f' 15» UterMtliiil Cismlc Ray Cinfiniet
CONFERENCE PAPERS VOLUME 1 OG SESSION
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PLOVDIV, BULGARIA AUGUST13-B8,1977 1Stt International Cosmic Ray Conference
CONFERENCE PAPERS VOLUME 1 OG SESSION
BULGARIAN ACADEMY OF SCIENCES PLOVDIV, BULGARIA AUGUST 13-26,1S77 Ill
PREFACE
The present publication contains the proceedings of the 15th International Cosmic Ray Confe rence, Plovdiv, 13-26 August. 1977. This Conference is to be held under the auspices of the Inter national Union of Pure and Applied Physics, organized by the Bulgarian Academy of Sciences The publication comprises 12 volumes. Volumes from 1 to 9 include the original contribu tions, which have arrived at the Secretariat of the National Organizing Committee by May 26, 1977. Papers which have been declared but not submitted by that date have been represented by their abstracts. Volumes from 10 to 12 include the invited and rapporteur lectures, as well as late origi nal papers. Volume 12 contains the general contents of the volumes, an authors' index and other references. All papers included in the present publication are exact reproductions of the authors' original manuscripts. The Secretariat has not made any corrections or changes in the texts. The original contributions have been accepted and included in the programme after a decision of the Interna tional Programme Advisory Board of the 15th ICRC on the basis of their abstracts. The full texts of the papers, however, have not been refereed by the editorial board of the present publication. The first nine volumes have been organized in accordance with the classical headings adopted at the cosmic ray conferences, which also coincide with the sessions. Volume 1 - OG (Origin) Session Volume 2 - OG (Origin) Session Volume 3 • MG (Modulations and Geophysical Effects) Session Volume 4 - MG (Modulations and Geophysical Effects) Session Volume 5 - SP (Solar Particles) Session Volume 6 - MN (Muons and Nutrinos) Session Volume 7 - HE (High Energy Physics) Session Volume 8 - EA (Extensive Air Showers) Session Volume 9 - T (Techniques) Session The National Organizing Committee is indebted to the invited reporters and rapporteur lec turers, as well as to all authors of original papers, who, by their hard and highly qualified work, have contributed to the success of the Conference and have made possible the publication of the present proceedings. We also express acknowledgement to the members of the Organizing Committee and the Se cretariat of the Conference, as well as to the Publishing House of the Bulgarian Academy of Sciences, without whose diligent work the publication of the proceedings would have been im possible.
Acad. Christo Ya. Christov Chairman of the National Organizing Committee IV BULGARIAN NATIONAL ORGANIZING COMMITTfcfc
Honorary Chairman - Acad. A. Balevsky, President of the Bulgarian Academy of Science* and Member of the State Council
Executive Chairman: Ch. Ya. Christov Vice-Chairman: P.N. Markov Secretary: B.L. Betev Members: M. Borisov, I. Todorov, G. Nestorov, K. Serafimov, Ts. Bonchev, Ts. Petkov, D. Pari- kian, N. Balabanov, J. Stainenov, L. Popova, St. Kavlakov, T. Stanev, N. Ahababun, S. Ushev, Ch. Tcherne", T. Palev, I. Kirov, J. Georgiev, L. Katsarsky
MEMBERS OF THE COSMIC RAY COMMISSION OF IUPAP Chairman: Professor A.J. Somogyi (Hungary) Secretary: Profeisor S. Miyake (Japan) Members: Professor A.E. Chudakov (USSR), Professor R.R. Daniel (India), Professor R. Gail (Mexico), Professor B. Peters (Denmark), Professor K. Pinkau (FRG), Professor H. Reeves (France), Professor C.J.Waddington (USA), Professor A.W. Wolfendale (UK)
INTERNATIONAL PROGRAMME ADVISORY BOARD Chairman: Professor Ch. Christov Secretary: Dr B. Betev Members: Professor A. Chudakov (USSR), Professor H. Elliot (UK), Professor S. Miyake (Japan), Professor S. Nikolsky (USSR), Professor K. Pinkau (FRG), Proftssor A. Somogyi (Hungary), Professor C. Waddington (USA), Professor G. Yodh (USA)
The 15th International Cosmic Ray Conference is organized by the Bulgarian Academy of Sciences under the auspices of the International Union of Pure and Applied Physics.
ADDRESS OF SECRETARIAT Institute for Nuclear Research and Nuclear Energy Sofia 1113, 72 Blvd Lenin Telephone: 7341 Telex: SOFIA BAN 22424 V ТЛ1Н.К ОГ CONTENTS VOLl'ME 1 - ORIGIN
DIFFUSE COSMIC AND GALACTIC GAMMA-BAYS page
OG-1 COS-Б Observations of Galactic Gamma-itay Emission * (Large Scale) J.A. Paul, K. Bennett, G.F. Bignami, R. Buccheri, W. Hermsen, G. Kanbach, F. Lebrun, H.A. Mayer- Hasselwander, G. Picclnotti, L. Scars.i, F. Soroka, B.U. Svanenburg and R.D. Wills
CG-2 Observations of Gamma Radiation from the Galactic Center Region ' J. Samimi, A. Jabbari-Azad, B. L. Kinzer and G. H. Share (Abstract)
OG-3 Gamma Bay Spectrum from the Galactic Disk ° R.K.Sood, G.K.Rochester, P. G. Clayton and B.T.Thomas
OG-4 Observation of Very High Energy Gamma Radiation in Direction of Galactic Plane 12 V. P. Fomin, B. M. Vladimlrsky and A, A. Stepanian
OG-5 Gamma Radiation from Central Region of the Galaxy and Diffuse Background at Ef>100 MeV *7 A. I. Beljaovsky, V. L. Bokov, V. K. Bocharkin, I. F. Bugakov, G. H. Gorodingksy, E,M. Kruglov, E. V.Mjakinin, G. A. Pjatigorsky, E. LChujkln, V. S.Juferev, .J. A. Shlvanov, E.N.Kolesnikova and B. A. Beloborodko
OG-6 Search for Primary Gamma Bays in the Medium Energy Range 5 MeV - 50 MeV 23 E. Jonfi».:d, B. Parlier, B. Mougin, J. Andrejol, J. C. Courtois, A. Lecomte, M. Gorisse, B. Agrinier, J. M. Lavlgne, C. Doulade, M. Nlel, R. Palmeira and K.R.Rao (Abstract)
GG-7 Gamma-Rays of 3 to 25 MeV from the Galactic Anticenter and Pulsar NP 0532 24 R.B.Wilson, S.Moon, J.Ryan, A.D.Zych, P.S.White and B. Dayton
OG-8 Comparison of High Energy Gamma Rays from (b)>30 with the Galactic Neutral Hydrogen Distribution 30 ill. E.Ozel, H.Ogelman, T.Turner, C. E. FIchtel, R.C.Hartman, D. A. Kniffen and D. J. Thompson VI
PW
OG-9 Electromagnetic Gamma Ray Spectrum from the Galactic Disk R. Schllckelser and K.O.Thlelhelm
OG-10 A Teat for the Origin of Galactic Gamma Rays •13 R. Schlicketser and K.O.Thlelhelm
OG-11 Galactic Rays and the Origin of Cosmic Rays 19 D.Bodds, A.W.Wolfendale and V..C.I-'-.'foung
OG-12 Gamma Ray Production in Dense Molecular Clouds 55 A. W. Strong and J. Skilling
OG-13 Cosmic Ray Penetration into Molecular Clouds V G1 C. J. Cesarsky and H.J. Votk
CG-14 The Spectrum of Compton-Scattered Galactic Gamma Rays 66 R. Cowslk and W. Voges (Abstract)
CG-15 Bremsstrahlung Gamma Radiation from the Galaxy 67 P. G. Shukla and C. J. Gesarsky
CG-16 Atmospheric Gamma-Ray Angle and Energy Distributions from 2 to 25 MeV 73 J.MRyan, S.H.Moon, R.Wilson, A.D.Zych, R.S.White and B. Dayton
OG-213 iiocket Altitude Atmospheric Х-Hays from Magneto spheric Electrons at Subauroral Latitudes 79 J. G. Luhmann and J. B. Blake
OG-17 Observational Constraints on the Possible Existence of Cosmo- logical Cosmic Rays 85 T. Montmerle
OG-18 The Contribution of Discrete Sources to the Flux of Extragalactic Gamma Rays 91 A.W. Strong, A.W.Wolfendale and D.K.Worrall
CG-19 The Gamroa^Ray Tnminosity of Spiral Galaxies, Its Evolution and Its Contribution to the Diffuse Background Above 100 MeV 95 G. G. Lichti, G. F. Btgnami and J. A. Paul
OG-20 Cosmic Diffuse Gamma-Rays at Medium Energies 100 R.S.White, 5. H.PIoon, J.M.Ryan, Н.Л.Wilson, A.D.Zych and B. Dayton VII GAMMA-RAV SCUHCES
OG-21 COS-B Observations of Gamma Ray Emission from Pulsars IOC The Caravane Collaboration (Abstract)
OG-22 COS-B Observations of Discrete Sources of High-Ewe rgy Gamm» Radiation 107 The Caravane Collaboration (Abstract) OG-23 Electromagnetic Processes In Pulsars Under Strong Electric and Magnetic Field Conditions 108 S. Ayasli, A.Hacinliyan, H. B. Ogelman and J. K. Daugherty
OC-24 Pulsed High Energy Gamma Rays from Pulsars 114 S.K.Gupta, P.V.Ramana Murthy, B. V. Sreekantan and S. C.Tonwar (Abstract)
OG-25 Contribution of Pulsars to the Flux of Galactic Gamma Rays 115 A. W. Strong and A. W. Wol'endale
OG-26 Gamma Rays from the Directionof the Galactic Centre 119 A. W. Wolfendale and D. M. Worrall
OG-27 Supernova Shell Around Young Pulsars as - Gamma 122 Neutrino and Neutron Source V. S. Berezinsky and C. F. Prilutsky
OG-28 A Possible Contribution of Discrete Sources to the Diffuse Gamma Ray Distribution Over Galactic Longitude 127 A. A. Stepanlan OG-29 Periodical Gamma-Radiation from the Х-Ray Source L'YG -C-3 131 A. M. Galper, V. G. Kirlllov-Ugryuniov, A. V.Kuroehkln, N. G.Lelkov, B. LLuchkov, Yu.T.Yurkin, V. P. Fomin, Yu. I. Neshpor, A. A. Stepanjan and B. M. Vladtmlraky
OG-30 Gamma-Ray Emission Spectrum of the CYG X-3 Source and Its Possible Origin 135 A. A. StqJanlan, B.M. Vladimir sky, Yu. I. Neshpor and V. P. Fomin
OG-31 Gamma Rays from CYG X-3 and CYG X-l 141 Krishna M. V. Apparao (Abstract) 11 12 OG-32 The Mount Hopklnfi Sky Survey of Sources of 10 to 10 eV 142 e Gamma Raye T.C.Weekes, H.F.Helmken, J.E. Grlndlay and E.Horlne (Abstract) OG-33 Coamlc Tamroa Ray Observations In. the 0.1 - 4 MeV Energy Range 143 P. Mandrou, M. Nlel, A.Dupont and G. Vedrenne (Abstract) VIM
X-RAY ASTRONOMY
OG-34 Diffusive Soft Х-Ray Background from the Interaction of Magnetosphertc Electrons with the Upper Atmosphere 144 J. G. Luhmann and J. B. Blake (Abstract
OG-35 Artel 5 Measurements of Sources with Power Law Spectra and Faet Time Variations 145 M. J. Coe, A. R. Engel and J. J. Quenby
OG-36 Ariel V Measurements of Sources with Power Law Spectra an1 Faet Time Variations 151 M.J.Coe, A.R. Engel and J.J.Quenby (Abstract)
OG-37 Ariel 5 Observations of Extra-Galactic Х-Hay Sources and Implications of Hard Х-Ray Galactic Source Measurements for the.Galactic Disc Gamma Ray Emission 152 C. S. Dyer, M. J. Coe, A. R. Engel and J. J. Quenby
OG-38 Recent Ariel V Measurements on Hard Х-Ray Bursts and the Implications for J-Burst Origin 157 J. J. Quenbv. M. J. Coe and A. R. Engel (Abstract)
OG-39 Observation of the Hard Х-Radiation from the Galaxy 158 R.L.Aptekar, S. V. Golenetskll, Yu. A. Guriyan, V. N. Il'lnskll, E. P. Mazets, V. N. Panov and I. A. Sokolov (Abstract)
OG-40 Hard X Ray Observations of 3U1820-30 159 K.Hurley (Abstrao*^
OG-41 Observations of Cosmic Х-Ray Sources Above 20 KeV from OSO-8 160 B.R.Dennis, 0.J.Crannell^J.F.Doiau, K.J.Frost, L.E.Orwig, J. H. Beall and G. F. Maurer (Abstract)
оЪ-42 Time Variation of Hard Х-Ray from CYG X-l 161 N. Nakagawa, H. Sakural and M. Uchlda
OG-4£ Rocket Observations of Celestial Х-Ray Sources at Soft X-Ray Energies 166 M.S.Radha, D.P.Sherma, T.M.K.Marar, V.S.Iyengar, K.Kastu Ix-angan and U.R.Rao (Abstract)
OG-210 Preliminary Results from the Stgne 3 Satellite 167 M.Nlel, A.R. Bazer-Bachl and G. Vedrenne (Abstract)
OG-44 Some Aspects Concerning Accretion In Binary Systems je8 G.M. Filho (Abstract) IX
GAMMA-RAY BURSTS
OG-45 Cosmologtcal Cosmic Rays and Nuclear Gamma Ray Emission }gg T. Montmerle
OG-46 Spectrum, Time Structure and Direction of Incidence of August 16, 1976 Gamma Ray Buret 173 M.Sommer, D.Muller, H.Ilorstman and L.Bassani
OG-47 Observation of Gamma Kay Burst at Balloon Altitude 179 J. Nlshtmura, M.Oda, S. Miyamoto, Y.Ogawara, M. Fujil, T.Yamagami, T.Tawara, M.Yoshlmori, H.I'urakami, K.Nakagawa and
OG-48 High Energy Gamma Ray Bursts 180 S.K.Gupta, P.V.Ramana Murthy, B. V. Sreekantan and S. C. Tonwar (Abstract)
OG-49 A Search for Gamma Ray Bursts from the Explosive Evaporation of Black Holes 181 T. C.Weekes and N. A. Porter
OG-50 Search of the Cosmic Gamma Radiation Bursts with the Energy Greater Than 100 KeV on the Date from the Satellite 'Cosmos-731' 187 A. I. Beljaevsky, V. L. Bokov, V.K. Bocharkin, I. F.Bugakov, G. M. Gorodlnsky, B.A.Dmltrlev, E.M. Kruglov, E. V. Mjakinin, G. A.Pjatlgorsky, E.I.Chujkln, V. S.Juferev, J. A.Shlbanov, E. N. Kolesnlkova and B. A. Beloborodko (Abstract) 13 14 OG-51 A Search for Bursts of 10 to 10 eV Gamma Rays Using Spaced Cosmic Ray Stations 188 G.O'Sulllvan, D.J.Fegan, B.Mc Breen and D.O'Brien OG-52 -Search for Cosmic Muon Bursts from Relativistic Dust- Grains 194 H. Muramoto, S. Yamamoto, T. Takahashi, Y. Teramoto, S. Htgasht and S. Ozakl
OG-217 Gamma Ray Production In Interstellar Space 158 G. D. Badhwar and S. A. Stephens
NUCLEAR COMPOSITION OF COSMIC RAYS
OG-53 Balloon Investigation of Spectra and Composition of Singly Charged Component of the Cosmic Radiation with a Magnetic Spectrometer 203 E. A. Bogpmolov, N. D. Lubyanaya, V. A.Romanov and S. V. Stepanov (Abstract) X
OG-218 fh# Proton *na tielLum Rigidity Spectra from 10 to 50 GV 201 G.D. eadhwar, R.R.Daniel, T. Clsghorn, R. L. Golden, J. L. Lacy, S. A. Stephesn and J. E. Zlpse
OG-219 The Cosmic Ray Antiproton Flux: an Upper Limit at the Predicted Level of Observation 209 CD.Badhwar, B.R.Daniel, T.Cleghorn, B.L. Golden, J. L. Lacy, S. A. Stephens and J. E. Zlpse 12 14 OG-54 Estimation of Primary Proton Spectrum Between 10 and 10 eV 214 T.K. Galsser, F. Slohan and G, B. Yodh (Abstract) OG-55 On the Proton Spectrum Measurements Above 2000 GeV 215 G.B.Yodh, R.W. Ellsworth, A. tto, J.MacFall, F. Slohan, R. E. Streltmatter, S. C. Tonwar and V. K. Balasubrahmanyan (Abstract)
OG-56 The Iron /Proton Primary Ratio and Relative Cosmic Ray Composition at 1012 - 1013 eV 216 J. E. Grlndlay, D. Hartman and T, C.Weekes (Abstract)
OG-57 Atmospheric Corrections for Charge Spectra and Isotope Ratios of Heavy Nuclei Measured at Balloon Altitudes 217 W.Heinrlch (Abstract)
OG-59 Measurements of the C/O and B/N Abundance Ratios in the Primary Cosmic Radiation in the Energy Range 0.5-2.0 GeV/ Nucleon Made During September 1976 21P J, H, Derrickson, T. A, Parnell and J. C, Gregory (Abstract)
OG-6P Measurement of the Li/B Abundance Ratios above 650 MeV/Hucleon 219 B. B'yrnak, N. Lund, I. Lundgaard Rasmussen and M.Rotenberg
OG-61 The Relative Abundances of the Elements Silicon Through Nickel in the Low Energy Galactic Cosmic Rays 224 M. Garcta-Munoz, G. M. Mason and J. A. Simpson
OG-6Z Charge and Energy Spectra of heavy Cosmic Rays at Inter mediate Energies 230 M. Garcta-Munoz, G. M.Mason, J. A. Simpson and J. P.Wefel
OG-63 Nuclear Composition of the Cosmic Rays Between 0. S"and~ 10 GeV/Nucleon 23G B.Byrnak, N. Lund, I.L.Rasmusien ana M.Rotenberg (Abstract) XI
OG-65 The Implication for Galactic Propagation.from a Measurement 237 of the Charge (Z=5-26and Energy (0. 5-5 BeV/nuc) Spectra of Cosmic Rays J. A. Leznlak, W.R.Webber, J. C. Kiah and G. A. Simpson
OG-66 Cosmic Ray Composition between 5 and 300 GeV/Nucleon 242 J. F. Arens, V. K. Balasubrhamanyan, J. F. Ormee, W. K.H.Schmidt, M.Simon, F. Slohan, H. Spiegelhauer and G. B. Yodh (Abstract)
OG-67 Charge Composition and Energy Spectra of Cosmic-Ray Nuclei at Energies above ~ GeV per Nncleon 243 John H. Caldwell and P. Meyer
OG-68 Chemical Composition of Cosmic Ray Iron Group Nuclei at About 1 GeV/nuc. 248 P. Meyer and G. Mlnagawa
OG-70 Cosmic Ray Abundances from Nitrogen to Zinc Using a Cellulose Nitrate Plastic Detector 253 G. S.Kalnth, V.S.Bhat«a and S. Paruthl (Abstract)
OG-71 VH Cosmic Ray Measurements with a 6.6 m Sr Electronic Detector 254 J.Tueller, P.Love, J.W.Epstein, M. H. Israel and J. Klarmann
OG-72 Ultra Heavy Cosmic Ray Measurements with a 6.6 m Sr Electronic Detector 258 P.Love, J.Tueller, J.W.Epstein, M.H.Israel and J. Klarmann
OG-73 Average Abundances of Galactic Cosmic Rays with Z 50 from Studies of Meteoritlc Olivines 262 V. P. Perelygin, S. G. Stetsenko, D. Lhagyasuren, O.Otgon- suren, P. Pellas and B. Jakupl
OG-74 UH Cosmic Rays: Possible Origin in Massive Stars 268 J.P.Wefel D.N.Schramm and J.S.Blake
OG-75 On the Dependence of r-Process Yields Upon Experimentally Unknown Nuclear Parameters Ultra-Heavy Cosmic Hay- Abundances 274 J. B. Blake and D. N. Schramm XII
OG-76 The r-Process Path In the Superheavy Region ?N> K. L.Haln*b*ch, J. B. Blake and D. N. Schramm
OG-77 Theoretical Calculation» of the Fe and Co Abundance* In Galactic Cosmic Ray* гм R.LKucnetaova. A. V.FIaenko and A.K.Lavrukhina
OG-78 Monopole, Antinucteua or Superheavy Nucleua ? 290 W. Z.Oaborw, L.S.Plnaky, P. B.Price, Е.К. Shirk ard R.Hagatrom (Abstract) 2 OG-79 The Hadron Energy Spectrum at a 10 G/cm Depth In the Stratosphere 29 i V.V.Abulova, S. B. Ignatyjv, M. A. Ivanova, K. V.Mandritskaya, I.V.Rakobolskaya, G. P. Sazhlna, N. V.Sokolskaya, N.LTull- nova, A. Ya.Varkovltskaya, E. A. Zamchalova, G.T. Zatseptn and V. L Zatsepin (Abstract)
ISOTOPIC COMPOSITION OF COSMIC RAYS
OG-80 Measurements of Helium Isotopic Composition In Primary 292 M. G. Jodko, V. K. Karakadko and V. A. Romanov
OG-81 The Isotopic Compositbn of LI-Be and В 295 F, В. McDonald, J.H.Trainon and W.R.Webber (Abstract) 10 OG-82 Cosmic Ray Age and a Measurement of the Abundance of Be 296 Using a Magnetic Spectrometer C.D.Orth, A.Suffington, P.M.Lubin, T.S.Mast and G.F.Smoot OG-83 The Isotopic Composition of Galactic Cosmic Ray Lithium, Beryllium and Boron 301 M. Garcia-Munoz, G. M. Mason and J. A. Simpson
OG-84 The Isotoplc Composition of Galactic Coamic-Ray Beryllium and the Cosmic Ray Age 307 M. Garcla-Munoz, G. M. Mason and J. A. Simpson
OG-86 Measurement of the Isotopic Composition of the Primary Cosmic 313 Radiation for the Elements B-Ne C.Bjarle, N.Y.Herrstrom, L.Jacobsson, CJonsson and K.Kristiansson OG-87 Mean Mass of Cosmic Ray Ne, Mg, SI at 1,2 GeV/amu 318 R. Dwyer and P. Meyer
OG-88 A Study of Cosmic Ray Charge and Isotopic Composition Using a Charged Particle Telescope 324 S. Dev Verma X!l!
