taira, T., Aizu, E., Hiraiwa, H., et d., ApJ. 238, 394 (1980), Nishimura, J., Kobayashi, T., Komori, Y., and Tateyama, N., 24th /CRC (Rome) 3, 29 ( 1995) Ozaki, M. Private ConzrlzurlicLlti[>tl( 1996) Plucinsky, P.P., Snowden, S.L., Aschenbach, B., Egger, R., Edgar, R.J., et cd., ApJ. 463,224 (1996) Ptuskin V.S. and J.F. Ormes , Proc. of’24th ICRC(Romu), 3,56 (1995) Reynolds, S.P., ApJ. 450,13L ( 1996) Stephens S.A. and Balasubramanyan, V.K., J. Geophy. Res. 88,781 1 (1983) Shen, C.S. ApJ L.162, 181 ( 1960) Shen, C.S. & C.Y. Mao, A.stroph~~sical Letter 9, 169, (197 1) Sofue,Y., ApJ. 431 L9 1 .( 1994) Tang, K.K., ApJ. 278, 881 ( 1984), Torii, S., Nishimura, J., Kasahara. K., Tamura, T., Tateyama, N., etd., Proc. of24th ICRC(Roma), 3, 57.5 (1995) Van der Walt, D.J. Mon. Not. R. Astr. Sot. 251, 143 ( I 99 1) Willingale. R., West, R.G., Pye, J.P.. and Stewart, G.C., MOIZ. Not. R. Astr. Sot. 278,749 (1996)
Pergamon
OBSERVATIONS OF HIGH ENERGY PRIMARY ELECTRONS AND THEIR ASTROPHYSICAL SIGNIFICANCE J.
Nishimura”, T. Kobayashi**, Y. Komori*** and K. Yoshida*
* Faculty of Engineering, Kmagaw University, Yokolzama, Japan ** Aoyama Gakuin University, Tokyo *** Kanagawa Prefectural College, Yokohnnaa
ABSTRACT Cosmic-ray electrons lose their energy by synchrotron and inverse Compton processes during their propagation through the Galaxy. It has been recognized that the spectrum of electrons will bring us unique information about the propagation and acceleration of cosmic-rays. First, we review the past observations of high energy primary electrons beyond 100 GeV, and discuss the inherent technical difficulties in the direct observations of electrons in this high energy region. Next, discussion is made on the observations and their astrophysical significance to the electron spectrum in the TeV region. By these observations, it is promising to find the major contributing sources and to determine the parameters in models of propagation of cosmic-rays. Astrophysical importance is emphasized in the observations of electrons with more statistical significance in TeV region. Q 1997 COSPAR. Published by Elsevk Science Ltd. INTRODUCTION Many theoretical and experimental efforts have been made to study cosmic-ray electrons since their discovery in 196 1. Early attempts were made to see the relation of the Galactic radio waves and the confinement time of the cosmic-rays, since the electron loses its energy by the synchrotron and inverse Compton processes. Low energy electrons below 1 GeV are now also studied in relation to the recent gamma-ray observations. In this paper, we focus on the astrophysical significance of the electrons beyond the 1 TeV region. Those electrons lose their energy within a time of order of IO’ yrs, and cannot travel far distances from the sources. Therefore, the number of sources contributing to the electrons decreases progressively with electron energy, and one would expect large fluctuations in the spectrum at high energy end and also in the anisotropy of the electrons. We can discuss the spectral shape referring to the recent data of the supernova remnants and Pulsars. The spectral shape also depends on the propagation parameters, and the observations of high energy electrons are important to resolve the problem of the propagation models. One of the difficulties in studying such high energy electrons is the detection of those electrons. The flux is low, and we need a detector of large SR,with a high rejection power against hadron showers. Here, we first review the problem of the detection of high energy electrons, and show the reason why visual detectors such as Emulsion Chambers have been successfully used in high energy region. We also discuss the possibilities of other direct observations of high energy electrons in future. OBSERVATIONS
OF HIGH
ENERGY
ELECTRONS
We illustrate the observed spectrum of electrons in Figure 1. The flux of primary electrons is about 3kn’str.day’ beyond 1 TeV. The ratio of electrons to protons decreases progressively with energy, and becomes 10.’ at around 1 TeV. We need a detector of large 32 with high rejection power of at least IO” to observe electrons in this energy region. The highest energy electrons observed with electronic detectors are by the Chicago group (Tang et nl. 19841. Their detector is one of the most sophisticated, consisting of transition radiation and shower detectors. The transition radiation detector discriminates the electrons from the protons using Lorentz factor. Beyond a few hundred GeV, however, this rejection power decreases. Electronic detectors usually have a small solid angle. This also prevents the observation of high energy electrons due to the low flux by the counter type detectors. The pioneering work using emulsions to detect high energy electrons W;IS first made by the Tata group (Daniel & Stephens 1965). They used emulsion stacks, and detected the electrons up to several hundred GeV. They mentioned that the emulsion stacks have several merits over the electronic detectors, namely;