oG-89 Isotoplc Composition of Low Energy Cosmic Ray Particles with Charges Z=5 - 8 329 R. Beauiean, H. Sagebiel and W. Enge
uG-90 Fe - Isotopes In Cosmic Rays 334 G. Siegmon, K. P. Bartholoma and W. Enge
OG-91 Cosmic Ray Iron Measurements with Fragementatlon Controlled 340 R.Scherrsr, W. Enge, R.Beaujean. S. Hertzman, K.Kristiansson andK. Soderstrom (Atalract)
OG-92 3tudy of the Be/B Ratio as a Function of Energy and Иа Implications for the Age of Cosmic Rpys 341 W.R.Webber, J.A.Lezniak, J.C.Kishsnc C.A.Simpson
OG-93 Isotoplc Composition of Low-Energy Cosmic-Ray Nuclei with Z=10-14 347 G. A. Simpson, J. Barbary, J. Kish, J. A. Lezniak and W.R.Webber (Abstract)
OG-94 A Measurement of the Isotopto Composition of Cosmic Ray Fe and Other Nuclei with Z=20-25 348 G.A.Simpson, J.KIsh, J. A. Lezniak and W.R.Webber
OG-95 A Detector of Cosmic Ray Isotopes 354 P. S. Freler, C. M. Gilman, W. R. Scarlett and C. J.Waddington (Abstract)
OG-96 Isotoplc Composition of Iron Group Nuclei at 320-500 MeV/amu 355 S.P.Ahlen, B. G.Cartwrlght, P. B. Price and G. Tarte (Abstract) OG-221 A Study of the Isotopes of Cosmic Ray Fe Nuclei In Lexan Polycarbonate Exposed on Skylab 356 P.B.Price, E.K.Shtrk, C.O'Ceallalgh, D.O'Sulllvan and A. Thompson (Abstract)
ELECTRONS AND POSITRONS
OG-97 The Stjectrum of Cosmic Electrons with Energies between 6 and 100 GeV 357 C. A. Meegan and J. A. Earl OG-98 Cosmic Ray Electrons: A Discussion of Recent Observations 360 D. Miiller and T. Prince XIV OG-99 A New Measurement of the Cosmic Ray Electron Spectrum from 10 GeV to 300 GeV G. Hartmann, D. Mutler and T. Prince
OG-100 High Energy Primary Electron Spectrum Observed by the Emulsion Chamber E.Alzu, H.Htralwa, M.Fujll, J.Nlshlmura, T.Tiiira. T. Kobayashl, K.Nlu, J.J. Lord, R. J.Utlkcs nnd R. Golded (Abstract)
OG-101 Galactic Radio Emission, the Emlsstvtty In the Direction of H Reglone and the Local Interstellar Electron Spectrum 373 VT.R. Webber
OG-102 On the Composition of Electron Component of Cosmic Rays Accelerated by Pulsars 379 V. V. Usov (Abstract)
OG-v.03 Radloemission of the Galaxies'a Halos and Electrons of Cosmic Rays 380 V.A.Dogiel (Abstract)
OG-104 Origin of Cosmic Ray Positrons 381 M. Gller, J.Wdowczyk and A. W. Wolfendale
OG-105 Interpretation of the Distribution of Synchrotron Radiation In the Galaxy Using a New Model of Galactic Spiral Structure 386 C. Brindte, D. K. French and J. L. Osborne
OG-215 Cosmic Electrons, Galactic Radio Background and Cosmic Ray Confinement 392 G. D. Badhwar, R. R. Daniel and S. A. Stephens
OG-216 Secondary Positrons and Electrons In the Cosmtc Radiation 398 G. D. Badhwar and S. A. Stephens
OG-220 Measurement of Negatron and Positron Spectra from 5 to 40 GV Using a Magnetic Spectrometer 404 G.D.Badhwar, R.R.Daniel, T.Cleghorn, R. L. Golden, J. L. Lacy, S. A. Stephens and J. E. Zlpse (Abstract)
ORIGIN AND TRANSPORT OF COSMIC RAYS
OG-106 The Nature of the Observed Cosmic Ray Spectrum 405 L. I. Dorman
OG-107 Systematic and Stochastic Acceleration of Cosmic Rays by Radiation 422 V. M. Charugin and Yu. P. Ochelkcv
OG-108 Cosmic Rays and Clumpy Magnetic Fields 428 P.Freier and C. J. Waddlngton (Abstract) XV OG-109 Fermi Acceleration and the Structure of Interstellar Turbulence * 29 J.R.Jokiptt
OG-110 Cosmic Ray Propagation In a Closed Galaxy B. Peters and N. J. Westergaard 43s
OC-111 Fluctuations of 10 eV Cosmic Rays G. Erdbs, T. Gombosi, J. Kota, A.J.Owens, A.J. Somogyl 441 and A. Varga бсш
CUo-Ь OUSLRVATIUNS Of GALACTIC GAHMA-RAV MISSION (LARCl SCAU.) J.. . TJUI, K. Bennetf, G.F. Blgnaml, R. Buccherl, W. Ilermeen, G. КлпЬлсп? Г. Lebrun, H.A. Мауег-НлчзсIwanderl G. Plccinottl? L. Scirai) r. sorok*^ 4.U. Swjnenburg! and R.U. Wills' THE CARAVANE COLLABORATION 1. Cosmic Ray Workinj ^roup, liuygons Laboratory. Leiden 2. Laboratorio dl Flfica Cosmica е Tocnoloqie Kolatl.e del CI.R, Univernti di Milano 3. Laboratono di Flsica Cosmica е Tecnologle Rolativc del CNR, L'niveruiti di Palermo 4. Max Planck-Instltut fur uxtraterre&trlsche Physik. Carching uel MUnclien J> Service d'Electronique Physique, Centre d'Etude» Nuclealres dc Saclay 6. Space Science Department oi the European Space Agency, LSZ'EC, Noordwijk ABSTRACT Most of the reqlons of the sky already observed by COS-Е contain part of the galactic disc. Detailed analysis has been completed of the data from many of these observation periods. In several cases a manual scanning, effective in removing the background events has been incorpo rated in the processing sequence. The data collected during most of these periods have been merged in order to reveal the structure of the large- scale emission. This large-scale emission features a ridge centered on the galactic equator, showing a variable latitude thickness and signifi cant enhancements along the ridge. While such emission is very bright to wards the inner galaxy (320* < lrI < 40°) it becomes v*»ry wean, and in places not apparent, outside the solar circle. The energy spectra of the large scale emission from selected regions have been measured. The spectra derived for different longitude oands do not differ, the form of which resembles neither a pure n° nor a power law form. 1. INTRODUCTION The continuing success of the COS-B mission has enabled a systematic survey of the gamma-ray emission from the Galaxy to be under taken. At the time of writing, data from 10 of the observations periods in vhtch COS-Е was pointed tcv/ardE the galactic plane have been analysed. The details of the observations are given by Bennett et al., 1977b. The same alphabetic designation of the periods is used in this paper. Only 5 of them have been subjected to manual scanning, in these cases events of classes 2 and 22 were selected (Scarsi et al., 1977). For the remainder only automatically analysed data are available) for these only class 22 events were used. With both these selections, the instrumental non-gamma- ray background is negligible, while the angular resolutions and energy responses are also comparable. It has been determined that the sen»itivi- ties of the instrument are approximately equal in these cases. No sensi tive correction factors have so far been introduced to *ake account of the variation of the detector response as a function of time, which is governed mainly by the performance of the spark chamber between flushing operations. As reported by Bennett et al. (1977b) these variations are of secondary importance. With the exception of PSR0833-45 and the 4 sources in the anticentre region (Bennett et al., 1977a), no attempt has been made to ex clude contributions from localised sources discussed in detail by Hermsen et al (1977). From the above considerations it is clear that this paper presents only a status report of the investigations into the large-scale average emission from the Galaxy.
2. OBSERVED LATITUDE STRUCTURE OF THK GALACTIC EMISSION For each Obser vation period the I11 profile for gamma-ray emission of energy greater than 100 MeV averaged over longitudes within 15° of the pointing direction, has been derived. A detailed comparison of the latitude profiles derived from edited (manually scanned) and unedited data Indicates that the angu lar resolution is independent of this aspect of the analysis (Bennett et al., 1977b). trlbutlon of зллмпа-гау inten slty was derived for a band of 12 degrees In latitude centred on the equator. Data from longitudes within IS degree» from the detector axis-were used. The average background level! In the longitude distribution were determined from the high la titude fluxes In the distri butions of Figure 1. A composite longitude pro file of the total data Is shown in Figure 4. E'en though the data have not «f (Г been fully corrected, the MUCTIC ISMItun agreement in overlapping Fig. 2: The variation with longitude of the obnervations is seen to be latitude at which the galactic gamma-ray emis good. The trapezoidal pro sion (qreater than 100 MeV) is half the maxi file of the bulk emission mum value in the longitude intervals. In the in the longitude interval dotted region the Vela source has been sup within 60» o* the gal*ctic pressed. centre and «.«• relatively low level of diffuse emission in the rest of the disc is in good accordance with the now familiar picture of the galactic emission derived from the SAS-2 data (Flchtel et al., 1975). The higher V) sensitivity of COS-B resulting from long t— observation periods, combined with the * z effective angular resolution of 2 deg 3 (Scarsi et al., 1977) has permitted 2 ^_ degree longitude bins to be adopted. fc This presentation clearly highlights < many of the regions of intense localised £ emission discusseu in detail by Hermsen :— et al. (1977). The amplitudes of the реакз due to possible sources in these tj regions should not be taken to reflect "— the absolute or even relative lntensi- > ties since the restriction of |Ь"| < 6° — in the event selection has inevitably \Q excluded different fractions of the £ counts. jT Within 30» of longitude from galactic Z centre the avera^- flux is in good agree-- ment with the SAS-2 value, when account у is taken of the fact that about 20% of o the total intensity will be excluded by t the restriction |b"| < 6°. At longi- S tudes greater than 60° from the galactic i centre the width of the emission begins to increase resulting in a greater under estimate of the total disc emission in these regions. Therefore, even taking account of the localised emission, the
Fig. 3: Latitude profiles of the inten sity of gamma-rays of energy > 300 MeV, with Gaussian fits to the 6 bins of highest intensity. 3
The latitude profit»» for each uf the Ю observations Are presented In Figure 1 In which the Intensity scales are the SUM for each period, subject to the qualification regarding sensitivity varia tions already mentioned. (The enhancesient at high negative latitudes in the profile for period Е should be disregarded: it la caused by terrestrial albedo radia .w tion detected during part of the orbit, when the earth was in the field of view.) The narrow latitude profile observed for periods B, G. H, and M is consistent with that expected from the super- position of GauMmian distri butions of different widths, corresponding to the varyinq • « flf«t*«lM angular resolution of the experiment over the wide range of energies. At larger angular distances from the galactic centre, the emission region increases in width. In the anticentre region (period A) there is no evi dence for any disc emission. In Figure 2 the observed width of the disc is plotted as a function of longitude. The lines mark the latitude at which the uxuC iliteiiMity averaged over the longitude interval indicated is half I m,
» v Л 5
flux In the» region» 1» generally higher than that iwaaurad by SAS-2. It should ba borne in Mind, howavar, that In rations of weak broad oatlaalon It la difficult with tha praaant data to aatl- mata tha high latitude background. An undar-estiaata In tha background level result* In an over-estimate of the ga ID lactic flux. 4. THE ENERGY SPECTRUM The energy spec • true of the galactic amisalon has been derived for 4 of the observation periods for which Manually edited data vera 'E available- In order to t«ke advantage оГ tha batter energy resolution for gamma- raya Incident at small angles to the detector axis, tha fluxes were inte grated over longitude intervals within IS* of the pointing directions. The method adopted has been described by Bennett et al., 1977. The resultinq spectra for longitudes centered on'11" ЗГ-, а" ж 75», x11 - 125* and l11 - 320° are remarkably similar in shape, espe cially vhen it is realised that contri 0.1 1 bution! from point-like sources are photon energy (GeV) present to a greater or lesser extent in the different regions. A typical spectrum is shown in Figure 5 which Fig. S: Energy spectra of the is compared with various model spectra gamma radiation from the galac fitted to the data at 600 MeV. The tic disc. The dashed and dotted broken lines in Figure 5 are n» spectra lines are the n* spectra of as predicted by Stecker (1970) and by Stecker and of Badhwar and Ste Badhwar and Stephens (1976). The solid phens respectively. The solid li«<* is the br*m*firr«hlimg sp*ctriun line is the bremsstrahlung predicted by Higdon (1974). The re spectrum Dredicted bv Hiadon. sponse of C0S-B is such that the input spectrrja does not differ significantly from the observed spectrum for events within 15"of the detector axis in the case of either the n* rr E~2 examples (Scarsl et al,, 1977). 5. CONCLUDING REMARKS The sensitive survey of the galactic emission by C0S-B is already providing detailed information of its spatial and spec tral nature. The wirde emission in regions away from the galactic centre suggests strongly that the production region is relatively close. The Intense emission in the central region is narrow and exhibits signifi cant fine structure. The component of the central emission at the positive latitudes in the longitude range 345» < i11 < 70* is reminiscent of the band of dark nebu lae detected in this region by Lynds (1962). In particular the position of the emission maximum localised in the range 25* < I11 < 35* correlated with a wide and very dark complex in the Lynds survey. The energy spectra at different longitudes are indistinguishable from each other, thus preventing recognition of different source mechanisms at different regions of the disc. The spectra resemble neither a single power law nor pion-decay form. If the spectrum is to be interpreted as a mixture of bremsstrahlung and n* decay the contribution of bremsstrah lung is substantially higher than suggested by Fichtel et al. (1976). In fact.the data are consistent with a pure bremsstrahlung (Higdon, 1974). COS-B has now completed the systematic survey of the galactic plane and the data already acquired will permit the remaining gaps to be filled. Planned repeated"observations of several regions will increase the sta tistical significance of the data, important at high energies to resolve 6
localised affect*, and will provide a direct determination of trie varia tion of the detector response. This will allow the fine correction fac tors to be determined thus permitting more precise merglnq of datrt from overlapping regions.
REFERENCES Badhwar, CD. and Stephens, S.A., 1976. presented at Gamma Ray Symposia. G.S.F.C. Bennett, K.j Blgnaml, G.F., Bonnardeau, M., Buccheri, R., Hermaen, W.. Kanbach, G., Lichtl, G.G., Mayer-Hasselwander, Н.Л., Paul, J.A., scarsi L., Stlglitz, B., Swanenburg, B.N., and Wills, R.D., 1977a, Astron. and Astrophy*. 56, 469. Bennett, K., Bignami, G.F., Buccheri, R., Hermsen, W., Kanbach, G., Mayer-Hasselwander, H.A., Paul, J.A., Plccerottl, G., Scarsi, L., Soroka, F., Swanenburg, B.N., and Wills, R.D., 1977b, Proc. 12th ESLAB Symposium, Frascati. Fichtel, C.E., Hartman, R.C., Knlffen, D.A., Thompson, D.J., Bignami, G. F., Ogelman, H., Otzel, M.F., and TUmer, T., 1975, Astrophys. J. 198, 163. Fichtel, C.E., Kniffen, D.A., 'Thompson, D.J., Bignami, G.F., Cheung, C.Y., 1976, Ap.J. 208, 211. Hermsen, W., Bennett, K., Bignami, G.F., Boella, G., Buccheri, R. , Hic- don, J., Kanbach, G., Lichti, G.G., 'lasnou, J.L., Мауег-Hasseiwander, H. A., Paul,.J.A., Scarsi, L., Swanenburg, B.M., Taylor, B.G., Wills, R.D., 1977, Proc. 12th ESLAB Symposium, Frascati. Higdon, J.D., 1974, Ph.D. Thesis, University of California. Lynds, B.T., 1962, Ap.J. Suppl. 2' No. 64, 1. Scarsi, L., Bennett, K., Bignami, G.F., Boella, G. Buccheri, R., Hermsen, w., Koch, L., Mayer-Hasselwander, Ei.A., Paul, J.A., pfeffermann, E., Stiglitz, R., Swanenburg, B.U., Taylor, B.G., and Wills, R.D., 1977, proceedings 12th ESLAB Symposium, Frascati.
Stecker, F.w., 1970, Astrophys. and Space Science 6., 377. 7
OBSERVATIONS OF GAMMA RADIATION FROM THE GALACTIl. CLNTEH"REGION J Samimi arid A. Jabbari-Azad Ferdowsi University, Mashhad, Iran R.L. Kinzer and G.H. Share U. S. Naval Research Laboratory, Washington, D.C.
Theoretical Ц Experiment»! Щ Bo,h D 3alloon-borne observations of gamma-ray emission from the galactic center region with the NRL-Mashhad telescope were made in 1971 ind 1975; analysis of the 1971 data is completed. The galactic emission detected above -15 MeV appears concentrated in an -2* wide band lying along the galactic equator. Assuming uniform emission from the plane, the 15-70 MeV and 70-200 MeV diffuse galactic fluxes (F) are (1.3+0.6) x Ю"6 and (1.0+0.5) x 10~6 —2 —1 -1 -1 cm sec rad MeV , respectively. However, because the emission appears to be concentrated in the interval 350*<£II<360*, it may be nisleading to compare these diffuse fluxes with those from other experiments when attempting to determine the galactic energy spec trum. The ratio F(15-70 MeV)/F(70-200 MeV)-l for a ire-decay spectrum; the data, although consistent with this ratio, suggest a somewhat softer spectrum. Preliminary results from a 10 hour exposure obtained in 1975 with a modified telescope are also presented; this observation has an area x time x efficiency factor 10 x that of the 1971 flights.
Coordinates:
OG-1.1 (Diffuse Cosmic and Galactic Gamma Rays)
Mailing address: Robert L. Kinzer Code 7128.3 Naval Research Laboratory Washington, D.C. 20375 GAMMA RAY SPECTRUM FROM THE GALACTIC DISK
R.K. Sood. G.K. Rochester, P.G. Cla/ton, B.T. Thomos Bl«ck«tt Laboratory, Imperial College, London SW7 2BZ
A spectrometer with area 1200 cm2 sensitive to gamma rays In the energy range 200 MeV < Е < 8 GeV has been flown on balloons over Australia fn 1973 and 1975- One of the flights lasted 60 hours and a total exposure of 15 hours to the Galactic plane In the region 330° < lu < 30° - 10e < b11 < 10° was obtained. Preliminary results from the 1975 flights when added to data from, the 1973 flights give a clearer picture of the gamma spectrum.
1. Introduction
Observations of gamma ray emission from the central region of the Galac tic plane In the energy range 200 MeV to 8 GeV have been made using balloon- borne Instruments flow;) over Australia. The first flight on 21st November 1973 lasted 8.hours at altitude and the results were reported at the cosmic ray conference in Munich (Sood et al. 1975). Since then a further two flights of identical detectors have taken place at residual atmospheric depths of between 3 mb and 5 mb, the second of which floated from 7-30 a.m. on November 18th, 1975 until 5.30 p.m. on November 20th.
The detectors were calibrated in the tagged photon beam at DESY. Hamburg, which was set up for COS В and which has also been used to calibrate SAS II. Comparison of the results obtained on the balloon flights with those obtained by the satellites therefore allows deductions to be made about the spectrum of the Galactic emission up to photon energies of 5 GeV.
2, Method and Results
The experimental technique employed has already been described together with an outline of#£he more Important response functions of the Instrument (Sood et al. 1975), *e attention here Is restricted to presentation of the results.
Figure 1 shows a contour map In Galactic Coordinates of the part of the sky scanned during the three flights. All flights covered essentially the same region, the exposure on the 58-hour flight being restricted in the analysis so that only that period during which the balloon was floating at a relatively constant altitude during daylight hours Is Included. In the case of the long flight this analysis is not yet complete so the results quoted here are derived only from the flights which took place on 21st November 1973 and 7th November 1975. Figure 2 shows the observed flux in arbitrary units plotted as a function of Galactic 'latitude In the energy range 600 < Е < 1400 HeV,- in which the detectors were most sensitive, for all longitudes covered by the exposure sjtown In Figure tf(a) for the 1973 flight alone and (b) with data added from the first 1975 flight. Addition of this extra data reduces the 2a upper limits and flux previously quoted (Sood et al. 1975). The new values are compared with a differential TV° decay gamma spectrum In Figure 3. 9
300 60 GALACTIC LONGITUDE
Fig. 1 Exposure map for the balloon flight on 21.11.1973- All the flights covered essentially the same region of the sky.