J. Nishimum et al.
168
The shower profiles at the starting point are I 1 llllrl, 1 I lll111, I a 111111, 1 “llm. ourdata accurately determined, and discriminate -7 _ tl Golden er al. 1994 > clearly between hadron and electron showers. 0 Tang1984 ! 10’ r The emulsion stacks have a large acceptance Golden et al. 1984 0 angle, giving a large SSZ even with a ti P relatively small detector. . ;,,,,,,(,,,,,,,,,,,, itjhpvp;r + “~web~re~a’.‘~io)1g80~ Energy determination is made by counting “E 2 the shower particles within a small radius from g 102= $ the shower axis. The total depth for full +$t tt+ 1 - 0 development of the shower and thus the g I -P weight of detector is reduced. t , , ,,,,,,, ! , ,,,,,,, ( , (,,,! The cost of emulsion stacks is high, and this ol prevents very large size detectors. 10’ ’ ’ “““’ IO4 10’ 10’ 10’ loo The emulsion chamber is essentially a stack of ENERGY f GeV ) lead and nuclear emulsion plates, and has the following merits over the emulsion stacks Fig. 1. Observed Spectrum of Cosmic Electrons. (Nishimura et al., 1980). The cost is much reduced, since the volume of emulsions occupies less than a few % of that of the detectors. The scanning problem is resolved by placing high sensitivity screen-type X-ray films (phosphoric plates with photographic films) on nuclear emulsion plates. We can detect the showers by naked eye scanning as a dark spot on the films, and locate the shower on the adjacent emulsion plates, and inspect by a microscope. Lead plates are used to develop the shower, and the weight is reduced to almost half of pure emulsion stacks. Identification of the electrons is also made by inspecting the starting point of the shower (see Figure 2). One of the difficulties of emulsion detectors is the accumulation of the background tracks, which limits the exposure duration. Also it has usually no timing information, and cannot be used to observe the anisotropy of the electrons. To resolve this problem, a new detector is developed with scintillating fibers by replacing the emulsion plates in the emulsion chamber. However, emulsion plates are placed at several layers to check the performance of these scintillating fibers for a while. The series of test flights were performed, and the results are promising. This can be used for observations of long exposures (Torii ef al., 1995). Other possible detectors for high energy electrons are also proposed. (Stephens etnl., 1983, Nishimura, 1994). They proposed to observe the synchrotron X-rays in the geomagnetic field or to utilize upgraded future Atmospheric Cherenkov telescopes having high rejection power for the detection of the gamma-ray sources. l
l
n
l
l l
l
l
ASTROPHYSICAL
SIGNIFICANCE
OF HIGH ENERGY ELECTRONS
The energy loss by synchrotron and inverse Compton process is proportional to the square of the electron energy, and the life time, T, of electrons of energy E is known as T=2.3.1 O5yr/E(TeV)=l/bE, where we assume < B’>=6.7pG, taking account of the depression of the cross-sections of the Compton effect at high energies (Van der Walt, 199 1). The density of electrons, N, at a distance R from a point source where electrons accelerated t yrs ago is: R’ (I- btE)Y-’ N= Exp(-E) EY 7 (4$:! where B is the time average of the diffusion constant of D for the particles of energy, and can be expressed as B=Do(E/GeV)s( 1-( 1-bEt)‘-‘/( (I-S)bEt ). We see from this formula, the average distance of R traveled by an electron during time t is: R-(4B t) I” Taking the case of E=lTeV, and assuming Do -1 OZR(E/GeV)’ cm’/sec, R is given by R-300pc for 6=0.3, R- 1kpc for 6=0.6. Then it is clear that only nearby sources contribute to the high energy electrons in the Solar system. Soectrum of Electrons from nearbv Sources
Assuming the electrons are accelerated at supernovae and the Pulsars, we can estimate the possible contributions of those nearby sources (Shen, 1970, Nishimura et (II., 1979, 1980, Aharonian et&., 1995). Supernova remnants and pulsars at a distance within 1kpc with ages less than lo” yrs are shown in Table 1. Loop 1 was once supposed to be the strongest candidate contributing to TeV region. However, its estimated age changed from -1O’yr to -10’ yr, and the contribution is much reduced. Furthermore there are some arguments Loop 1 could be a star burst near the Galactic Center. (Sofue, 1994). If this is the case, we cannot expect contributions from Loop 1 for high energy electrons. Recently the estimate of the age of Vela is also revised (Lyne et ~1. 1996), and the contribution of Vela becomes more dominant in TeV region than we had suspected at the time of Rome Conference (Nishimura et al.. 1995).
Observationsof High EnergyElzctrcm
Proion tan 8 = 0.43
Fe tan 0 = 0.60
#30184 Electron 100 GeV tan 9 = 0.60
Fig. 2. Starting
Profile of Electron
showers
Observed
in Emulsion
Chamber.
D=10”s(E/GeV)0.3
10’
10’
10’
10”
cm’/sec
Yt-
Fig. 3. Contributions to the high energy Electrons from each Source. The shaded region is corresponding with the location where the contributions of the supernovae are hard.
IO’
Fig. 4. Possible
contribution
of nearby
Sources
to the high energy
electron
spectrum.
J. Nishimuraet al.
770
Table. 1 List of Nearby Supernova SNR Pulsar SN 185 s 147 G65.3+5.7 Cygnus Loop Vela Monogem Loop 1 Geminga B105552
Distance 0.95 kpc 0.8 0.8 0.77 0.5 0.3 0.17 0.3 0.92
Remnants and Pulsars
Aqe l.8.1O’yr 4.6103 2.0104 2.010” 2-3’1 o4 8.6 10” 2.0.105 3.4.10s 5.3.10s
Emax 130 TeV 50 12 12 8-12 2.7 1.2 0.7 0.3
Ref.
(Lyne, et al., 1996) (Plucinsky, et al., 1996) (Eggar & Ashenbach, 1995)
We illustrate the degree of contribution of the nearby supernova remnants referring to the data shown in Table 1 (Figure 3). Only a few supernova remnants contribute to the high energy electron flux, and one would expect large fluctuations of the spectrum in this energy region. Contributions from each source also depend on the propagation parameters. We illustrate the change of the spectrum with diffusion constant, as shown in Figure 4. If we have enough statistical data in TeV region, we can judge which supernovae give major contributions. Moreover, we can estimate the probable value of the diffusion coefficient in this energy region. If the diffusion coefficient D is about lO’“cm’/sec around at TeV region, our observation is already at the high energy edge of the electron spectrum. Then it is possible to find the effect of the nearby sources by increasing the statistics of electron data in the few TeV region. Relating to the supernova origin of cosmic-rays, recent observations of X-rays reveal evidence of acceleration of electrons of a single power low spectrum up to 100 TeV in the shock of SN1006 (Koyama et&, 1995). The total amount of energy of electrons accelerated in SN1006 is estimated as -104”erg (Reynolds, 1995, Ozaki, 1996, Willingale et al., 1996), which is enough to explain the absolute flux of the observed electrons. AnisotroDv
of high enerev
electrons
Anisotropy of high energy electrons has been proposed as due to (Earl & Lenchek, 1969 Shen & Mao, 197 1). Aligned magnetic field line along the Galactic arm; Diffusion from nearby sources. As indicated by Ptuskin & Ormes (1995), and other papers, the anisotropy from a single source is given by
l l
,_3DVN_3R c N 2 ct’ Ptsukin & Ormes indicated that Vela is the most promising candidate to show a large anisotropy of 20%. This value could be reduced by a factor of 2-3, by taking the recent values of age. We would be able to observe this anisotropy, if the diffusion constant is large to give the enough flux from Vela as shown in Figure 4. References Aharonian, F.A., Atoyan, A.M., and Volk, H., A&A. 294 L41. (1995) Daniel, R.R. & S.A.Stephens, Phys. Rev. Letters, 5615,769. (1965) Earl, J.A.& Lenchek, ApJ.157,87 (1969) Eggar, R.J. & B. Aschenbach, A&A, 294 ( 1995) L25 Koyama, K., Petre, R., Gotthelf, E.V., Hwang, V., Matsuura, M., et d., Nuture, 378, L25S (1995) Lyne, A.G., Pritchard, R.S., Graham-Smith, F., and Camilo, F., Nuture 381,497 (1996) Nishimura, J. Proc. of .3rd Air Chrre~lkov Telescope Meeting, ICR. University of Tokyo 1 ( 1994) Nishimura, J., Fujii, M., and Taira, T., 16th ICRC(KJoto), 1, 488 (1979) Nishimura, J., Fujii, M., Taira, T., Aizu, E., Hiraiwa, H., et d., ApJ. 238, 394 (1980), Nishimura, J., Kobayashi, T., Komori, Y., and Tateyama, N., 24th /CRC (Rome) 3, 29 ( 1995) Ozaki, M. Private ConzrlzurlicLlti[>tl( 1996) Plucinsky, P.P., Snowden, S.L., Aschenbach, B., Egger, R., Edgar, R.J., et cd., ApJ. 463,224 (1996) Ptuskin V.S. and J.F. Ormes , Proc. of’24th ICRC(Romu), 3,56 (1995) Reynolds, S.P., ApJ. 450,13L ( 1996) Stephens S.A. and Balasubramanyan, V.K., J. Geophy. Res. 88,781 1 (1983) Shen, C.S. ApJ L.162, 181 ( 1960) Shen, C.S. & C.Y. Mao, A.stroph~~sical Letter 9, 169, (197 1) Sofue,Y., ApJ. 431 L9 1 .( 1994) Tang, K.K., ApJ. 278, 881 ( 1984), Torii, S., Nishimura, J., Kasahara. K., Tamura, T., Tateyama, N., etd., Proc. of24th ICRC(Roma), 3, 57.5 (1995) Van der Walt, D.J. Mon. Not. R. Astr. Sot. 251, 143 ( I 99 1) Willingale. R., West, R.G., Pye, J.P.. and Stewart, G.C., MOIZ. Not. R. Astr. Sot. 278,749 (1996)