This spectrum has been normalised to an integral flux of 2.10"1* photons cm"2sec"1rad"1 > 50 MeV corresponding approximately to measurements reported by Fichtel et al. (1975)- Discussion of the connection between these results and the satellite observations will follow completion of the analysis of the data from the 58-hour flight. 10
Fig. 2 Observed flux as a function of Galactic latitude. (a) Data from f1ights on 21.11.1973 and 21.11.1973 + 7.11.1975.
'40 -20 0 20 40 60
(b)
«0
Galactic Latitude (degrees) It
^1 *
«r» -
• J—1— 1 11 i J . Г
3 .0-
\ :
(MB 04» (H M 03 M 1 131
Photon Energy (GeV)
Fig. 3 Differential spectrum of the Gamma flux observed from the Galactic Plane.
References
Sood, R.K., Bennett, K., Clayton, P.G. and Rochester, G.K., Proc.1<»th Int.F.R.Conf. Munich I, 35. 1975 FIchtel, C.E., Hartman, R.C., KnlTfen, O.A., Thompson, D.J., Bignami, G.F. and Ogolman, H., Prod'tth Int.C.R.Conf. Hunich 1, 29, 1975 штт 12 OBSBRVATIOIS ОГ VBRY HIGH HZRGY OAMU RADIATIOM II DIHBCTIOI OF QAL1CTIC PLAH.
V.P.Vomin, B.M.?ladlmiraky, A.A.Stepanian Crimean Astrophysloal Obeervatory, Crimea, 334413, USSR
The results of four-years of observations of the galactic plane by means of atmoapheric Cerenkov techni que* сос to one* The average amplitude dered longitude* ie found to be - 0.7* with confidence level (1-10-5). it is auggeated that the galactic diak is the source of very high-energy -f -emission which ia ge nerated by the inverse Compton aoattering of cosmic ele ctrons on the light of stare. The decrease of the inten sity in the direction of the galactic equator results from the decrease of photon density in the vicinity of the equator, because of the absorption by interstellar dust»
Very high-energy f-ray observations were made in the Cri mea using Cerenkov flash detection of BAS to aearch for disc rete a*ourcea, ae early aa 1969 (the results were described in fl]). Particular attention waa paid to observations of the ga lactic plane. So the galactic equator has been creased in eight galactic longitudes during these measurements. The counting rate profile was obtained in the vicinity of the equator, and the decrease of counting rate of BAS waa revealed for the la titude Interval \&>~\& 1°»5. The preliminary results are ahown in [2,3]. The dependence of BAS oountlng rate on galactic latitude is shown in Fig.1 for Cassiopeia (о~аб1°30() and Cygnua (£* • -35°20') regions, as well aa the summary ourve for the other six longitudes. The statistical weight of the measurements is taken into account to calculate this curve. The mean amplitude of the decrease is 0.7* +0.2J6. Decrease of intensity at the equ ator is observed for each longitude taken separately. The pro bability of having negative amplitudes seven times by chance 13
dees not exceed 10 . The -I 1 I 1 I I Г- analysis made in £>J has CassJq^o shown that such counting 5*4 1.00 rate behaviour oould not 0.ff v be connected with the ap-• paratus, м it ii i real f 'Clonus variation of the intensi < tot ty. The observations of other authors, who used "^ toi 3 100 the seme type of the equi- S Q99 upaent, were studied. The results of scanning the -15 -10 SO 5 10 IS objects in the Yicinity of the galactic equator are Gafacttc tatitude,td«pus presented In Table 1. Щ. i.
Table 1,
Object effect error energy Beferenoes % e* ет 6*
Cas A -0.6 ±0.43 113° 5#1012 -1e.5 ChudakOT et •1 [4] 11 Cas A -0.67 ±0.85 _«_ 2»10 —«» Weekefts a t el S»1572 +0.22 ±1.19 120° 2-1011 +1e.5 CP0329 -0.89 ±1.06 144° 2»1011 -1e.2 -»- Galac tic -1.48 10.96 0° 2»1011 o* Grindlay et al center f6j GZ340+0 -1.20 •1.1 340» 2-1011 0" —n_ GX1+4 -2.47 ±2.14 1° 2«1011 4° —»_
One can see that, apart from supernova remnant 1572, all the sources have negative effects. SSR 1572 observations have poor statistics. The scanning for this case is made nearly along the galactic equator, so the expected effect is small. The probabi- и
OaCactic lonQitu.de , degrees
'* 2- lity of obtaining negative deviations by chance is only 6*10 for all except one. We may conclude that there is real decrease of EAS intensity in the direction of galactic equator, as reve aled by our data and the data of the other authors. The decrease of EAS intensity appears to depend on galac tic longitude. Thie is shown in Pig.2 for all the measurements. The errors are large in each case and it is not possible to derive any decisive conclusion. One can see, however, that the effect is observed for all directions. The decrease is present in a narrow interval of galactic latitudes. For this reason, we asaume that it is connected with gamma-ray quanta of the energy >1012 ev. Several mechanisms for the above described phenomena were considered. A similar distribution of gamma-ray quanta with the energy c£ 10 ev was observed' in [7]. The authors assume that their results may be explained by the distribution of discrete gamma-sources. Here we show that the effect revealed in the 19 10 ev range may be interpreted as the emission of the galac tic disk in very high-energy gamma-quanta, generated via inver se Compton scattering of cosmic ray electrons on star light. The decrease of the f -ray intensity in the vicinity of the equator is caused by diminishing thermal photon number densi ties, because of interstellar absorption caused by the dust. 15 Calculation* of the spatial diatribution of thermal photons relati ve to the galactic pla § ne were made, taking in to account the diatri •V) bution of «tare of dif ferent spectral types fеt and the distribution of o galactic dust. Typical o results are presented 3 - in Pig.3. It may be seen 4 that a large decrease of i4 photon density is obser ved near the plane with § a halfwidth-250 pc. The maximum photon den ro too looo sity is reached at a dis "Distance from qotacticptanelpc) tance from the galactic 4 3- plane of about 600 pc, and the density decreases again at lar ger distances. The difference between energy density at the equator and the maximum energy density is as large as 2 • 3 ev. cm . Although some approximations were mads In this calcula tion (e.g. radial distributions of the stars and the dust are not used) the qualitative picture of photon distribution is, we hope, correct. Using the distribution obtained and a well known equation for the intensity of f -quanta produced by in verse Compton scattering of electrons [в] , one may calculate y'-emission intensity as a function of galactic latitude. It was suggested that the power of the cosmic electron spectrum is 2.7» up to the energies 10 ev, and the density of these electrons is the same as near the Earth. As a result, у-ray surplus flux reaches the largest value near/£"(— 5° • 10°. If the field of view of the Cerenkov detector is 7*10~4 eterad the flux is equal to 4*10 quanta cm sec eterad . In the direction of galactic plane jt l<1°,5 the flux is 2«10~11 quanta cm sec 'eterad '. This change of flux corresponds app 16
roxiaately to 0.5% of ooaaio ray background. This value la la good accordance with the observational data preeented. Thua the daoraaae of counting rat«a of KA3 at the gelec- tic equator may be explaned if we aaauae that inveree Coapton •f -quanta make conaiderable contribution to the "f -emission of galactic diak.
References
1. A.A.Stepanian, B.M.Vladimireky, Yu.I.Heshpor, V.P.Pomin 1975, Ap.Space Sci., 2§, 267. 2. A.A.Stepanian, B.M.Vladimireky, Yu.I.Veahpor, V.P.Pomin 1975, Izv.Crimean Aatrof.Obe. Ц, 29. 3. V.P.Fomin, 1976, Izv.Crinean Aetrof.Obe. Jj6, 29* 4. A.E.Chudakov, V.L.Dadykin, B.I.Zateepin, F.M.Heeterova 1964, Trudy (Proc.) Lebedev.Phys.lnst., vol 26, p. 118.
5. T.C. Weekes, G.G.Fazio, H.P.Helmken, EcP.O*Mongain, G.H.Rie- ke, 1972, Ap.J., ГЦ, 165. 6. J.E.Grindlay, H.P.Helmken, R.H.Brown, J.Davie, L.R.Allen, 1975, Ap.J., 201, 82. 7. C.P.McKechnie, K.E.Mount, D.Ramsden, 1976, Ap.J.Lett., 207. L151. 8. V.L.Ginsburg, S.I.Syrovatekii, 1963, Origin of Cosmic rays, Moscow, Academy of Sci.Publ.House. SCK&JY$
GAsUA HADIATIOsT FBOM THK СИГГЕ11 HEGIOB" OF ТШС GALiXT AHD DIFFUSE BACKfiOOID AT I,>100 MB? A.I.BelJaevsky, Y.L*Bokov> V.I.Bonharlrln, I.F.Bugakov, G.M.Gorodinsky, I.M.Kruglov, B.V.Mjakinln, G.A.PJatigorsky, B.I.Chujkln, V.S.Juferev, J.A.ShiDanov, E.V.Kolesnikova and B.A.Beloborodko.
A.F.Ioffe Physical-Technical Institute, Academy of Sciences of the USSR, 194021 Leningrad, DSSB
A spark chamber gamma-telescope was flown on the "CoaBOs-561" satellite on May 1973* The telescope axis was oriented all the tine to Zenith. A band of the sky with incxination 64,5° and located within coordinates S ж ± 60* and «* from 160* to 310* and from -30* to 120* was investigated. The observation tine was 103 sec. In an analysis of the elektron pairs in the spark chamber 169 Quanta with the ener gy more than 125 Her were revealed. The investiga tion of the distribution of these events on the sky showed an excess of quanta coming from the band stretched along the equator with coordinates <* = + 10*, t* = -35* -_*25e. The flux from this band - X 1»77 ± 0.7* ),ЛО~+ quanta cnT^sec"'rad"1 - is fo und to be in satisfactory agreement with other mea surements. She differential gamma spectrum of this radiation with the energy from 125 U*r to 700 Her was obtained for the first time. The spectral index is - 2,52 ± 0,49» An estimation of the upper limit of the diffuse background gives the value of ( 7,1 ± 2 1 3t5 )»10"? quanta cm- see-ieter- • 1. Introduction. The purpose of the experiment was to search xor local sources of cosmic У-rays with a flux above 100 Mev exceeding 10~* quanta cn"2eec~1. A gamma-ray telescope was laun ched on board the "Cosnos-561" satellite on may 25, 1973* "Cos- aos-5bi" had an orbit with an inclination of 64-, 5*, an altitude ranging~froa 214 km to 316 km, and a period about 89,3 ainf'J . For seven days of flight the band of the sky depicted in Fig. 1 was scanned. The data related to the events registered during the gamma- telescope blooking periods have been published earlier ^. The existence of hard jr-radiation coming from the region contai ning the galactic equator with l' = -30* - +10° was established. The measured flux is f=^(>100 Mev ) = ( 8,3 * 2,7 ). 1Q~5 (s
quanta ca~2aeo""1. Upper liaits for the frequency of bard raya burata (< 2.10"** borate aeo-'eter,-1. .-1-'-) and for the energy par buret (< 10~5 erge c«~2 buret-1) tare Ъееа determined fro» an analysis of the dietribution of tiave Intervale with a siren nuaber of gaama events per Interval*
Fig. 1•- A aap of the sky in ooordinatea «*.t § Points denote the gamma-quanta registered*
2. у-ray telescope description. ТЬл у-ray telescope is sche- natiealy shown in Fig. 2 ^3]. д gaaaa-quantun converts in Pb converter into an electron pair which activates plastic scin tillator detectors D1, D2 and eherenkov counter Ch connected to a coincidence circuit* She plastic scintillator detectors 03» D4-, D5 are connected to an anticoincidence circuit and serve to discriminate against charged particles. The experimentally de termined efficiencies of these detectors were £3^ 0.995 «£*^
0,9996 , £5 5= 0,998 .
Fig.2.- Telescope arrangeaent. Fb- lead converter (260x200x1,5 an'), D1, D2- pla stic scintillator detectors measuring 260z2u0z50 and
The detector 01 threshold was established at auch a level that 1)1,1)2- telescope registered approximately 13% of cosalc^- mesons at the saa level. The alx gap spark chaaber S1 т Sb le triggered by «Tents that satisfy the logic D1 иг Uh itf w u> . Such events may be due to у-quanta or soae interfering pneno- mena and are called у-quantum -like «rente. The information concerning spark coordinates, double and triple coincidence rates, Ch amplitude is stored in an opera- tiTe memory during a blocking tine t. = 3?i55 sec. The gamma- quanta are registered during this time but without triggering the spark chamber, ill the detector rates and other telescope parameters were registered in-flight. For in-rlight calibration, the modes of detector and electronics operation could be chan ged. The full weight ox the gamma-ray telescope was 190 kg, and electronic power supply •• 2? w , 3. Telescope characteristics» 'line effective gamma-ray teles cope ax-«et c«n be expressed as a product
of Z0(Er) -an effective area at vertical incidence of gamma- quanta ( 9t = Ол вг *0) and q (et.&i,Bf)- a directional diagram, where 0< and 0& are projections-!, angles, and 4(0,0,Ef) = i.
If a power law is valid for у-ray flux dP/^£r~£f tnea 5tff (9<,B*,Et) = SeC~)-/2 Denote WL(£f,u) a counting characteristic that can be ex pressed as follows ana ,/ W« (Er,<*Jc/Er = £ . 20 Hi* dependencies or £,If»-) and WcfFr,oO on the energy are mom In JTlg. 3. In Tig. 4 S. £°0 la represented. Fig. J,- Dependencies of the effec tive uu S„(fc"V)- S,ff(0,0,£r) at -vertical gamna-q.uan'ca incidence, and counting characteristics tt<^ • 1 and oi. ж 2 on 'Che energy. «WW W Pig. 4. - $.(<*) та. index <*• An experimental angolar distribution of gamma-quanta along the telescope axes was used to determine ^ (0 E> Mer 100 150 200 300 500 700 1000 8 degree 7,7 5,9 4,6 5,2 2,5 1|5 1,5 4. Data reduction and results. Approximately 10 spare cham ber triggerings hare been registered during the flight . And among them 7000 - in a mode of jf -quanta registration for lot 4 60°. Full observation time of these 7.10^ events (excluding blo cking time) was 10^ sec • All the events have been derided into three classes . The first 21 of them containing 33* of the «rente includes "false trlggexdmg" 2u«ee «rente were «ranly distributed oror tne tine of flight . The absence of «park» in the most of gape «ma characteristic fer these «rents* To the second class containing 60% of the «Teats belong eronts with a single track in four or more gaps. Booh eases could be produced partly by charged particles going from below, which than stopped in M or РЪ-еоптег*ег. Sis third class includes «rents in which more than one track was obserred. Pram 650 such «rents a pair at least in one projection has been «b- eerred in 236 cases, and 969 pairs hare satisfied all the cri teria of gamma-quanta selection in both projections • ill the observed region of the sky was dirided into 816 5°x 5e bins oriented along the satellite trajectory (fig. 1) • Ихе solid angle for each bin can be expressed as 0,0078.eos 6* ster. An exposition integral has been calculated for each bin • Hear the celestial equator it is T * 140.5v(°0emrsee star. Ill the аЪоте 169 ©rente are shown as points in Pig. 1. I-\ the low er left part of Pig. 1 along the celectial equator region 1 can be seen. It is bordered by a thick line and contains bins in li mits The intensity Jr from the sources situated along the galac tic equator within llи 325° f 23" has been obtained as a diffe rence of the intensities from the abore regions 1 and 2 and has 4 1 a ralue lv (>100 Iter) = (1,77 * 0,74).10" quanta cm^sec-^rad" which agrees with The background gamaa-radidtion was latitude-dependent, more probably due to D3 damage which did not function during the flight. Therefore the intensity for diffuse background radiati- . on 3rdif (*-100 Mer) = (7,1 ± 3.5).10-5 quanta cm-2see"1etsr"1 that has been determined for regions of geomagnetic latitude where the rigidity is ^ 16 Ger/c must be considered only as am upper limit for this flight • In Fig. 5 an integral spectrum of у-rays from region 1 is shown * It was obtained taking into account the contribution of the diffuse background and correction for S.(£*0 • In error boun dary shown by the dashed line in jrig. > is drawn accordingly to errors along both axes f5? . Xhe power law index of the integral 22 speotrua in the rang* 100 M«r < В j- < 700 Мет was found to be U. • 1,59 1 0,*9 . fig» 5»- Ganma-spectrue from region 1 of th* galactic equator shown in Tig. 1. No sources with a flux Fy(>100 Мет) аЪоте 1,5.10"5 quanta cuT^see*"1 for i&l ^- 20° and with a flux аЪоте 3,6.10~5 quanta cm"2sec~1 for /5") = 20° f 40° were found in all the ob served region of the sky at a significance lerel ^ 2<5 References» 1. "PraTda", 26 May, 1973» 2. l.I.Beljaersky, Т.Ь.Вокот et.al., Dorlady AH USSR,.219. 830, 197*. i Pis'ma т ZhETP, 21, 729, 1975» 3* А.1.Ве13аетвку, V.L.BOXOT et.al., Izrestija AH USSR eer.fiz., 28, 1838, 1974 ; A.I.BeljaeTsky and G.A.PJatigorsry, Preprint MI AN USSR N 291, Leningrad, 1970. 4. C.E.Fichtel, R.C.Hartman et.al», Ap.J», 198. 163, 1973 • 3* V.V.Fedoror "Theory of optimal experiment" ( in Russian), M., 1971. 23 SEARCH FOR PRIMARV GAMMA RAYS IN THE MEDIUM ENERGY RANGE 5 MeV - 50 MeV E. Bonfand, B. Parlier, B. Mougin, J, Andrejol, J,C. Courtois, A. Lecomte, M. Goriaee, B.Agrinier Service Electronique Physique, Centre Etudes'NucKair е • de Saclay, France J.M. Lavigne, C. Doulade, M. Niel, Centre d' Etude* Spttiales det Rayonnements, Toulouse, France R. Palmeira, K. R. Rao, I.N. P. E., Sao Jose- doa Campo», Brasil Theoretical Q Experimental [x] Both Q In a balloon flight at South latitude (-23*S B'-asil) performed the 26th of November 1976 we monitored the medium energy gamma rays at altitude ~ 3 mb with'a 1. m^ area spark chamber telescope (Agathe). About 20% of the 40, 000 pictures taken during ascent and ceiling are pairs and Compton tracks occuring in the 15 x 0.007 radiation length plates. After substraction of the atmospheric component we have looked for the Galactic Central region emission and for the Extragalactic component when the telescope axis was far away from the Galactic disc. Preliminary results of these searches are given in this paper. Coordinates: OG 1.1 Diffuse cosmic and galactic gamma rays Mailing address: B. AGRINIER DPh/EP/ES - BSt. 28 Centre d1 Etudes Nucle'aires de Saclay B.P. 2 - 91190 Gif-sur-Yvette (France) 24 GAMMA-RAYS OF 3 TO 25 MEV FROH THE CALACTIC ANTI-CENTER AND PULSAR NP 0532 Robert B. Wilson, Shin H. Moon, James M. Ryan, Allen D. Zych, R. Stephen White Physics Department and Institute of Geophysics and Planetary Physics University of California, Riverside, California 92S21, U.S.A. and Bruce Dayton Physics Department California State University, Los Angeles, California 90032, U.S.A. ABSTRACT. Gamma-rays of 3 to 25 MeV are reported from the galactic anti-cen ter region and the Crab Pulsar NP 0532. The observations were carried out with the UCR double Compton scatter gamma-ray telescope on a balloon launched from Palestine, Texas on May 13, 1975. Gamma-rays from the Galactic Anti- Center were observed as the Crab Nebula passed overhead within 10° of the zenith. Pulsed gamma-rays from NP 0532 were observed at a 4.4a significance level. Our total flux from 3-25 MeV is (4.9±2.0) x 10~3 photons/cm2-sec._ The pulsed flux from NP 0532 from 3 to 25 MeV is 4.3±2.6 x 10_l* photons cm~2s_1. The ratio of the total to the pulsed flux from 3 to 25 MeV is 11+8. 1. INTRODUCTION. Extensive observations of the gamma-rays from the Galactic Anti-Center Region and the Crab Pulsar NP 0532 have been carried out with con siderable accuracy using scintillation crystals at energies below about 1 MeV. At energies above 20 MeV there have been a number of observations with spark chambers. However, in the intermediate region of 1 to 20 MeV few observations exist and these have rather large uncertainties. For the rest of the paper, we will use the term "Crab Nebula" instead of "Galactic Anti-Center Region" although there appears to be no experiment to date, including ours, in the energy region of 1 to 25 MeV with the angular resolution to unambiguously separate the Crab Nebula from the other sources in the galactic anti-center region (Thompson e_t^ al., 1977). At low energies the distribution of total gamma-rays from the Crab Nebula fits the power law 4.0 x 10-^ E-2'* photons cm~2s~^MeV~l and the pulsed flux from NP 0532 fits the line 6.2 x 10-* E~2-2 photons cm-2s-1MeV-1 (Walraven et al. 1975). McBreen et al. (1973) prefer the fit to NP 0532 of 5.5 x lO'^E-2-1 that is only slightly different. Alternative fits have been given by Laros et al. (1973). There is some indication that most of the flux is pulsed above about 100 MeV (McBreen et al. 1973; Kniffen et al. 1974; Thompson et al. 1977). In the region of 1 to 20 MeV, the total fluxes of Baker et al. (1973) at 1.6, 3.8 and 8 MeV are factors of 5, 20 and 40 above the total flux line of Walraven et al. (1975) but have large uncertainties. In the region from 1 to 10 MeV the data of Walraven et al. (1975) are scattered with the point at 6 MeV a factor of 15 and 2.6a above the line and the point at 4 MeV an upper limit only. Above 35 MeV the total flux has been measured by Kniffen et al. (1974) with the SAS-2 spark chamber. The flux of pulsed gamma-rays from NP 0532 measured by Kurfess (1971) at 2 MeV fits the pulsed line of Walraven et al. but with a rather large uncer tainty. The point of Kinzer et al. at an energy of 10 MeV is a factor of 10 above the line but has a large uncertainty. The point of Albats et al (1972) la a factor of 10 and 2a above the line. The higher energy data scatter somewhat but are in reasonable agreement. The UCR double Compton scatter IS gamma-ray telascopa It designed to take advantage of Compton scattering that is the dominant Interaction at gamma-ray energies of 3 to 25 HcV. 2. OBSERVATIOHS. The observations were carried out with the University of California, Riverside, UCR, double Coaipton scatter gamma-ray telescope (Herzo et al. 1975; Zych at al. 197S) on a 20 ail Hon cubic foot balloon launched fro» Palestine, Texaa on May 13, 1975. Measurements were made for 6 hours on either aide of cloaeat approach as*the Crab Pulsar passed overhead within 10* of the zenith at 2035 U.T. on May 13, 1975. The balloon waa at an altitude of 3.5 g/cm2 of residual atmosphere during these measurements. The flux of Incident gamma-rays and their energies and zenith anglea were measured by Compton scatters that occur in each of the two liquid scin tillator tanks, SI and S2, shown schematically In Figure 1. SI is 100 са х 100 cm by 12.5 cm and S2 is 100 cm x 100 cm x 20 cm. The tanks are spaced 100 cm apart from center to center. The incident gamma-ray, Y, scatters from an electron, ej, in SI through the angle, 6, continues on as gamma-ray, Yi, scatters from an electron, ег, in S2 and loses all or part of its energy in S2. If YJ deposits all of its energy in S2 the kinematics are uniquely defined and the energy and scattering angle of the incident gamma-ray are determined. However, Y2 escapes sometimes without depositing all of its energy so a distribution energies and angles results. The errors in the as signed energy and angle are evaluated in Figure 2 of Ryan et al. (1977), paper OG-16. For a gamma- ray of 6.1 MeV the energy resolution varies from a HWHM of 24Z at a scattering At ANTICOfNCIOENCE angle of 10° to 5Z at a SHIE10 scattering angle of 62°. The angular resolution has a HWHM that varies from about S" to 12° at angles from 10" to 62". The tanks SI and S2 are each divided into 28 cells for better angular resolution and each cell is observed by a separate photomultiplier. The outside 12 cells are 25 cm x 25 cm and the inside 16 cells are 12.5 сш х 12.5 cm. Plastic scintillator 0.6 cm thick completely sur rounds each tank with its SZ UOUID SCINTILLATION OETECTOm (2в) photomultipliers to veto FIG. 1. Schematic section through the tele charged particles. The scope. The incident gamma-ray, y, scatters pulse heights in SI and S2, from an electron, ei, in scintillator, SI, the times of flight and the continues on as Yi, scatters for electron, cell identifications are SI and S2 are each divided into 28 cells digitized and telemetered to iz- and each cell is viewed by a separate photo- a receiver on the ground multiplier for better angular resolution. SI where the data were stored and S2 are each separately surrounded by on magnetic tape. plastic scintillator to veto all charged particles. 26 For thlt paper data la pre sented only fro* vert leal palra of cells In SI and S2. In this case, the scattering angle la the saaa as the zenith angle so the zenith angle la directly measured. The flux in a given zenith angular Interval la plotted aa • function of tlae. When a aource enters this zenith angle interval an increased counting rate is observed. From the increased counting rate and the telescope efficiency, using vertical cells only, the gamma- ray flux Is determined. A vectron CO-20V-5 oscilla tor with a stability better than 1 ^art in 109 per day assigned the tlae of. each event. This time wat digitized to an accuracy of 50ps then transmitted by telemetry along with the other information. -Prior to launch 1400 1600 ЙО0 2000 2200 2400 0200 1973 I3MAY I и MAY it was synchronized with Loran UNIVERSAL TIME I C at the launch site at Palestine, Texas. Unfortunately, during the launch preparation, the clock was inadvertantly reset so the synchronisation time was lost. FIG. 2—Total flux from the Crab Nebula Our relative timing, however, from 2 to 25 MeV during its passage through is good to SO us. the telescope aperture. Vertical cell pairs only are used. The flux is shown The NP 0532 pulse period for three angular aperture* of 10*-20*. and change In phase for the 20°-30e and 30*-40\ The aolid lines are epoch of our flight at the. the expected fluxes on the basis of the position of the telescope was measured angular resolution (Ryan et al. kindly furnished by John M. 1977, paper 06-16). Rankin of the National Astronomy and Ionosphere Center, Axecibo, P.R. This Information was incorporated Into our program to determine the pulsed flux and was updated every ten minutes for the new period and pfiase. 3. RESULTS. The flux is calculated using vertical cell pairs only. Data from 4 sets of four small adjacent cells are combined so that the 16 'small center cells in each tank form the equivalent of 4 large cells. The double scatters from the 16 equivalent large pairs are combined in energy and divided into zenith angle intervale of 10"-20*, 20*-3O* and 30*-40*. These are plotted in Figure 2 versus universal time for the epoch Hay 13, 1975. The flux of gamma-rays increases in the Interval of 10*-20* for one hour on either side of the maximum at 2035 U.T. when the Crab Nebula is nearest the zenith. A broader peak is seen in the 20*-30° flux as expected with a higher background caused by the larger solid angle for atmospheric gamma-rays. The 30*-40> fluxes give the peaking at earlier and later times as expected as the Crab Nebula passes through that angular aperture. It also has the 27 highest background flux. The solid line* give tha fluxta expected on tea bails of telescope resolution a* tha Crab Mabula paaaaa thromgh tha thraa celcacopa apcrturaa. Ho aource flux члш obaarvad In tha 0* to 10* apartura. Thla la expected aa tha angular resolution, particularly at email angles, cuta off aharply below tha maximum but haa a sloping tall Coward larger angles (Ryan at al. 1977, paper OG-16). The aourca flux la found from tha difference of the average total flux In tha 10*-40* Interval when the source Is present and whan It la absent. The result la 4.8 t 2.0 x 10"J photons caj~2a~1 at energlea fro» 3 to 25 HeV. Thla Is 5 times the flux obtained by integrating the power law 4.0 x lO'1!-1'1 photons ca"2s~2Ma\r1, suggested by Walravan et al. (1975) at lower energies, over energiea from 3 to 25 MeV. Four points on the energy distribution of the total flux from the Crab Nebula, found In the same way, are given in tha energy distribution of Figure 3 along with reaulta from other observations. Our valuea are below those of Baker et al. (1973) and Gruber (1974) but above the upper limita of Schonfelder et al. (1975) In the caae of the assumed power law B"2-3. B. Pulsar NP 0532. The time diatrlbutlon of the number of events occurring as a function of the pulsar phase is given in Figure 4. Data from vertical cells only and for the time of 2.5 hours when NP 0532 la moat nearly overhead are used. When gamma-rays with angles between 15* end 25* only are retained, the signal to noise is Increased. The primary 20 I I I I I I I I I I I I I I I I I I I I I l II I l В— —I— PULSED— H В \l f S ScMWMv «01119791 Gnikv (I9T4) /Я ЧЛ Tnompaon « or. (ttrr) ь* Mondrao o> «1.(I97T) Fit of Wefrovon «t «I. 0973) 1' ' I • • I I•I I II Л10Г Ura. of ol. (1973) ' ' I III и! i i 11 iin.L i 11 mill. I I Mil < • r) if « «пгоггммазо» Ю 100 PHASE INTERVAL I mi ) GAMMA RAY ENERGY (MoVI FIG. 3—The energy distribution of tha FIG. 4—Number of events per total gamma-rays from the Crab Nebula. phase Interval for the pulsar Vertical cell pairs only.are used. Our NP 0532. Data vara used from results are compared with other obser vertical cell palra only for 2.5 vations. hours when NP 0532 was most nearly overhead and for incident gamma rays of 15* to 25*. 28 Tha «teas* counts above background In the region of and betwaen tha two peaks ia combined with tha total observation time and efficiency of tha telaacopa to give the time averaged flux of pulaed t -i-r-yi The back ground la determined from tha region В of Figure 5. Our reaulta and other observations of Pulsar MP 0532 are given In Figure S. They are In reasonable agreement with tha line 6.2 x 10-1,E~2'2 photona ст_1вес"1 MeV"1 auggeated by Walraven et al. (1975). Our flux Integrated from 3 to 30 MeV Is 4.312.6 x lO-* photons ст~2ввс~1, a factor of 3 higher than that obtained by Integrating the line 6.2 x 10-|,Г"2,г photons свГга-1 MaV-1 (Ualraven et al. 1975) over the ease energy Interval. The ratio of the total flux to the pulsed flux from 3 to 25 MeV la 11±8. We thank NASA, NSF and ONR for their support and NSBF for the launch, flight and recovery. FIG. 5—The energy distribution of pulaed gamma-rayе from NP 0532. Our point is compared to other observations. REFERENCES Albats, P., Frye, G. M., Jr., Zych, A. D., Mace, 0. B., Hopper, V. D., and Thomas, J. A.-1972, Nature 240, 221. Baker, R. E., Lovett, R. R., Oxford, K. J., and Ramsden, D. 1973, Mature • Phy». Scl.. 245. 18. Browning, R., Rameden, D., and Wright, P. J. 1971, Nature Phya. Scl.. 232, 99. Gruber, D. E. 1974, Proc. Int'l. Conf. on X-rays in Space 2. 875. Hayaca, R. C, Ellis", D. V., Fiahman, G. J., Kurfaa, J. D., and Tucker, W. H. 1968, A£. J. (Letters). 151. L9. . Herco, D., Koga, R., Millard, W. A., Moon, S., Ryan, J., Uilaon, R., Zych, A. D., and White, R. S. 1975, Hud. Inst. Methods. 123. 583. Kluxer, R. L., Share; G. H., and Seeaan, N. 1973, A£. ±,, 180. 547. Knlffen, D. A., Hartman, R. C, Thompson, D. J., Bignaml, G. P., and Fichtel, C. I.. 1974. Nature. 231. 397. Kurfess, J. D. 1971, A£. jTTtetters). 168. X39. Laros, J. C., Matteson, j7 L., and Pelling, Ri M..19ja, Katura Phy». Set.. 246. 109. • :ч Kandrou, P., Nlel, M., Vedrenne, C, , and Dupont, A, 197?, A£. J.. JJ2> ?04. McBreen, B., Ball, S. E. , Jr., Campbell, M., CreUen, K., and Koch, D. 197), A£. J_. , ljW. 571. Orwig, L. E., Chupp, E. L. , and Forrest, 0. J. 1971, Nature Phya. Sc 1 • . .'31, 171. Ryan, J. M., Dayton, B., Moon, S. H., Wilson, R. B., Zych, A. D., and White, R. S. 1977, 15th Int'I Cosnlc Ray Conf•. Paper OC-lb. SchOnfelder, V., Lichtl, G., and Moyano, C. 1975, Nature. 257, 375. Thonpaon, D. J., Fichtel, C. E., Hartman, R. C., Knlffen, D. A., and Lamb, R. C. 1977, Ap_. J_. , 212, 252- Walraven, G. D., Hall, R. D., Meegan, C. A., Coleman, P. L.. Shelton, D. H., and Haymes, R. C. 1975, Ap_. JL , 202, 502.- Zych, A. D., Herzo, D., Koga, R., Millard, W. A., Moon, S., Ryan, J., Wilson, R., White, R. S., and Dayton, B. 1975, February, IEEE Trans. Nuc. Scl., Vol. NS-22. 605. 30 COKPAUSQ* OF Ш1СЯ ЖКШЖОТ GAMMA tATS ГШОН |b | > 30° HITS TO GALACTIC WUTtAL ПОКОСШ DISTKnOTIOM H.E.flsel. H.Ogelman, T.TOmer Middle Eaat Technical Dnlvanity, Ankara, Turkey C.E.Mchtel, R.C.Hartman, D.A.Knlffen, D.J.Thompson HASA/Goddard Spree Flight Center, Greenbelt, Maryland, USA Abstract: '. High energy gamma ray (E > 35 MeV) data of SAS-2 satellite haa been used to compare the intensity distribution of gamma rays with that of neutral hydrogen (HI) density along the line of sight, at high galactic latitudes (|b| > 30°). A model has been constructed when the observed gammi ray intensity has been assumed to be the sum of a galactic component proportional to the HI distribution plus an Isotropic extragalactlc emission. The x2-te*t of the model parameters indicate that about 301 of the total high latitude amission may originate within the galaxy. I. Introduction Cosmic ray-galactic diffuse matter interactions are expected to contribute the high energy gamma ray flex at significant levels. For «ample, Schliekalaer and Thlelbelm (1976) calculates that the magnitude 4i of local galactic contribution to high lattltuda gemma ray amiss*— 1* C(> 40 MeV) - 21 ** X °* the total Intensity +21 (1> C(> 100 MeV) - 51 "• Z " " " " -14 With thaaa Ida*» In mind, a two-component model (galactic ami axtraaalaoHc) for the observed high lattltude gamma ray emission baa been developed and compared with the SAS-2 data. II. Method and Model The observed intensity variation of the high energy (> 35 MeV) gamma rays in 10° x 10° ( .030sr) equal solid angle bins is presented In Figure 2 in galactic coordinates, ( i , b ) . This map was obtained by asking use of the gamma ray data from the SAS-2 experiment whose detailed description is given in Derdeyn et al.,1972. In order to avoid as much tim possible the effects of the galactic plane which is an Intense gsmms-ray source, only the regions With |b| > 30° will be considered in the analysis. As the representative of the Batter distribution st |b| > 30° the N1 distribution surveys of Heiles (1975) were used. The galactic HI map used in this analysis is given ля Fig.)* This figure is divided into same equal solid angle bins as Fig.2 end HI-density for each bin is represented by the average value of four 5°x5° bins forming that 10°х1в° bin. The essence of the model fitting is a X2 test between the expected gamma ray intensity fif and the observed intensity gi, at the seme 10°xl0° bin, summed over the 81 available data bins. •In our model, the expected relative gamma ray intensity value, f^. 32 for each bin, la assumed to ba a aim of two tinu aa HI CR ft - C + K NJ Ni (2) where C stands for the laotroplc «xtragalactlc component conatant for HI all f-blna and the aecond tera repreaenta the galactic component; N. , CR N. being the HI ami cosmic ray (CR) denaitlea for that bin. Furthermore, it was aaaumed that the CR density to be proportional to HI density: if~<«f>e O) This assumption has as its basis the Idea that interstellar matter, csomic rays and magnetic fields are coupled Co each other forming a hydrostatic equilibrium under the galactic gravitational attraction (Parker,1966, 1969). Then, we can write, HI 1+a HI b fi - C + K(N^) - C + K(^ ) (4) The X2-test statistic is defined as 2 f bf*0' XZ - I 5 (5) i-1 Lfi where Af^ - (t^Js^y- and gi is the observed relative intensity.ej is the sensitivity fector of the i'th bin. One further condition on the statistic is that 28l - Zf± (6) That is the observed and expected total intensities are equal. This condition determines the constant K for a chosen set of C and b. The most probable values of the model parameters C and b for 33 different significance levels, can ba obtained In a two-dlaansloaal X2 plot In tha C-b plana aa given In Fig.3. A contour of significance level la defined In line with Lampton at al.(1976)t x2 < x2/-) - w+ fa 0> where x£(a) la the value of xZ diatribution for p-degreea of freedom, p bates the number of parametera adjusted In minimising Г (2 In our caae) and.la also tabulated In Lampton et al. (1976). A rejection of the model la recommende only If Х.Ш > 4-p (a " 102) (в> In our case, N • M-l since (6) has already been adjusted. III. Results From Fig.3 the following Inferences can be made: 2 (1) The pure extragalactic model corresponds to b • 0. The value of x for this case is X2(b - 0) - 112 which should be rejected on the baais of (8) since -the rejection limit la sig nificantly lower: 4*,-,. «-"»-» Several combinations of C and b have better fits than the b - 0 (C • 100 X) cam*. (2) The minimum contour corresponds to хъш"88 The 90 Z confidence limit is X2 (90 Z) - 93 There are a large number of Cb-comblnationa within theae liaita. Among these we should** further restrict ourselves to Ъ-лго1иев of b > 1, The cases with b 4 1 ' и doct net aeem to be physically meaningful since It requires a decrease In CX denalty for an increase In HI danalty. Such a ralatlonahlp la contrary to (*) and to the hydrostatic equilibrium arguments thought to be valid In the galactic dlak. This way, C ia further raatrlctad to a region whara C > 501 of the total high lattltude amission. As a plausible apaclal caa* for a uniform CR denalty in dependent of N*11 correapondlng to b - 1 the galactic contribution la: G<> 35 MeV) - (32 t 14) X of the total (10) at 90 Z confidence level. A alailar analyala for Е > 100 MeV gamma raya, leada to (Siel,1977) G(> 100 MeV) - (44 ± 16) X of the total (11) at, again, 90 Z confidence level. A uniform CR density around solar neighbourhood with |b| > 30° regions of the galactic disk is a plausible assumption since a relatively small region of the disk (about 10-5 of the volume of the disk) has to be considered. Then (10) and (11) can be taken as the magnitude of the galacti contribution to the high lattitude y-rays. These result are also in accordance with the theoretical estimations given in (1). IV. Conclusions and discussions: These results can be applied to the previously reported diffuse gamma ray intensity results of SAS-2 (Fichtel at al., 1975), giving I(> 35 MeV) - (5.7 ± 1.5)xl0~5 gamma raye/cm2-s-sr I(> 100 MeV) - (1.2 ± ,4)xl0-5 " " " " " for the extragalactic part of the observed diffuse emission. Its possible implications are discussed elsewhere (Ozel, 1977). V. References 1. Derdeyn, S.M., Ehrman, C.H., Fichtel, C.E.,Kniffen, D.A.,Ros8, R.W., 1972, Nucl.Ina. and Meth., 98, 557 . 2. Fichtel, C.E., Bartman, R.C., Kniffen, D.A., Thompson, D.J., Bignami, G.F., Ogelman, H., Ozel, M.B., Turner, T., 1975, Ap.J., 198. 163. 3. Heiles, C, 1975, Ast.Ap.Supp. И>. 37. 4. Lampton, M., Margon, B., Bowver, S., 1976, Ap.J., 208, 177. 5. Parker, E.N., 1966, Ap.J., 145, 811. » 6. Parker, E.N., 1969, Sp.Sc.Rav., 9_, 651. 7. Schlickeiaer,ltaad Thielheim, K., 1976, Nature, 261. 478. 8. Oxel, M.E., 1977, Ph.D.Thesia, METO, unpublished. IV •»«"^'«.e.,Krmm*m»mr,W.I..,1965,Sp.Sc.B.T.,i,123. 35 20 -2 Fig.l HI densities «long the line of eight,in units of 10 H-om ,in galac tic coordinates,for |b|ju.O (Heiles,1975KParts with broken lines are adapted from Garmire and Kraushaar, 196b» - o» - o. i* - V4 П « •- 3 N 0J Tip.2. Relative дашта ray intensity in galactic coordinates for |b|>30 as ob served by SAS-2.The values correspond to gj_ sCni/sjJxlOOj n^: the observed nucb;r of ijaiii'iA rays,s-J sensitivity in units of H'j>'60 s.(-l* stands for regions unobserved by SAS-2. 36 mx *£• . «- C»ftfltlt- P. Schlickeiser, K.O. Thielheim Instltut fiir Heine und Angewandte Kernphysik, Universitat Kiel, 2300 Kiel, Fed. Rep. of Germany Cosmic gamma ray production by bremsstralJ.ung in the interstellar gas involves cosmic ray electrons with energies below 1 GeV. In tiiis energy range the electron energy spectrun is rather poor known and quite uncertain. Vfe estimate the influence of this uncertainty to the gamma ray bremsstrah- lung productionrate. Especially it is shown that the integral production- rate above 100 MeV is uncertain to a factor 8. The consequences for the galactic gamma ray spectrun are discussed. 1. Introduction Measurements of the energy spectrum of high-energy (>50 MeV) diffuse gamma rays from the galactic center direction (Fichtel et al. 1972, 1975; Sood et al. 1975) have shown that this spectrun is softer than a pure 77 -spectrum and contains a component with an integral IL_-behaviour which probably originates in electromagnetic interactions of relativistic cosmic ray electrons (Stecker 1975), namely non-thermal bremsstrahlung in the inter stellar gas and inverse Compton scattering of low-energy photon gases like starlight and 2. 7°K blackbody radiation. In order to arive at well-founded conclusions on the spatial distribution of interstellar gas and cosmic rays which we expect from ganma ray astronomy it seems interesting to discuss the prcductionrates of these processes together with their accurateness. Stecker (1973) estimates the uncertainty in the integral TT-productionrate due to the experimental errors of the cross-section measurements and inter stellar cosmic ray spectra to be about £15 percent at 100 MeV photon energy. Bignami and Piccinotti (1976) have performed a detailed discussion of the scattered starlight contribution. Recently, Schlickeiser and Thielheim (1977) have shown that the uncertainties of the cosmic ray electron spectrum measurements above 100 GeV imply an additional uncertainty in the integral (>100 MeV) source function of inverse Compton scattered 2.7 K blackbody photons of at least a factor 6. In this work we discuss the remaining process, bremsstrahlung of cosmic ray electrons in the interstellar gas. 38 2. Nan-thermal brems itrahlung In the interstellar gas flw klnsmkclcs of this process have widely bean discussed In the literature mil—ileal and Qould 1970). Ooemic gamna ray production via bramsstrahlung Involve* m—lir ray electron» with energies below 1 GeV. Assuming an interatellar helium to hydrogen ahuwHanoa ratio of 10 peiuenl the Integral mn1 directional bnusstrahlung productionrata par equivalent hydrogen atom is given by oo oo У -f*c (1) where I (E) denotes the demodulated differential spectral flux of i rise! r ray electrons in the solar vicinity. In the relativJstic case E, E_» 0.511 MaV the differential cross-section is given by (Bethe and Heltler 1934) «W-tg-lU'b-irtH -&-Ш ш(2 ) where at is the fine-structure constant, r the classical electron radius and o are £| and Ф2 energy-dependent scattering functions determined by the piopax- ties of the involved targets. Blumenthal and Gould noted that for photon energies Ey.J <%* MeV - z is the atomic number of the target material - the strong-shielding limit is applicable which means that the scattering functions are nearly constant S 5 t i ^ (3) If the cosmic ray electron spectrum 1(E) - NE-p(p>1) is characterized by a single роиш: law we derive from (1) -tf.4) (5) a result given first by Stecker (1977). 39 In figure 1 ме ahov the orient knowledge at the lntanfaajlar ooaatc па/ electron apectrua below 10 0*/ derived free deanrtnlafrlnn of вааяшашг*» in the aolar vicinity and froa the —мtared frequantlal dlatributlon at galactic eynchrotran radiation, tte aae that with eMoaptlon of the upper lint given by Oaetinga at «1. all ргцюваД differential apactra aay ba ruaaaantad by a ГК.Е r Е «5c (6) In this caae the integral braaaatrahlung productionrate ia given by 4 sir * r/ ,.- (7) Stc for K Ъ h 40 I I МИЯ МТОТПЦХМ «РРОТШПМ. •ч I MMIII I" • чччччт чг" I I мщ •-• 1 •* V : t „«• пои мою DATA TZZ ми» ет - MLDSTON ет оемпоииьтсо: -- НИШИ ЕТ ALimi -* •'•••"• * ' ' Ч лют wis: 30 ЮО — COWSK AND wees i«*> FMOTON емеиог/i SHUKL* AND MUL <Ч711 30 «0 300 «D0 3000 ELECTRON vmmt /n*v I ч< * •* "'"цищ *» «t «i. tr» i.i i гл жоа I.H i тл» :ю • to** •* R«ii.im i.« > i.9 JOCD I.M J m (i-o-i.ti • «Г* *rt «bfM itH i.« i a,« JOOD О.Ш I МШ l.t» ' Ю"* nr««i. im i.i i J.* JCOD I.*» I »H.M 2.10 - Ю** «••ма im I.M t i.fi JOOO t.M i HW.X 1.И • »"* **> мнуш tea г UMC t •»«* *1 1.«.«г* 41 Jith (7) we calculate far the various proposed electron spectra the bremr- itxahlung pEOductionratM which are shown In the table. It oan be aeen that Леее results differ at least by a factor 8 depending on the choice of the ilectron spectrum parameters. i. Oansequsnces for galactic qasma ray astronomy Bignarai et al. (197S) and Stacker (1977) have pointed out that gamma ray production with energies grea- :er than 100 MeV In typical regions of the galactic disk Is dominated by Uan-decay and bremsstrahlung interactions. Only within the galactic nuclear region a considerable contribution from starlight scattering is expected uhereas microwave scattering becomes important at directions with higher jalactic latitude (Schlickeiser and Piielheim 1977). Since gama ray intensi ties from U^decay and bremsstrahlung are scaled by the sane line-of-sight integral assuming that the spatial cosmic ray electron density varies like the cosmic ray nucleon density the spectral behaviour of the integral galactic gamma ray intensity is propirt Irani to the sun of the integral picnic and bremsstrahlung productionrate which is shown in figure 2 for gamma ray ener gies between 70 and 2000 MeV. Ihe picnic productionrate is taken from the work of Stacker (1970) applying the * 15 percent uncertainty at all energies whereas the bremsstrahlung productionrate is calculated by using the lower and upper limit of the .electron spectrum of Cunnings et al. .At 100 MeV we find for the sun of the omnidirectional picnic and bremsstrahlung productionrate per equivalent hydrogen atom 4 Prior estimates of this quantity by Fichtel et al. (1976), Ramaty and Wester- gaard (1976) and Stacker (1977) lie within our error bars. As an important consequence of equation (в) we find that the conclusions on the interstellar matter distribution and origin of cosmic rays which are provided by garnna ray astronomy are uncertain to the same extent. These astrophysical important questions can be answered by gamma ray astronomy only then when our present knowledge of the productionrates is seriously improved. Our work has shown that much better knowledge of the interstellar cosmic ray electron spectrum is needed since this implies the major part of the un- 42 certainty. On the other hand it i» clear fro» figure 2 that gasae ray spectral ssaasurenants at energies lower than 100 MBV nay help to find the true Inter stellar coeelc ray electron epectniB. Actoowladgasants: One of us (R.S.) acknowledges a helpful discussion with Or. R.C. Hartnan on cosmic ray electron measurseents below 10 OeV. Numerical calculations were performed on the PEP 10 caspubsr of the Rschsn- zentrum of the University of Kiel. References: Bethe.H.A., Heitler,W. 1934, Prcc.Roy. Soc. (London) Л 146, 63 Bignami,G.F., Piccinotti,G. 1976, Astron.Astrophys. 52, 69 Bignami,G.F. et al. 1975, Astrophys.J. 199, 54 Blumenthal,G.R., Gould,R.J. 1970, Rev. of Modem Phys. 42, 237 Cowsik,R., Vbges,M. 1974, in Ohe Context And Statue of Gamna-ray Aatiunsy, ed. B.G. Taylor, ESRO SP-106, p. 229 Cumiings,A.C. et al. 1973, Prcc. 13th Int. Conf. Cosaic Rays(Denver),1^ 335 Daugnerty,J.K. et al. 1975, Astrophys. J. 198, 493 Fichtel.C.E. et al. 1972, Astrophys. J. 171., 31 Fichtel,C.E. et al. 1975, Astrophys. J. Jt98, 163 Fichtel,C.E. et al. 1976, Astrophys. J. 20B, 211 Goldate.tn,M.L. et al. 1970, Phys.Rev.Letters 24, 1193 Ranaty,R., Westergaard,N.Y. 1976, Astrophys. Space Sci. 45, 143 Schllckeiser,R., Thielheim,K.O. 1977, Astrophys. Space Sci., in press Shukla,P.G., Paul,J. 1976, Astrophys.J. 208, 893 Sood.R.K. et al. 1975, Prcc. 14th Int. Conf.Oosmic Rays (Munich) 1., 35 Stecker,F.W. 1970, Astrophys. Space Sci. 6_, 377 1973, Astrophys.J. .185, 499 1975, in Origin of Cosmic Rays, ed. J.L. Osborne and A.W. Wblfendale, Reidel, Dordrecht, p. 267 1977, Astrophys.J. 212, 60. 43 A TEST FOR ТЖ ORIGIN OF GALACTIC GNMt RAYS R. Schllctoi—r, X.O. Ihielhelm Inetitut fUr Rein* und Angowandt* Kkmphysik Uhiversltat Kiel, 2300 Kiel, Pad. Rap. of Germany Fran measured CQS-B gamna ray data it la poaaihle to derive the variation of the ratio of the galactic ganma ray flux in the two energy intervals 70-300 MeV and 300 - 2000 MeV as a function of galactic lati tude for two longitudinal regions. It is shown that this energy ratio is a good test between a more proton and a more electron induced origin of galactic ganma rays. 1. Introduction Diffuse galactic ganma radiation with energies greater than SO MeV originates in nuclear and electromagnetic interactions of primary cosmic radiation with interstellar matter and cosnic radiation fields: (i) the decay of neutral plans produced by inelastic encounters of cosmic ray protons and ос-particles with ambient Interstellar gas atoms and molecules, (11) non-thermal bremsstrahlung of relativistic cosmic ray electrons in the Interstellar gas, (ili) inverse Oampton scattering of low-energy photon gases like starlight and 2. 7°K microwave blackbody radiation by relativistic electrons. It is ccnuonly believed that more than 70 percent of galactic ganma radiation with energies greater than 100 MeV originates by the IT -decay (Stacker et al. 1974, Dodds et al. 1975). However, recently a nutter of authors have discussed the possibility, that the major part of high-energy (>50 MeV) gamna rays originates in electromagnetic interactions. OowsUc and Vbges (1974) have shown that with appropriate parameter values of the frequential and spatial starlight distribution inverse Oompton scattering of starlight photons explains, the major part of ganma radiation from the galactic center direction. Ramaty and Westergaard (1976) haw. shown on the basis of the closed galaxy propagation model of Raanuseen and Peters (1975) that the brerosstrahlung productlonrabe (E^.> 100 MeV) may be considerably greater than the corres ponding plan productionrate. In a recent work we have noted that due to the ,44 large acale height of the coamlc ray conflnaeant region (Joklpll 197C) Onwntnn •oetberlng of adcrowev* blackbody photona doalnabea at •»<••• and hlah latitude* (Schlickeisar and Thlelhmlai 1»77). Therefore It аша to be reasonable bo diecuM again «apiricel —thoda to dieeriaunate batman coamlc ray electron Induced цоиеаaaa (bremaatrahlung, dnverae .Ooaptaan scattering) and tha m—1r ray proton Induced ff*-decay. 2. Definition of tha energy ratio In the 006-B satellite «xperlaent the latitude profile of the galactic gauna ray aniaalon for the energy regions 70 - 2000 »fcV and 300 - 2000 MeV will be measured for two galactic longitu dinal intervals (Bennett et al. 1976): a) 350°< ln< 20°, hereafter called "center region", b) 244° < ln<284° excluding the contribution of the Vela point source, hereafter called "tangential region". Therefore we have investigated the ability of the following dimensionless energy ratio F(W-300fty;t') Ktb/" F(300-2W»*l'•,^I,) to discriminate between various possible galactic gamma ray origin processes. 3. Theoretical predictions Taking into account the gamma ray generating processes (1) - (ill) and assuming that the shape of the energy spectrum of primary cosmic radiation does not change throughout the galaxy the integral spectral flux of galactic ganma radiation from the direction (1 , b ) is given by (xvj1 sec" spv 45 where the integral productlonretes and source functions par sberadian f±[ >E-) «re taken fron the work of Stacker (1970) (1T"-dec*y), Blusenthal and Gould (1970) (DreBSJStrahlung, B), Bignaml and PiocinotU (1976) (starlight scattering, S) and Schlickeiser and Thielne!» (1977) (microwave scattering,»). ТЪе spatial dependent factors c^ (lU, b11) are determined by the following line-of-sight integrals: .. ' M C/wc <4 -a. <•", C/vvb (?>, cs (д. ,ь J — j«r ^ M CVi*. f6K O, is the density of the interstellar gas, w . denotes the energy density of starlight and N and N represent the cosmic ray densities of nucleons (cr) and electrons (e). №e subscript s denotes the corresponding quantities in the vicinity of the solar system. By averaging over a certain longitudinal interval {"/,< , tt J one obtains from (2) the latitudinal profile (7) Furthermore the final angular resolution of modem ganma ray detectors like the OOS-B one has to be taken into account. It has been pointed out by Pinkau (1973) that due to the use of spark chanters in experimental technique this angular resolution in general is energy-dependent. In our calculation c we have similar this effect by folding the laUtudinal profile P(>Ef; i ) with a Gauss function centered at b with a photon energy dependent standard deviation б* (F^): 46 4У - <»-/)' Ub br-c Bar o" (By.) MUM б*- 6° for Ey. *[70, 100 Hev] , 5° for *y «fu». ISO MsvJ 4° for Е ь [150, 300 Mev] , 3° for E-.6 [ЗОО, 500 Mev} , 2.5° for Г > 5O0 !•» since the accuracy of the angular resolution increases with Increasing gsam ray energy. By using equation (8) we have calculated the energy ratio R. (• ) tor the central and the tangential region for each of the four game ray generating processes ()X , B, S, M) separately. In ourreticulation we assume: I Die spatial distribution of the Interstellar gas is given by i of the 21-cm emission line of atonic hydrogen and the induced amission from carbon monoxide molecules from dense cool clouds (Burton and Gordon 1976). II The starlight photon energy density is proportional to the density of stellar matter. III Regarding the spatial • distribution of primary cosmic rays wa dlffnui two galactocentric cylinder synmetrlcal (-Г fe) models! concerning z: total thickness 2 Kpc concerning r: a) proportional to the distribution of supernova remnants (Kodaira 1974) b) constant. But, as may be seen from figure 1 and 2, results are rather Independent from both alternatives. 4. Results In figure 1 and 2 results are shown for the central and tangen tial region respectively. In the upper part for model (a), In the lower for model (b). In both cases we see that the energy ratio provides a good test 4? TNCCWETKAL ™ "••ЙД:ДЮ I I I I 2U* i i 1 I lUm*P<2U.m (Ь) =c^zp=!=H _l 1 Figure 2 -20 -IS -1'0 -5 10 . 15 20 b*(D«grml 4Я Ьацаип the nuclear (1Г) and the ali.tn:—I/MIIC (B,S,M) gamaa ray generating prooaaea•. Roughly spoken, a hard gaarae ray spectrum like tha picnic спа liarta to aawll energy ratloa whereas aoft apactxa like the electron Induced laply rather large ratloa. A further discrimination betiwen tha three electromagnetic processes seams to be not conclusive since tha differences between their energy ratloa are boo anmll. Therefore as the consequence of this work we propose to group empirical data in this way in order to get conclusive results an the origin of galactic ganma radiation in specific regions of the galaxy. Acknowladgeaents: The numerical calculations were carried out on the PEP 10 rqipiter of the University of Kiel Rechenzentrum. References: Bennett,K. et al. 1976, in: The Structure And Content Of Our Galaxy And Cosmic Bays, ed. C.E. Fichtel and F.W. Stacker, NASA X-662-76-154, p. 23 Bignami,G.F., Piccinotti,G. 1976, Astron. Astxophys. 52, 69 Blunenthal,G.R., Gould,R.J. 1970, Rev. of Modern Phys. 42, 237 Burton,W.B., Gordon,M.A. 1976, Astxophys, J. 207, L 189 Oowsik,R., Voges,M. 1974, in: The Context and Status of Ganna-Ray Astronomy, ed. B.G. Taylor, ESBD SP-106, p. 229 Dcdds,D. et al. 1975, J. Phys. A 8, 624 Jokipii,J.R. 1976, Astrophys. J.. 208, 900 Kcdaira,K. 1974, Publ. Astron. See. Japan 26, 255 Pinkau,K. 1973, in: Proc. of the Int. Synp. and Workshop en Ganma-Ray Astrophysics, ed. F.W. Stecker and J.I.Tzonbka, NASA X-641-73-180 Ramaty,R., Westergaard,N.J. 1976, Astrophys. Space Sci. 45, 143 Rasnussen,I.L., Peters,B. 1975, Nature 258, 412 Schlickeiser,R., Thielheim,K.O. 1977, Astrophys. Space Sci., in press Stecker,F.W. 1970, Astrophys. Space Sci. 6, 377 Stecker,F.W. et al. 1975, Astrophys.J. 188, L 59 •«> GALACTIC KAYS AMD THE ORIGIN OF COSMIC RAITS D. Dodds and A.W. Vfolfendals, Physics Department, University of Durham, England, and E.C.M. Young, Physic» Department, University of Hong Kong. A critical appraisal ic made of the dittribution of neutral and molecular hydrogen in the Galaxy, both towards and away from the Galactic Centre and the data are used to make a detailed analysis of the distribution of cosmic ray nuclei in the Galaxy needed to explain the observed gamma ray flux. The inferred cosmic ray distribution is compared with that of pulsars and SNR. The effect of varying the magnitudes of the less veil known parameters is examined. 1. -.Introduction. The problem of determining the distribution of low energy cosmic rays in the Galaxy from measurements of the distribution of Galactic y-rays is a continuing one. A number of problems arise: the target material in the I.S.M. with which the C.R.s collide to produce the parent «"-mesons is somewhat uncertain in mass and position, largely due to problems with the distribution of molecular hydrogen; there are difficulties with the penetration of the dense molecular clouds in which so many ©f the Hj molecules reside and, as a general point, the cosmic rays may be confined to regions not occupied by much molecular hydrogen. Further difficulties arise concerning the contribution to the measured Y-ray flux from discrete Sources, such as pulsars, and from interactions of the electron component. In what follows the various problems are considered. 2. Gamma Ray Emissivity. A prerequisite for derivation of the C.R. intensity (by which is meant the intensity of the nuclear component) is the emissivity. Strong and Worrall (1976) have derived the relative emissivity of y-srays above 1СЮ MeV, W(R), from the intensity v» longitude results given by Fichtel et al. (1975) and this hss been binned in double cells by us with the result shown in Figure 1. It is to be noted that the region beyond 10 kpc is taken as the datum. 3. Distribution of Target Material. A number of analyses of neutral hydrogen have been made and two of molecular hydrogen,- the latter being by Scoville and Solomon (1974, 1975,referred to as SS) and by Gordon and Burton (1976, referred to as GB). There is a difference in absolute magnitude by a factor of ^ 2 and both workers agree that this it probably the order of accuracy at the present time (1976, private communication). The potential importance of molecular hydrogen is shown by the fact so 20 " 1 IS <> «ft) 10 : : 5 - 1 - ik12 —i—i—i—i—i—i—1_ Figure 1. Relative y-ray emissivity (Ey > 100 MeV) as a function of galactocentric radius after Strong and Worrall (1976). that the surface density of gas is increased over the value for neutral hydrogen alone by the following factors: 5, 5, 3, 1.7 and 1.2 at R » 4, 6, 8, !10 and 12 kpc respectively (R is the Galactocentric radius). The values are from GB. It should be stressed that, whereas the H2 inhabits rare dense clouds,the data presented have been smoothed and it is usual to assume radial symmetry. He ignore these problems and allow the CR to 'see' all the gas, in the initial analysis. 4. Contribution from non-CR. Pulsars come into this category because the emission of y-rays is presumably by way of electrons accelerated in the pulsars. Some pulsars (Crab, Vela, FSR 1818-04 and PSR 1747-46) have been recognised as y-ray sources and their contributions subtracted from the measured intensity. Following the work of Dahanayake et al. (1977), which suggests that the pulsar contribution is < 10Z, this component is ignored. It remains to be seen whether the estimate is valid although, since the form of W(R) -for y-rays from pulsars is probably similar to 51 И(М Figure 2. Ratio of CR Intensity at distance R, I (it) to the local value, 1(10). The surface densities of SNR and Pulsars, o(R), are necessarily very approximate: SNR from Pimley (1975) Pulsars from Lyne (1976, private communication). the total (see later), the presence of a quite appreciable component might not be too serious. A number of estimates of the relative contributions from electrons (principally via bremsstrahlung) and П°- mesons, have been made. We adopt the figure of 2031 (for y-rays above 100 MeV) given by Dodds et al. (1976) although attention should be drawn to the latest value of 38Z given by Worrall (197?). In.deriving the distribution of CR from the emiesivity divided 6y mas') density allowance must be made for the contribution of secondary electrons as distinct from primaries. Whereas the primary electron contribution is a constant fraction of the total the secondary contribution increases relatively as the CR intensity itself. For an increase in cosmic ray intensity over the local value by a factor of 3 we estimate that this factor should be reduced by 7Z to allow for secondaries. 5. Derivation of the R-dependence of the CR intensity. It is assumed, that there is radial symmetry in the distribution of cosmic ray intensity, 52 Figure 3. Ratio of CR intensity at 6 kpc from the G.C., 1(6) to the local intensity 1(10) as a function of the fraction of H2 seen by £ 3 GeV C.R. The H2 data of GB were adopted. although likely perturbations should be borne in mind. Figure 2 shows the ratio of the CR intensity at distance R, X(R) to that near the earth, 1(10). The emissivity of Figure 1 has been divided by the l.S.M. mass distribution given by GB and the yield factor of Dodds et al. (1976) has been adopted. A correction has been applied for electrons, both primary and secondary (see §4). The value of I(12)/I(10) indicated in Figure 2 comes from the analysis of Dodds et al. (1975) which relates specifically to the general direction of the anticentre. A more recent analysis by Strong et al. (1977) using the COS В data indicates an even lower CR intensity at R>10 kpc; in fact, there appears to be virtually no evidence for 1-10 GeV C.R. in and beyond the Perseus arm. 6. Comparison with the distribution of SNR and Pulsars. SNR and pulsars are likely CR sources although flare stars too may contribute significantly at energies near 1 GeV. Figure 2 shows best estimate lines for the distribution of SNR and Pulsars. It will be appreciated that the uncertainties on these linas are large, typically of the same order as the S3 uncertainties on I(R)/I(10). Within the present errors the distributions are all rather similar giving COM support to tba suggestion that coamic ray» in tha range 1-10 G«V originate largely in evolved objects. A necessary corrollary ia that rhey do not diffusa far fro» thair point of origin (typically { 2 ape). Before developing this idea further it ia necessary to examine tha sensitivity of tha predictions to uncertainties in the input data. 7. Sensitivity of results. Tha problem* of tha absolute magnitude of the n°2 danaity and molecular cloud panatration were mentioned in ll. Figure 3 shows tha ratio of I(6)/l(l)> to tha fraction f of H2 'seen' by the Ct. As reaarked earlier, the analysis has used tha G.i. H2 densities; if the SS results are adopted tha ratio falls and tha datum value of f, f,, at which 1(6)/1(10) falls to unity (i.e. tha situation for no CK gradient) ia reduced from 4 ± 1.5 to 2 ± 0.8. Bearing in mind the convents about tha uncertainty in H2 densities it ia aean that data towarda the G.C. could conceivably be consistent with no gradient at all. Turning to the penetration probability, Skilling and Strong (1976) estimate that, at 3 GeV, exclusion starts in regions having t> 3 z tha local CR energy density at cloud densities of mere than about 3 z the average density. Thus, we would expect only a snail reduction in y-ray flux. However, uncertainties in the cloud properties and in the model predictions raise doubts and it would not be surprising if f were smaller than unity and 1(6)/1(10) correspondingly higher. Another factor leading to the same trend is that some clouds may well be in interara regions where the CR intensity might be smaller than in tha arms themselves, The y-ray flux in the A.C. direction still offers the best opportunity to determine the cosmic ray gradient. Here H2 densities are small and tha H distribution is known rather well. More precise experimental y-ray data would be useful. 8. Conclusions. On balance, there is evidence for a cosmic ray gradient in the Galaxy such as would arise from 1-10 GeV particles having been produced in evolved objects (and perhaps, such objects as novae and flare stars at these low energies) with attendant small diffusion. Tha outlook for models involving easy diffussion into a Galactic halo does not appear to be good. Acknowledgements. The authors are grateful to Miss D.M. Worrall, Professor C. Dahanayake, Dr. J.L. Osborne, Dr. J. Skilling and Dr. A.W. Strong for useful discussions. We are indebted to Pr fee'ore W.B. Burton, N.C. Scoviile, P.M. Solomon and Dr. A.G. Lyne for helpful correspondence. References. Dahanayake, C, Dodds, D. and Wolfendale, A.W., 1976, y-ray Symposium, Goddard S.F.C., 126, IX-662-76-154). Dodds, D., Strong, A.W., and WoJ.femiale, A.W., 1975, Hon. Hot. K. Astr. Soc, 171, 569. 54 Dodda, D., Wovcayk, J. and Uolfemdala, A.W., 1*7*. Мм. Not. K. Aatr. See., 17», 345. Fichtel. C.I. at al., 1»75. Ay. J., IN, 163. Cordon, M.A., and Burton, W.B., 1976, Ap. J., 176. 597; 20», 346. rinlay, K., 1975, H. Sc. thaaia. University of Durham. Scoville, H.Z., and Solomon, P.M., 1974, Ap. J., Lett. 117, L67 (and private coammication). Skilling, J., and Strong, A.W., 1976, Aatron. and A»trophya., S3. 253. Strong, A.V., Bennett, X., Will», K.D., and Wolfendale, A.V., 1977, 12th ESLAB Synpoaium, Fraacati (in the preaa). Strong, A.W., and Horrall, D.M., 1976, J. Phya. A., 9. 823. Worrell, D.M., 1977, Ph.D. theaia, Univeraity of Durham. / ss GAMMA-RAY PRODUCTION IN DENSE MOUCUT-AR CLOUDS Л. W. Strong Department of Physics, University of Durham, Durham City, England. J. Shilling Department of Applied Mathematics and Theoretical Physics, Silver Street, Cambridge, England. The movement of primary and secondary electrons into and "dut of dense molecular clouds is inhibited by resonant Alfven waves generated by cosmic ray nuclei and also by the cosmic ray electrons themselves. The waves cause the electron density within a cloud to diminish, and this causes a corresponding reduction in the expected flux of bremsstrahlung gamma-rays from clouds. He estimate the reduction for plausible parameters. 1. Introduction Skilling and Strong (1976), hereafter referred to as S a S, showed that Alfven waves can be sustained by cosmic- ray nucleons at the boundaries of dense molecular clouds, as a result of the aniaotropy set up by energy losses of the cosmic rays across the clouds. The effect becomes Important for nucleons of energy below ~ 1 GeV, so that it is only of marginal importance for pion-decay gamma-ray production. We now extend the theory to a mixture of cosmic-ray species, including electrons. We consider particularly the relatively low energy («j 100 n^V) electrons which are expected to be the major source of low-energy (~ 20 MeV) gamma rays in the Galaxy (Flchtel et aL 1976). Any inhibition of the propagation of external electrons into, and internal electrons out of, the clouds will modify the expected bremsstrahlung emission. 2. Theory Our analysis and notation follow S * S, but now we have three distinct species of energetic particles, namely electrons (e~ and e+), protons and helium nuclei. As before, we require a 1:1 correspondence between wavelength 2 if/k and particle rigidity R. - cps/ I zs«l (suffix s labels the different species, s = е for electrons, s > p for protons, в - * for helium). In the presence of any density gradients, each species contributes independently to the intensity growth rate of waves propagating into the cloud. For each rigidity value, the corresponding waves then quickly reach an equilibrium intensity 3 3 ГИД i''''11" X"; (l) at which damping, probably due to ion-neutral friction, at a rate T balances the growth rate, I0(IC) = particle spectrum outside (inside) the cloud. The effect of the waves is to cause a flux 56 (3) of particles to enter etch of the two aides of the cloud (cf. equation В of S ft S). Here x - Coaptoa-Gettlng factor and 0 - 4ри-/ЗГ»По4 - spatial diffusion coefficient. In a steady state, conservation of particles gives *^(eJ«E - -Л/(ЕО;П4,(Е)Д(Е)) - Л^(Е)ЛЕ (3) for a cloud of thickness Л atoms св~ within which particles lose energy in a mean free path of X, atoms cm'2. This is equation - (13) of S I S with an extra source term qs(E)d£ particles atom ! s~l to take account of secondary electron production within the cloud. Eliminating ja and Os yiejds «,1„ ± + JbL fr(le»-W + _A^ -^i(Uci) (4) for the response of the particle flux density Is to a self- consistent level of waves. The singularity. In (4) when j -* 0 and correspondingly Zp'v do-Ic) -» ° ™y be circumvented by noticing that the flux n js(E) In (2) may not exceed we(no8~ ce)/ I (e) - l-s-* Ю"Г (E/l&ev)" c«"V sr"' MeV (6) with У* 1.Я below lGeV and У= 2.5 above 1 GeV (spectrum (a) of Goldstein et al. 1968). Electrons suffer ionization and brems- strahlung losses, which we take to be continuous, giving a 1 1 1 resultant mean free path Хв = (Хд.е~ + Abe" )" Ginzburg (1969) quotes 5? U Xu - |.4S-.IO /C4i(^/i%e^ + OJt] c«* (e) for ionization and bremsstrahlung losses. In this paper, we assuM that clouds contain 34 helium nucleons for «vary 100 hydrogen atoas (Allan 1973), so we correspondingly increase the ionization loss rate by a factor i.17 and the bremsstrahlung loss rate by 1.51. Plausible electron synchrotron losses have negligible effect on the results. Secondary production spectra for e~ and е in the general interstellar medium are given by Ramaty (1974). However, he took the local cosmic ray spectrum for the producing nucleons, whereas these may well be partially excluded from the interior of dense clouds. Since the dominant production mode of secondary electrons, ir'±decay, is numerically proportional to the decays which produce Y-rays of a few hundred MeV, we scaled the Ramaty spectra down in the same ratio that У-ray emission from the cloud is reduced. „ m.-.4...\ pica У -flux from cloud with waves «e = qe(R«Baty). £г5а у -flux from could without waves (9) The standard external proton and helium spectra were taken from Garcia-Munoz et al. (1975)f I - 4-1 * io* (е + 7*о «^(- г.гхю-^е ))'хле ^"VVM.V"' 5 г Гвл = l-lxio (o-XS E + U0 «^(->Г*1о~*Е)Г*'™«'" *"'$г*'м«/'<10> where Е is in MeV/nucleus. Nuclei suffer energy losses from ionization according to and from inelastic interactions according to Г х Vr ' K* - t-sx ю* ш- (12) Within the clouds, we set the nuclear source terms to zero. (13) The cloud thickness Л'was assigned the "standard" value 4.2 x 1022 cm'2. With the input data complete, the three coupled equations (4) or (5) (for e, p, *) were solved numerically. 58 Fig. I shows the proton spectrum In»Id* the cloud - the helium spectrum Is similar. These nuclear spectre show • pro gressive attenuation inside the cloud below about 600 MeV/nucleon, ss the Alfven waves switch on and become progressively more important. Fig. 1 Proton spectra outside and Fig. 2 Electron spectra inside a standard cloud. outside and inside a standard cloud. Fig. 2 shows the corresponding electron spectra inside the cloud. The electrons too are progressively and significantly attenuated towards lower energies due to the Alfven waves. The energy (rigidity) at which waves switch on is determined almost entirely by the protons, and hardly at all by the less powerful electrons, although at lower energies the actual intensity of the waves becomes largely determined by electrons. 4. Bremsatrahlunq radiation from clouds The bremsstrahlung emissivity for photons of energy Ey from electrons of intensity Ie(Ee) Is proportional to J* Ie(Ee)dEe (Stecker 1971). The reduction in bremsstrahlung flux because of the waves is therefore given by the ratio This is shown as a function of energy in Fig. 3. $9 100 Ky (KeV) Fig. 3 Reduction factor for bremsstrahlung gamma rays of energy Ey . For 20 MeV photons, the reduction is by a factor 0.5 for standard parameters (and by yet more significant factors if the cloud thickness is increased). This is the energy region for which Fichtel et al. (1976) have made predictions for the diatri- bution of Galactic У-rays, and the effect of the waves will clearly be considerably to modify their predictions. A proper evaluation of the effect will require a knowledge of the distribution of column densities for the clouds. If we assume standard parameters everywhere, then the Fichtel et al. bremsstrahlung flux should be reduced by at least a factor 0.6 for \t\ < 30° (the region where molecular clouds dominate the interstellar matter distribution). This underestimates the effect because Fichtel et al. assume a higher electron intensity in denser regions, so that the reduction will be further enhanced. 5. Conclusion We have shown that cosmic ray electrons will be excluded from molecular clouds, along with cosmic ray protons and heavier nuclei, by the action of instability-generated Alfven waves. This has the effect of keeping energetic electrons away from a large fraction of the interstellar material, thus decreasing the intensity of bremsstrahlung У-rays of around 20 MeV which will be produced within the Galaxy. The magnitude of this reduction in У-rays almost certainly exceeds a factor of 2 for emission from \l\ < 30°. Acknowledgements AWS thanks the Royal Commission for the Exhibition of 1851 for the provision of a Research Fellowship. D. M. Worrall is thanked for useful conmients. References Allen, C. W. 1973 "Astrophysical Quantities" Athlone Press. 60 Daugharty, J. K., Bartman, R. C. and Schmidt, P. J. 1975 Aatrophya. J. ,198 493. Flchtal, C. Ж., Knifxwn, 0. A. , Thompson, 0. <)., H71UI, C. F. and Cheung, C. У. 1976 Aatrophya. J. 20_l 211. Garcia-Munot, H., Kaacn, G. N. and Slapaon, J. A. 1975 Aatrophya. J. 202 265. Glnzburg, V. L. 1969 "Elementary Proceaaea for Coaaxlc-Ray A*trophyalea" Gordon 4 Braach. Goldatein, M. L. , Ranaty, R. .and Flak, L. A. 1970 Phya. Rev. Lett. 2± 1193. Ramaty, R. 1974 "High Energy Particlea and Quanta In Aatrophyaica" ed. McDonald and Fichtel, MIT Preae. Skilllng, J. and Strong, A. W. 1976 Aatron. Aa trophy a. 5_3 253. Stecker, F. W. 1971 "Cosmic Gamma-Raya" NASA. 6QWM&j 6i COSMIC RAY PENETRATION INTO MOLECULAR CLOUDS by C.J. Casaraky* and H.J. Volk** Centra d'Etudes Nucleaires da Saclay, B.P.2, ^+9U90 Cif-Sur-Yvette, Franca. Mex-Planck-Institut ftlr Kernphysik, 69 Haidalbarj, Pottfach I039BO, Germany. Assuming denae intaratellar clouda to b« thraadad by tha general galactic magnatic fiald, tha rola of magnetic fields in reducing the interior cosmic ray intensity ia inveatigated. Tb« diffuaive effects of small acale irregularities inaide and outside the cloud aa we]1 aa of large acale hydrosagnetic wavei are found to be negligible down to energiei of aoae tens of MeV/nucleon. There fore quasistatic magnetic mirrors which may arise at local conden sations can trap particles over time scales long enough so that the trapped cosmic ray population is absent due to losses. This ener gy independent effect can lead to an average reduction of у~с*У luminosity and hydrogen ionisation rate of tbout 50 per cent. At the same time density contrasts in the gas are considerably empha sized in the Y~light. I. Introduction. This work represents a detailed discussion of the effects of magnetic fields on the interior cosmic ray intensity and therefore the ionisa tion rate and the y-ray emission of typical molecular clouds. The galactic magnetic field is assumed to thread the cloud, being compressed by the con traction of interstellar material to actual cloud densities. This field is expected to be quite inhomogeneous. Its effects on cosmic rays of a given energy depend strongly on the spatial and temporal scales of the inhomogenei- ties and may comprise adiabatic energy changes, resonant diffusion inside and outside the cloud, nonresonant diffusion by random mirror fields, and trap ping in quasi-steady large scale mirror type equilibrium configurations. Such effects could prolong the particle residence time in the cloud and thus en hance collisional losses and result in a depressed internal intensity. Des pite the expected inhomogeneity of the field it appears meaningful to intro duce an average magnetic field Вд. Thin implies that the actual field 13 is not too severely tangled. Such a spaghetti plate type structure is indeed quite unlikely: As pointed out by Heiles (1976), the infrared and optical po larisation are observed to be large, which implies that the fields are rather more uniform than tangled. Secondly, the compression of a cloud with tangled fields would not allow preferential flow down the field lines but would ra ther lead to ah isotropic compression such that B^n2'3 (n - number density of the gas). The upper limits В < 50 uG reported for two clouds by Crutcher et al. (1975), however, appear to exclude such high fields except .perhaps for isolated small regions which are not our concern here. For definiteness we consider a standard cloud with „ ,-3 „ -3 „ ,_-3 „ -3 n, - 2.10 .N cm , n.» 2.10 .N. cm , , a. o ' i io В - 50BQUG,M • 10 .КД , H • 8.2.R pc, where n„ - 2^ + n^t is the hydro gen number density given by that of H2 and HI, n.is the ion number <" plasma) density, M is the cloud mass and R its radius.For the standard cloud N • N. » В -M «R -1. The lowest cosmic ray energies considered in i> IO O O O JO this paper are 10 MeV/nucleon. 62 1. Adiabatic «ю—ntua change, resonant and nonresonant diffusion in the cloud. Except (or possible •••11 Fermi acceleration effects in iaall seal* aagnetic coapressions, the cosaic ray energy in • cloud can be regarded as being un affected by the magnetic field. This follows because the free fall {iaa l3 ,/2 tf{ - (32 Свн^/З *)~ - 3,6.10 Ho~ sec is large compared to the largest collisional aoaentua loss lis* t. - 6,5.10 N sec in the cloud appropriate for relativiatic nuclei with an aaauaed range of 66 g ca (e.g. Nakano and Tadeaaru, 1972) for protona; here G and a„ denote the gravitational constant and the hydrogen mass, reapectively. Small scale random magnetic irregularitiaa on the Laraor radius acale L « 6,4.10 SY/B cm can resonantly scatter particles in pitch angle ( Be and у are the particle velocity and Lorents factor, reapectively). If the resulting diffusive mean free path A for notion along ^B is short enough so that the effective residence time t. is of the order of the loss tire t., then energy losses will be significantly enhanced and such particles will be severely depleted in the cloud. This can be true although-The unperturbed range of the particles may be large corresponding to the straight line column density 0,16 N R g cm-2 represented by the cloud. In this sense, resonant dif fusion in the cloud is important if , tD - t. , where tD - 6R A 8c This requires a certain level of fluctuations. From magnetic scattering theory (e.g. Hasselmann and Wibbcrenz, 1968) one roughly obtains,except for factors of 0(1): _ Ф x - rT A_ . „ L f * P 00 wnere L8,rW W 8" r, w L W is the required magnetic energy density in resonant fluctuation* with wave number k and power spectrum P(k). For relativistic particles we then get W - 3,5.10 6yB R N erg cm , whereas for nonrelativistic particles W w is smaller to reach W ww•* .1 0 W at 10 MeV/nucleon. These resonan• t waves are decoupled from the neutral gas for all particle energies of interest and are damped by friction with the neutrals.Using the Kulsrud and Pearce (1969) damping coefficient generalized to an arbitrary mass m.for the ion dominating frictional losses, we can calculate the total power input into waves of such scales that is necessary to offset their damping: dW 36 w . 4n D3 1,1.10 „ „ _ -1 ST ~ R я /- , <* BoRo erg sec The above rate is enormous. This is easily; appreciated if one compares it to the maximum rate 3GM /5Rt _ - l,8.1036 M2 N'^R-1 erg sec"1 of gravita tional energy release in free fall collapse of the cloud. What other sources of wive enery could be relevant? One of the most important should be newly formed expanding HII regions within the cloud. Following an estimate by Aarons and Max (1975), for the case of an 05 star we find for the maximum power input P(HII) into hydromagnetic waves P(HII) - 1,8.1033 В 5/3 N~7/6 erg sec"1 . 63 If we reaeaber that the resonant wavai in question cannot b« directly excited by motion* of tha neutral gas, than the snergy aust soathow be trantferred to scales of 0(r, ) down froa tha acalaa at which turbulence ii excited in the neutral taa. Assuaing gravitational inhoaogeneities at tha icala of I pc and, rather arbitrarily but conservatively, the relevant scale* of a I pc newly formed HI1 region to be about I0"2 pc due to preexisting inhoaogeneitie* of the gaa, then wa have to reduce theae acalea by factor* of 3.I0'BQ and 3.10°Bo, respectively, to reach the Larmor radius scale of a 10 MeV proton. 2 To sake the best case for turbulent tranafer, we take Wu"\.k~'/ correspon ding to the inertial range of a Kraichnan (1965) spectrua. Even in this case the maximum possible power input is about a factor of Ю smaller than re quired by wave dissipation and thus insufficient to drive significant cosmic ray diffusion. Nonresonant diffusion in the cloud by random mirroring can come about for par - ticles with large enough pitch angles if the correlation time tc ъ (kc.vA) ' of the random mirrors is smaller than the transit time 2R/Bc of a particle through the cloud.From this condition it follows first of all that waves with these scales have the same damping decrement as the shorter scale resonant waves. Secondly, if one requires, that at least 10 per cent of all particles should be affected by these waves on average, then one obtains (W /8irB2 ) > 10~' and thus W > 10~1A.B2 erg cm"3 independent of particle energy (Volk. 1975). The required power input is about 2.102.RQ NQBQ/YB times larger than the maximum rate of gravitational enery release. Thus, in contrast to the supposition of Hakano and Tademaru (1972; random mirroring is negligible too in typical clouds. 3. Cosmic ray induced diffusion outside the cloud.Recently, Skilling and Strong (1976) have pointed out that the loss of cosmic rays in dense clouds leads to a cosmic ray streaming into the cloud which, by arguments of con tinuity, may excite hydromagnetic waves in the ambient intercloud (IC) medium along magnetic flux tubes that connect onto the cloud. He point out here that the actual streaming speed v|c in the IC medium is by a factor BIC/B smaller than the streaming speed in the cloud, because this factor represents the ra IC IC tio of flux tubes areas. Taking B - 3 uG we obtain B /B - 0,06/Bo for the standard cloud. Assuming for the sake of argument the instability not to ope rate we assume an approximately uniform cosmic ray distribution N(p) every where. Then we have +R В 2 v {"•v-r -£ f 4- (Sf)coll -R where (dt/.Л ,, is the momentum loss rate due to two body collisions at J coll (e.g. Montmerle, 1976). Extrapolation to low energies the demodulated spec trum of Morfill et al. (1976) that is well approximated by a power law N^p-2,5 for proton energies £ 200 MeV, using the following IC parameters'. ngC - 0,2 cm-3, TIC - 104K, n*C - 0,03 cm"3, BIC - 3 uG, (Kulsrud and Cesarsky, 1971) and the dispersion relation of Wentzel (1968), the insta bility behaviour is shown in Fig. 11 Particles with kinetic energies \iD 2 corresponding to f(E^in) > 2,28.10" RN are unaffected by this mechanism. The constant 2,2e.l0~2RN is proportional to the cloud column density.For our stan dard cloud the energy of marginal stability is about 50 MeV/nucleon. Thus at least down to these energies cosmic rays are not selflimiting their access 64 to polecular cloud* as considered bar*. Siaple arguaanti (how that «uch a stataaerit should be even «ore true If th* cloud was labedded in a hot "coronal" aadi.ua. A. Cosaic ray trappini in in ternal airrora.Froa the re i(i-J sults of the previous sec tions it ia raaaonabla to look at the opposite case of a acatterfree situation. Than particles can be trapped i*> («.it*») in quasistatic equilibrium configurations. Due to two body energy losses the trapped particle population will ef fectively be absent. Repopula- tion would only be possible if the resonant mean free path Л would be shorter than the mirror size 1. Already for 1 mirrors we thall call "bottle*" of unperturbed density p0. Assuming the mass in each consecutive pair of ». • mirror and bottle of Fig. 2 to be equal, we can limply calculate the average depression and local con trast of the cloud's y-ray lumi nosity or volume ionization C»n««niali«A relative to the case of an un modified cosmic ray intensity. For the example given in Fig. 2, the average depression is about a factor of 2, whereas the local Fig.3 relative contrast between adjacent mirror and bottle regions is a factor of 3 to 4. Both effects might be observable in the y'light. 5. Conclusions.Molecular clouds, as defined by our standard parameters exhi- bit no additional screening of cosmic rays due to irregular magnetic fields down to about 50 MeV/nucleon. For this reason, large scale internal mirror fields lead to average depressions of y-ray luminosity and primary ionisation rate of about 50 per cent and to a roughly 3 to 4 fold increase of relative contrasts of both quantities. The result suggests in particular a hydrogen ionisation rate of about 10-17 sec-' in dense molecular clouds. The calculation neglects local sources of cosmic rays and effects due to secondary electrons. To attempt a-theoretical prediction of the effect of local sources does not appear very fruitful %t this stage. Secondary elec trons should not play a larger role at y-energies > 100 MeV than else where in the galaxy (Badhwar and Stephens, 1976; Shukla and Cesarsky, this conference). References Aarons.J., and Max, С.ЕГ., 1975, Ap. J. 196, L77 Badhwar,G.D.and Stephens, S. A., 1976, Phys. Rev. D JU, 356 Crutcher, R.M., Evans, N.J. Ill, Troland, T., Heiles, C, Ap. J. J98, 91 Hasselmann, K., and Wibberenz, 6., 1968, Z. Geophysik 34, 353 Heiles, C, 1976, Ann. Rev. Astron. Astrophys. _Ь4_, 1 Kraichnan, R., 1965, Phys. Fluids, 8, 1385 Kulsrud, R., and Cesarsky, C.J., 1971, Astrophys. Lett, j5, 189 Kulsrud, R., and Pearce, W,, 1969, Ap. J. 156, 445 Montmerle, Th., 1977, to appear in Astron. Astrophys. Morfill, G.E., VSlk, H.J., and Lee, M.A., 1976, J. Geophys. Res. j», 5841 Nakano, T., and Tademaru, E., 1972, Ap.J. 173, 87 Parker, E.N., 1963, "Interplanetary Dynamic Processes", Interscience Publ., New York, p. 165 * v Shukla, P.G., and Cesarsky, C.J., 1977, this conference Skilling, J., and Strong, A.N., 1977, Astron. Astrophys., to be published Volk, H.J., 1975,-Rev. Geophys. Space Phys. J^, 547 66 пщ згвсам or ееиргаивитшшо ашегж >чиаа г ». 0—1** ааа «. Tafaa* • Tat* laatltata af Гаа >•••>>! 1*м«к, а«ааа/-5, labia Мах FLaaak tatltat far axtrafcwrraatrlMba Pb/alk, **» Oaraaia*, k. Made*, ПО. Tbaoratleal D • D Tbaeratleal «ark oa taa arltlaa af galaetle gaaaa-ra/a Ьм ladieatad that tba g—• raya prodaoad by IIWH-Coaptaa aeattarla*; «Г atar-U«at fora aa lapcrtaat ooapcaaat at taa Пих free taa galaotle plana, though the axaet fraotlee of thla ooapoaaat la la oaaatian. lb an attaapt to elarlfy thla qaaatloa m praaaat bar* a oaraful oaleala- tioa of tbair apaetral abapa. At flrat, «a critically rariew both taa theoretical and axperlaental etadiee oa tba apaotraa of etar-light photon* and that of tba eoeaio-ray elactrone la tba Qalaxy. Wa tbaa «aa tba exact croea-aaetlona for Coapton-aeatterlag to oaloolate tba apaetraa of gaaaa-raya prodaead la their lataraetloaa. 9a dtaeaaa tale apeotraa la the context of tba SiS-ZI obaarratioBa of tba falaetle dlae. Co-ordiaateai CO 1.1 Mailing Addraaai Or. B. Canals Tata Xaatitate of Fandaaantal fteaearoh Baal Bhabha Koad Boabay 400005, bdla 67 BREMSSTRAHUJNG GAMMA RADIATION FROM THE OAUXY P.O. Shukla and C.J. Ceearakv Service d'Electroniqu* Physique, Centra d'Elude* NucMairea d* Saclay, France We preaent a detailed itudy of the bremaatrahlung gamma ray emiaaion from the galactic diak. We •how that there are large uncertainties in the production apectrum of photona of energy 10 to 70 MeV, due to our lack of knowledge of the interstellar electron apectrum. We discuss what can be expected, and what can be learnt, from future observations of gamma ray* in that energy range. I. INTRODUCTION The galactic gamma ray emission through electron bremastrahlung ha* been studied earlier by several author* but the most recent investigation* have been due to Fichtel et al. (1976) and Kniffen et al. (1976). A general conclusion from these latter studies is that while beyond ~ 70 MeV the emission mechanism through the generation and decay of neutral pions is dominant, at lower energies it is the electron bremsstrahlung that contribute* most of the gamma ray* created in the interstellar space. We have carried out a detailed •tudy of the gamma-ray emiasion by this mechanism and have in particular considered the influence of the uncertainties in our know ledge of the cosmic ray electron apectrum. In view of the possibility in the near future that the COS-B satellite data will provide at least a rough shape of the interstellar gamma-ray energy spectrum, we consider it important t» point out these uncertainties which have not been explicitly considered before. II. CALCULATION OF THE BREMSSTRAHLUNG EMISSIVITY IN THE SOLAR NEIGHBOURHOOD We first evaluate the gamma-ray emisaivity for regions of the interstellar medium which are close to the sun and tor which we have some knowledge of the cosmic ray spectrum, flux and composition. -) Electron spectrum. For our purpose, we are interested in the electron (negatron + positron) spectrum above 10 MeV. The h*gh energy part of the spectrum is directly accessible to observations from the earth and is essentially unaffected by solar modulation ; a glance at some recent review paper* (Daniel and Stephens». 1975, Lund 1975) shows that the election flux at 10 GeV is known -with an error of ~ 50%. Between 200 MeV and a few GeV, the interstellar spectrum can be derived from the galactic non-thermal radio spectrum (Goldetein et al. 1970). 6S Cummings rl al. (1971a) show that, while th« shape of the spectrum -a power law of Index 1.76- ia quite well determined, it» amplitude ia uncertain by a factor of ~ 4 above ~ 300 MeV, due to uncertaintiea of the interstellar parameter* involved. Direct fieaaure- menti of the ncgatron spectrum in this energy rr.nge are useful mostly to make a comparison with the interstellar spectrum and extract some knowledge about the solar modulation parameters. At lower energies, 10 MeV to 200 MeV, the radio results quickly become unreliable ui account of the heavy absorption of radio signals by unknown amounts of ionised gas along the line of sight. One possible way of determining the interstellar electron spectrum In this range of energies is to compare the observed positron spectrum at the earth, with calculations of the same in the interstellar medium. This comparison permits to evaluate the solar modulation parameters ; applying these to the observed negatron spectrum, one can, in princi ple, derive the interstellar negatron spectrum. Cummings et al. (1973b) apply this method to the measurements of Beuermann et al. (1969), and conclude that the interstellar negatroii spectrum must flatten consider»-- bly below 100 MeV. A similar analysis was carried out by Daniel and Stephens (1975) arriving at somewhat different results. In a more recent paper on this subject, Hartmann and Pellerin (1976) emphasize the difficulty of measuring the positron spectrum in the tens of MeV range in the earth's residual atmosphere in high altitude balloon flights. They consider their own results from two flights in 1974, as well as those of Daugherty et al. (1974) and of Beujermann et al. (1970) to represent only upper limits to the actual fluxes. between 10 and 50 MeV. If this is the case, Cummings et al. (1973b) have underestimated the effects of solar modulation, and the interstellar electron spectrum at energy < 180 MeV is above the one they derived. On the other hand, a fraction of the negatrons detected could have been accelerated locally, for instance in the Jovian magnetosphere (Teegarden et al. 1974) ; this would have an opposite effect on the derived electron spectrum. In summary, it is clear that the spectrum of interetellar electrons below 180 MeV is not well known. Hence, to study the bremsstrahlung emissivity, we have decided to consider several possible electron spectra, which are represented in fig. 1. They are : . ST, "Standard" spectrum, being the nominal interstellar spectrum of Cummings et al. (1973a), a power law of index 1.8, which has been linearly extrapolated to low energies up to 10 MeV. This spectrum is in good agreement with that given by Daugherty et al. (1974) beyond 50 MeV. Previous workers (Kniffen et al. 1976, Stecker 1976) have considered spectra of this kind only ; . CSV, being the nominal spectrum of Cummings et al. (1973b) with a turning over at low energies ; LB, UB, being the lower and upper bounds of the spectrum of Cummings et al. (1973a) based on radio continuum data, which have been extrapolated to low energies. As discussed above, it is difficult to assess the magnitude of the error of the spectrum at,1 say, 50 MeV; 6? DS, demodulated electron ipcitrum from Daniel «nil Stephen* (1975). b) Brcmsstrahlung production formulae. We hive used lorcmi- lae from Gould (19fc9) and Blumenthal and Could (1970) (or the mntri- bulion of interttellar H and Hc Komi to the bremastrahlung rmixian . for heavier «pecie», we applied the formula given by Gmiburn (19ъ9). Delaila will be given elsewhere (Shukla and Ceaaraky 1977). 111. GAMMA RAY EMISSION FROM THE ANTICENTER DIRECTION The bremsstrahlung differential production spectra in a medium with hydrogen denaity equal to 1 cm-' for the different electron spectra of fig. J, are displayed in fig. 2. In the same figure, we show *he local emissivity due to IT* decay, as calculated by Badhwar anu Stephens (1976). We assigned an error of 30% to it, due essentially to uncertainties in the demodulated cosmic ray nuclei spectrum. We also show the local inverse Compton emissivity for the "standard" electron spectrum of fig. 1 and for a light energy density as in the solar neighbourhood. We see that the inverse Compton emission is negligible with respect to the bremsstrahlung and the тт° emissions ; it is even less important if the mean local hydrogen densi ty is higher than 1 ст'Л In fig. 3, we represented the integral produc tion spectra, and in fig. 4 the total spectrum (Bremsstrahlung + TT° + inverse 'Compton) for some of the cases discussed. We note that, at 20 MeV for instance, the differential brems strahlung emiesivities predicted by our upper (UB) and lower (JLB) limits, extrapolated from Cummings et al. (1973a), differ by a factor of order 70. We see in fig. 4 that for the former, the "n° bump" in the y-ray spectrum disappears. The difference in the predictions of the spectrum CSV, with a bump at ~ 60 MeV, and the standard spectrum ST, is of only a factor ~ 0.6 at 20 MeV. The predictions of the standard spectrum are significantly different from those of DS : at 20 MeV, the difference is by a factor 4.5. At present, there are no reliable data concerning the galactic flux of gamma rays below 70 MeV but it is to be hoped that future experiments will provide this information. In the longitude range 90°< еП<270° and at low latitudes, the gamma »ays received (after exclu ding known compact sources) come mostly from closeby regions (see discussion in Cesafsky et al. 1977). For this case, the spectra of cosmic ray electrons and nuclei should be quite similar to those utili zed to calculate the emissivities given in fig. 2, and thus the gamma- ray spectrum received from the outer galaxy should resemble one of the spectra given in fig. 4. Gamma-ray observations in the range of energies 10-70 MeV could provide an important constraint to determine the interstellar low energy electron spectrum. On account of the convolution entering in ti . calculation of bremsstrahlung emissivity, the sensitivity of this approach is limited but with careful measurements this information might be useful for eolar modulation research. 70 n-J ,-1 r.1 E, (6»V) Е Ю"' „,-1 wu ю1 V XT •> tf Er (QtV) ;E, (GSV) Fit. 4. Total diffsrsntial production rates for three of ths electron spectra Fig. 3. Integral production rates, for ths of fig. 1 : differential rates of fig. 2. Tl • Bl + n + C (upper bound spectrum)' T2 » B2 + ff + C (standard spsctrum) T3 • B4 + и + C (lower bound spsctrum) 71 In the anti-center direction, it low latitude*, we expect the ga.mma.-ra.ye created by the inverse Compton (I, C. ) process to play a relatively unimportant role. On the other hand, if the acale heights of coimic ray* and star light were much larger than that of the gae, this contribution could become progressively more significant at higher latitudes. If the cosmic ray scale height is as given by Baldwin (1976) and the light distribution as given by Shukla and Paul (1976) we would find that at a latitude of 90*, the ratio of the inverse Compton contri bution to that due to bremeatrahlung and n* decay increases by ~ 1Л% (using the "standard" electron spectrum). If there is a galactic halo of cosmic rays, then the i. C. contribution would be even higher ; according to Worral and Strong (1976), uncer certain conditions, it could be responsible for most of the diffuse high latitude gamma rays at 100 MeV. In this case, it will be difficult to separate the galactic and extragalactic components at high latitudes. IV. CONCLUSION What will we be able to learn from future measurements of the galactic gamma-ray spectrum in the 10-100 MeV region ? a) A determination of the gamma-ray spectrum in the 10 to 100 MeV regions towards the anticenter region will ihed some light on the interstellar electron spectrum in the solar vicinity. If the low energy electrons received at the earth are of interstellar rather than interplanetary origin we will have an additional test for solar modula tion theories. b) At the galactic center : Several authors have pointed out that the starlight energy density at the galactic center is so high that the inverse Compton emission could be dominant there (Cowaik and Voges 1974), especially if the cosmic ray scale height is large (Stecfcer 1976). In fact, if the starlight distribution is as given by Shukla and Paul (1976), its density above the galactic center decreases quite rapid ly as the height increases. Consequently, increasing the scale height of cosmic ray electrons does not have a strong effect. In particular, if we adopt a large scale height as proposed by Baldwin (1976) rather than the lowe.r scale height of Paul et al. (1976), and the electron density on the galactic plane of Paul et al. (1976), we find that the inverse Compton emission from a beam of 5* centered at the galactic center, is increased by only 23% ; the total emission is then increased by ~ 2%. It mayi be, of course, that the starlight distribution in the galactic center is not as assumed by Shukla and Paul (1976) or that the electron to proton ratio in the cosmic rays is,much higher there than elsewhere. In that case, inverse Compton emission due to inter actions of electrons with starlight could become important. This hypo thesis can be- tested if data on the high energy (E > 1 GeV) part of the spectrum becomes available as n* and bremeatrahlung spectra would be power laws, while an inverse Compton spectrum with starlight dominating largely over the blackbody background would decrease' steep- 72 ly beyond 1 G*V (the onaet of the decay of the starlight inverse Compton apectrum can be a«en In fig. J). c) On the whole, we expect that the gamma-ray spectra from different direction* in the galactic plane will be similar, but who know* what aurpriae t* in atore 7 If spectra from different direction* did turn out to be different it might mean that point aource* in the inner galaxy contribute a larger fraction of the radiation , although with th* data available at pre*ent, this hypothesi» aeem* unlikely (tee Ce*ar*ky et al. 1977). If the diffuae radiation u largely dominant every where, then Kniffen et al. (1976) point out that the ratio 1(10-10 MeV)/ I(>100 MeV) it a meaiure of the ratio of electron* Ne to nuclei Np in cosmic ray*. Thi* i* only true if the interstellar electron apectrum i* well under our upper bound (UB) ; also we note that the information acquired would concern the ratio Ne(~ 30 MeV)/N_(few GeV) and that it cannot be extrapolated in a atraightforward manner to Ne(few GeV)/ Np(few GeV) if only because, of ionization losses. This i* too bad, because a comparison between gamma ray flux and synchrotron emis sion can yield some knowledge of the interstellar magnetic field (Paul et al. 1976) ; however, the relevant radio data concern electrons of a few GeV, the lower energy one* radiating at frequencies for which the interstellar gaa ia opaque. So, one can compare the high energy gamma- ray flux (>100 MeV) to the radio data aa in Paul et al. (1976) and make assumptions on the ratio Ne(few GeV/N_(few GeV). Or, when it become* available, one can u*e the low energy flux (10-30 MeV), but then it will be necessary to make some gueaae* concerning the electron spec trum in different parts of the galaxy. We thank Michel Ca**e\ Jean-Paul Meyer, Jacques Paul and Thierry Montmerle for discussion*. REFERENCES Badhwar.G.S. .Stephens, S. A. 1976. Phys. Rev. D»ugherty, J. K. et al. 1975, Ap.J. 196,493. D14. 356. Fichtel, C.E. et al. 1976, Ap, J, 208. 211. Badhwar.G.S., Stephen», S. A. 1976, Preprint Glniburg, V. L. 1969. Elementary processes NASA. for cosmic ray astrophysics, (Naw-Уогк : Baldwin. J.D.I 976, NASA Preprint X-662-76- Gordon & Breach). 154, p. 206. Goldstein, M. L. et al. 1970, Phys. R»v. Lett. Beuermann, K. P. et al. 1969. Phy». Rev. Lett. 24., 1193. 22,412. Gould, R.J. 1969. Phv». Rev. 185,72. Beuermann, K. P. 1970, Acta Phys. Acad. Sci. Hartman, R. C., Pellerin, C. J. 1976. Ap. J. 204. Hungaricae,2_S, euppl. I,p. 173. 927. Beuermann, K. P. 1 974. Proc. 9th Ealab Symp. Kniffen, D. A. et al. 1976, NASA Preprint p.259. X-66Z-76-154, p. 341. Blumenthal, G. R. .Gould, R.J. 1970, Rev. Mod. Lund, N. 1975, Proc. ICRC(M„nich), U., 3746. Phys. 42,237. Paul, J. etal.1976, Ap, J. 207, 62. Cowsik, R. , Viget, W. 1974, Proc. 9th Eslab Shukla. P.G., Paul, J. 1976. Astr. J. 208. 893, Symposium, p. 229. Shukla, P. G. . Cesareky, C. J. 1977. submitted Cummings, A. C. et al.l973a, Proc. 13th to Astron. and Aatrophys. ICRC, Denver, l_, 335. Stecker, F. W. 1976. NASA Preprint X-662-76- Cummings.A.C. 1973b, Proc. 13th ICRC 154, p. 357. Denver, 1^ 340. Teegarden, B. J. et al. 1974, J. Geophya. Cesarsky.C. J. et al. 1977 Astr. Ap., in Res. 73,4909. press. Worral, D. M. , Strong, A. W. 1977. Astr. Daniel, R. R. , Stephens, S. A. 1975, Space Sci. Ap., in press. Rev. 17,45. ?J ATNOSPHEUC GAIttt-RAT ANCU UD DOTtCT DlSmWJTIOHS ГМИ I TO IS NaV Лама N. «у an. Shin a. Noon, Robert ». Wlleoe. Alia* D. Zyca. t. St «Има waits Phyalca Department aad laatltuta of Gaopayalca aaa* Planetary Payalea Ualveralty of Caltforala, Rlverelde, California 92)21, U.S.A. aad Bruca Daytoa Phyalca Dapt., California Stata Untvaralty, Loa Angeles, CA 900)2, U.S.A. Abatract. Angle and energy distributions ara reported for gamaa-raye near the top of the ataoaphere for energlea of 2 to 25 MaV and for tealth eagle* of 0* to 50* (downward moving) and 180* to 130* (upward moving). Theee distribu tions were obtained with the UCI double Coapton acatter gaaaa-геу telescope flown on a balloon to 3.0 g/ca2 realdual ataoaphere froa Palaatlaa, Texas oa May 13, 1975 (Geomagnetic cutoff 4.5 GV). Growth curvea froa 3.4 to 100 g/ca* of realdual ataoaphere are uaad to determine the downward aovlng ataoapaarlc fluzea. Reaulta are given for 6 energy intervale froa 2 to 25 MeV and for 5 angle intervala froa 0* to 50*. 1. OBSERVATIONS. The gaaaa-ray obaervatlona were aade with the Univeraity of California, Riverside, UCR, double Coapton acatter teleacope (Herzo at al., 1975; Zych et al., 1975). The teleacope waa carried to 3.0 g/ca2 realdual ataoaphere by a 0.54 z 106 a3 balloon launched froa Palestine, Texas on Hay 13, 1975. Downward moving gaaaa-rays were obaerved during aacaot and at float and upward moving gaaaa-rays were observed primarily during ascent. Details of the telescope and launch may be found in Hllaon et al. <1977), paper OG-7. The flux for a particular energy Е and zenith angle в is calculated directly from the counting rate using the efficiency for detection of vertical cell type events. The detection efficiency has been calculated analytically and by Monte Carlo aethods and has been measured by callbrationa with gamma- raya of known energies. Discussion of the calculatlone aay be found in Ryan et al. (1977). The telescope was calibrated at the Van de Graaff accelerator at Call fornla State University,' Los Angeles (CSULA) in March and July of 1976. The detector was exposed to gamma-ray beams of 4.4, 6.1 and 12.1 MeV. Figure 1 shows measured angle and energy distributions of the detector for 6.1 MeV gamma-rays at various incident angles. The energy resolu tion HWHM varies continuously from 24Z at an incident angle of 10* to 5Z at 62*. The angular resolution HWHM varies from 8* at 10* to 12.5* at 62*. Figure 2 gives the ana lytical and Monte Carlo efficiencies FIG. 1—Measured angle and energy resolutions for vertical cell pairs with 6.1 MeV incident gamma-rays. Curves are drawn through the points 2 4 6 6 0 30 «0 <0 to guide the eye. MEASURED ENERGY VEWT1CAL CEU. «UKS OML> •r rf- I.I ** 6.1 M Й» RESIDUAL ATMOSPHERE (gm/cm4 12.1 i FIG. 3—Typical flight observation growth curves for the dowm.ard moving gaama-rays at energy intervals of 2-3, 5-7.5 and 10- D Monte Corto Colculalion 15 MeV for the telescope selected inci a Analytical Calculation dent angle of 25*. The solid line in • Mtoturtd each case is the least squares fit for • I • I i \ the atmospheric component. INCIOENT ANGLE tdtgritil along with measured efficiencies from the FIG. 2—Comparisons, of absolute CSULA calibration as a function of inci analytical and Monte Carlo dent angle °or incident energies of 4.4, calculated efficiencies with 6.1 and 12.1 MeV. A more complete de measured efficiencies for verti scription of the detector calibration will cal cell pairs as a function of be published later. It can be seen that the Incident gamma-ray angles. the analytical and Monte Carlo calcula The Incident gamma-ray energies tions qnd CSULA calibrations are in good are 4.4, 6.1 and 12.1 MeV. agreement up to the maximum calibration energy of 12.1 MeV. For the atmospheric gamma-ray analysis, the Monte Carlo efficiencies are used up to 15 MeV and the analytical efficiencies for 15 to 25 MeV. 2. RESULTS. A. Angular Distributions, (a) Downward Moving Gamma-rays. The downward gamma-ray flux near the top of the atmosphere consists of atmospher ic, extraterrestrial and neutron produced gamma-rays. Typical growth curves from 3.5 to 90 g/cm2 are shown in Figure 3 for the zenith angle of 20* to 30* and energy intervale of 2-3, 5-7.5 and 10-15 MeV. 75 TAILS I. DOWNVAKD MOVING CANNA-kAYS IN UNITS OF (Ю"1" photoae/cm'-eec-ater-MvV-d/em3)) Energy Zenith Angle В (MeV) 5* 15* 25* 35* 45* 2-3 6.9 1 1.3 14 i 1 3-5 6.4 i 0.9 6.3 l 0.5 6.5 s 0.5 9.5 l 1.0 5-7.5 6.7 t 5.3 3.8 l 0.4 3.6 l 0.4 4.8 1 0.6 e.o ; i.o 7.5-10 2.7 1 0.4 2.0 1 0.4 1.8 1 0.8 1.9 1 0.4 3.5 t 0.6 10-15 1.0 l 0.2 1.2 t 0.2 l.S t 0.3 2.9 0.5 15-25 0.5 1 0.1 0.7 ± 0.1 1.7 i 0.2 t TABLE 2. UPWARD MOVING GAMMA-RAYS AT 4.2 (/cm2 IN UNITS OP (10-3 photons/cm2-sec-ster-MeV) Energy Zenith Angle в (MeV) 175* 165* 155* 145* 135* 3-5 11 ± 4 18 ± 3 22 ± 3 5-7.5 8.6 i 1.5 10+2 16 ± 3 31 ± 7 7.5-10 1.4 ± 0.9 6.1 ± 2.0 6.7 ± 2.0 8.6 ± 3.0 13 ± 5 10-15 6.4 + 3.4 4.5 ± 2.0 4.0 ± 1.4 4.5 ± 1.9 5.4 ± 2.2 15-25 1.5 ± 0.4 0,8 ± 0.2 1.7 ± 0.4 2.1 ± 0.5 2.8 ± 0.8 We represent the total downward gamma-ray flux, T, as (хДсовв> T « a + bx + c в- where T is the total measured gamma-ray flux, a Is the altitude-independent gamma-ray background Induced by neutron interactions in the liquid scintilla tor near the top of the atmosphere (White and Sc'.Bnfelder, 1975), bx is the atmospheric flux and c is the extraterrestrial flux. The quantity x la the atmospheric depth, and Л is the interaction mean free path in the atmosphere for gamma-rays of energy E. The quantities a and c have units of (photons/ cm2-s-ster-MeV) and b has the units of (photons/cmz-s-eter-MeV-(g/cm2)). In principle a, b, c and X could be evaluated from the data. However, when the data are divided into many energy and angle bins, the statistics are not suf ficiently good to evaluate all four constants. We prefer to calculate a by the method of White and SchOnfelder (1975), use the total gamma-ray inter action cross section in the atmosphere to evaluate X and obtain b and c by the X2 fit to the data. The resultant values for the atmospheric gamma-ray fluxes b are given in Table 1. The cosmic diffuse fluxes c and the calculated neutron induced fluxes a and the details of their calculation appear in White et al. (1977), paper OG-20. Data from atmospheric depths of 10 to 100 g/cm2 contrib ute most to the determination of the atmospheric gamma-ray fluxes. Here, the cosmic diffuse and neutron induced ga-jma-ray fluxes are almost negligible compared to atmospheric gamma-ray fluxes, and their uncertainties contribute only slightly to the overall atmosfAieric gamma-ray flux uncertainties. (b) Upward Moving Genoa-rays. The upward moving fluxes of gamma-rays, at a residual atmosphere of 4.2 g/cn2, are given in Table 2. The contribution of the neutron produced gamma-rays is much less than 10Z. The UCR observed upward gamoa-ray fluxes, integrated over angles from 170* to 140* and energies of 3 to 25 MeV, are consistent with a constant flux of (1.65 + 0.06) x 10"1 photons/cm2-s, Independent of depth from 3.6 to 90 g/cm2 of residual atmosphere. Above 30 MeV, both the measurements and the calculations 6f Thompson (1974) show a gradual decrease with increase or depth in the atmosphere. His Monte Carlo model follows the 3-dimensional cascades 76 40 MO ZENITH ANGLE («4. 60 » 110 ZENITH UtdEttttJ FIG. 4—The measured angular distribution FIG. 5—The measured atmos- of the atmospheric gaaaa-rays at the pheric gamma-ray angular residual atmospheric depth of 4.2 g/ca2. distribution at 2.S g/ca2 residual atmosphere. The that produce the gaaaa-rays from the theoretical values of Graser initiating radiation. Ling's (1975) semi- and SchSnfelder (1977) are empirical theory for 3 to 10 MeV gives a given for energies of 3-10 gradual increase of the flux by a factor HeV and 10-30 MeV and the of 1.6 for an Increase of residual ataos seal-empirical values of phere from 3 to 100 g/cm2. Ling are given for a residual ataosphere of 2.2 g/ca2 and The downward and upward fluxes are energies of 3-10 MeV. plotted in Figure 4 for the altitude of 4.2 g/ca2 residual ataosphere. At' 3.5 g/ca2 and energies of 4 and 8 MeV, our downward fluxes are 3 and 8 times those of the semi-empirical model of Ling (1975). Deeper in the atmosphere the agreement is somewhat better. Extrapolating our atmospheric fluxes to 2.5 g/ca2, in Figure 5 we find upward fluxes over the energy range of 3 to 25 MeV at 15°, 25* and 35° about a factor of 3 larger than those predicted by Graser and SchBnfelder (1977). The ratios of the UCR upward fluxes to those of Ling (.1975) in the energy range of 3-10 MeV are .about 0.5 and 0.9 at 165* and 140*, respectively. While absolute values of the., fluxes are in reasonable agreeaent, our fluxes increase away from the vertical contrary to the decrease predicted by his model. Ratios of UCR upward fluxes to the upward theoretical fluxes of Graser and SchBnfelder (1977) are 2.1 and 3.6, respectively, at these angles in the same 3-10 MeV range. The ratios of our measured upward to downward fluxes, extrapolated to 2.5 g/cm2 residual atmosphere, are 12+1 and 9 + 2 in the energy intervals of 3-10 and 10-25 MeV, respectively. The theoretical ratios of Graser and SchOn- felder (1977) are about 15 at these energy intervals and those of Ling (1975) are 70 and 300 at 5 and 10 MeV, respectively, while Schdnfelder et al. (1976) give an experimental ratio of about 9 for energies of 1.5-10 MeV. Our gamma- ray atmospheric angular distribution is consistent with the distributions of Llchti et al. (1975) and Schonfelder et al. (1976) which peak near the horizon and decrease at larger and smaller zenith angles. B. Energy Distributions, (a) Downward Moving Gamma-rays. For comparison with other data, the downward moving atmospheric gamma-rays are averaged over 77 ""'"гтттт—i i r u"i i t i m»^ ^""4 5 1r (|T0*-I40>I l^f _ ~k ] 14 i + if v^ ~Ь KJa > • "•* _i i i i nil 10 Ю0 O Km» gf «L (1974) FIG. 7—The measured energy distribu fVMHm «.urn» 0 OoMtach* «t zenith angles from 10s to 30* where the angular distribution is reasonably flat. The resultant distribution, along with the results of other authors, is plotted in Figure 6. Our fluxes bridge the gap in the previous data ^ • between 10 and 20 MeV and are in reasonable agreement with the observation» both below 10 MeV and above 20 MeV. The theoretical curve'of Beuermann ° (1971) is a factor of 3 lower than our experimental results. Ling's «al- empirical calculations are a factor of 2 lower than the UCR results at the lowest energies and about a factor of 10 lower at 10, MeV. The theory of Daniel and Stephens (1974) is a factor of 3 below the UCR fluxes at 20 MeV and a factor of 20 below at 2.5 MeV. (b) Upward Moving Gamma-rays. The energy distribution of upward moving gamma-rays, as derived from the fluxes of 170* to 140*, is given in Figure 7. Also included are the power law fits of Imhof et al. (1976) at the equator and the pole of 4 x 10-2 Е-1*39 гиЛ 17 x 10""2 «Гьз* photons/cm2-e-ster-ltoV, respectively, at lower energies. They deduced the fluxes of the upward moving gamma-rays of 40 keV to 2.7 MeV with a germanium detector on a satel lite in a polar orbit, assuming the value of Pal (1973) of 1.25 x 10"2 (E/MeV)~2-i photons/cm2-s-ster-MeV for the cosmic diffuse flux.. Our values extrapolated to 0 g/cm2 are essentially equal to our values at 4.2 g/bm2. From Figure 7, our fluxes at the geomagnetic cutoff of 4.5 GV are a factor of 3.5 times their equatorial values extrapolated to higher energies and a factor of 0 7 times their polar values.' This agreement le considered reasonable. We thank NASA, NSF and ONR for their support and NSBF for the launch, flight and recovery. 78 REFERENCES Baueraann, K.P.. J. Geophya. Re». 76. 4291, 1971. Dahlbacka, G.H., P.S. Pr«tar and C.J. Waddington. Лр. J. 180. 171, 117). Denial, R.R. and S.A. Stephen», Ray. Caophya. Spac» Pnya. 12. 233, 197*. Graaar, U. and V. SchPnieldar, J. Caophya. Raa. 82. 1055. 1977. Herzo, D., R. Koga, W.A. Millard, S. Noon. J. Ryan. R. Ullaoo. A.D. Zych and R.S. White, Nuclear Inetrumante and Hathoda 123. 583, 1975. Iahof, W.L.. G.H. Nakano and J.B. Reagan, J. Caophya. Raa. 81. 2835, 1976. Klnzer, C.E., G.H. Share and N. Seaaan, J. Caophya. Raa. 79. «567, 197*. Lichti, C., C. Moyano and V. Schunfelder, Proc. 14th Int'l Conf. Co—.lc Ray. 1423. 1975. Ling, —«c; C, J. Ceophya. Raa. 80, 32*1, 1975. Pal, Y., IAU Synpoaiua, No. 55, X- and Caaau-ray Aatronoay, adltad by H. Bradc and R. Glacconl, Reidel, Dordecht, 279, 1973. Ryan, J.N., B. Dayton, S.H. Moon, R.B. Hllaon, A.D. Zych and R.S. Whit», to be publlahed J. Geophya. Re»., 1977. SchBnfelder, V., U. Graser and J. Daugherty, submitted to Aatrophyalcal Journal. 1976. SchOnfelder, V. and G. Lichti, J. Geophya. Res. 80, 3681, 1975. Staib, J.A., G.M. Frye, Jr. and A.D. Zych, J. Geophys. Rea. 79, 929, 1974. Thompson, D,J., J. Geophys. Res. 79, 1309, 1974. Valdez, J.V., P.S. Freler and C.J. Waddington, Acta Phye. Suppl. No. 2<», 1, 79, 1970. White, R.S., B. Dayton, S. Moon, J. Ryan, R. Wilson and A. Zych, paper OG-20 1977. White, R.S. and V. SchBnfelder, Astrophys and Space Sci. 38, 19, 1975. Wilson, R.B., B. Dayton, S.H. Moon, J.M. Ryan, A.D. Zych and R.S. White, paper OG-7, 1977. Zych, A.D., D. Herzo, R. Koga, W.A. Millard, S. Moon, J. Ryan, R. Wilson, R.S. White and B. Dayton, IEEE Trans. Nuc. Sci., Vol. NS-22, 605, 1975. 79 ROCKET ALTITUDE ATMOSPHERIC X-RAYI PROM MAONBTOBPHBRIC ELECTRON! AT SUBAURORAL LATITUDES J. G. Luhmam J.B.BU» Space Sciences Laboratory The Ivan A. Catting Laboratories The Aaroapaoa Corporation Loa Angeles, California 00001, USA Like auroral electrone, quaaitrapped magnstospherlc elae- trona mirrorinf in the upper atmosphere at low and middle latitudea will generate x-rays by the brenuatralunf and Ka line excitation processes. Theae atmospheric x-rays may contribute a diffuse background to rocket-borne astronomy experiments launched from White Sands, New Mexico and Kauai, Hawaii. Calculations of the atmo spheric x-ray spectrum at these launch sites baaed on observations of quasitrapped electrons indicate that this locally generated flux is comparable to the reported soft x-ray flux below a few keV. INTRODUCTiON Some atmospheric x-rays are generated in the electromagnetic cascades produced by cosmic rays in the atmosphere (Beuermann, 1971). This well-known component has a well-defined linear dependence on atmospher'-! depth above about 50 km (Imhof etaL, 1971). Data obtained during the rocket ascent gives the proportionality constant between dentil and the cosmic ray related atmospheric flux (Horstman, 1973). An additional but. lesser known atmospheric component results from x-rays generated by magnetospheric electrons. This component has a different altitude dependence than the cosmic-ray ч generated flux because it is not a cascade product, but the direct result of the interaction between magnetospheric electrons and the upper atmosphere. Cne important feature of this magnetosphere related, atmospheric component Is that the originating electrons are 80 not necessarily present at tht rocket whan tha «-ray» are praaant bsaauss «И of Им alactroM oan have thalr mirror points above tha rocket altitude. Figure 1 tUwatratea achamatieally how magnatoapheric electrons oan travel on trajectories that pass near and above tha racket but do not intersect it. However, photons that they generate by both bremsstrahlung (Berger and Seltzer, 1972) and the excitation of Ke line amission can penetrate to the rocket location. -ГТЛЦЕСТСИИ Figure 1 Kraushaar (1974) derived a rough estimate of the forementioned atmospheric background from an electron spectrum based on preliminary data acquired by the Low- Energy Electron Experiment aboard the ЛЕ-С spacecraft (Hoffman et al., 1973). Although Kraushaar found that these electrons could have generated only a few percent of the observed diffuse x-ray flux in the 0.10-0.28 keV energy range by bremsstrahlung emission, he also found that KQ line emission from collisionally excited oxygen and nitrogen could increase the x-ray flux below 1 keV substantially. In spite of the fact that this preliminary work suggested a potential problem for rocket-borne soft-x-ray experiments, to the authors' knowledge a more detailed evaluation of the atmospheric component from magnetospheric electrons has not been reported. The purpose of this report is to present the results of a calculation of the x-ray flux at rocket altitudes resulting from magnetospheric electrons. Specifically, the x-ray fluxes produced, by electrons in the residual atmosphere above the usual rocket apogee (~200 km) at the White Sands, New Mexico and Kauai, Hawaii launch sites are calculated from extensive satellite observations of energetic electrons (including energy spectra and pitch-angle distributions) measured at geomagnet'.cally quiet times at L values appropriate to the launch sites, and a previously derived method for calculating x-ray fluxes generated by electrons in the atmosphere (Luhmann, 1977, Luhmann and Blake, 1977). The latter, which was designed for calculations of auroral x-ray spectra, takes into account the 81 energy degradation of the electrons in the atmosphere as well as the radiative transport of the x-rays. The results of the present calculation are estimates of the background x-ray spectrum at rocket altitude for the forementioned sites. These background spectra are compared with the reported diffuse flux. BACKGROUND CALCULATION The radiation belt electrons that reach low altitudes at low and middle latitudes are part of the quasitrapped population that is lost in the atmosphere near the South Atlantic Anomaly every drift period (Roederer, 1970). Satellite observations indicate that during geomagnetically quiet times some of these energetic electrons are present to a minimum latitude of about 200 km over White Sands and Keuai at nigh*. This fact is derived from pitch-angle distributions. Figure 2 shows a typical pitch-angle distribution of downward directed electrons at ~ 750 km in the vicinity of White Sands (longitude =*253°E, L-1.8) (A. L. Vampola, private communication). In this figure, the local (750 km) pitch angles for which electrons would mirror at 500 km and at 200 km are indicated. These angles are derived from the principal of the conservation of the electron magnetic moment in the geomagnetic field. It is apparent from Figure 2 that a substantial fraction of these electrons was present in the atmosphere over White Sands to. minimum altitudes of ~200 km on this occasion. The pitch angle distributions observed near Kauai lead to a similar conclusion for that site. Since the electron observations are for typical quiet times, the configuration depicted in Figure 1 applies in general to experiments launched from White Sands and Kauai. The electron spectrum observed near the equator at geomagnetically quiet times on drift shells that pass over low-latitude ground sites at a few hundred km altitudes (White Sands is at L=»1.8 and Kauai is at L^1.2) can be approximated by exponentials with e- folding energies (E ) between 30 keV and 100 keV (i.e., see Lyons and Williams, 1975). In the present calculation it is assumed that an energy spectrum of exponential form with EQ = 30-100 keV describes electrons at all points on the L=l to 2 drift shells (the low- altitude spectral data support this assumption) and thus electrons with these energy spectra are present in the atmosphere above White Sands uid Kauai at the top of the practical atmosphere (-500 km). At that point, the electron flux is reduced, because of magnetic mirroring, by a factor of .015 at White Sands and .070 at Kauai compared to the equatorial flux. 42 We previously developed a method for calculating the atmospheric x-ray emission at arbitrary altitude generated by an electron influx with a specified energy spectrum (Luhmann, 1977, Luhmann and Blake, 1977). This method was employed to calculate the photon spectrum from the spectrum of magnetospheric electrons over White Sands and Kauai described above using the atmospheric density given by the (1000°K exosphere) COSPAR International Reference Atmosphere (CIRA, 1965). Because the pitch angle distributions show that the electrons do not penetrate below -200 Km and attentuation between 200 km and the rocket altitude is negligible, it is assumed that the x-ray flux at the rocket is given by the flux at 200 km. Figure 3 shows the resulting atmospheric x-ray spectra. The enhancement near 0.5 keV is from KQ emission by collisionally-excited oxygen in the atmosphere above the rocket assuming a detector bandwidth of 150 eV. The upper and lower limits of each cross-hatched band correspond to the Е = 100 keV and Е = 30 keV cases, respectively. The stipling shows the range of the reported diffuse galactic background spectrum (see Bunner, 1974; Cash et al., 1976). This comparison suggests that magnetospheric electrons that mirror above the rocket at the midlatitude launch sites provide a substantial fraction of the reported diffuse flux near kilovolt energies. 3, 10 10 ENEHGf IkiVI PITCH ANGLE Figure 2 Figure 3 83 CONCLUSIONS Method! developed for the study of auroral к-ray ipeetra have been applied Ip the calculation of atmospheric x-ray fluxes from quesitrapped mafnetoapheric electrona observed above two midlatitude sites. The sites of White Sands, New Mexico end Kauai, Hawaii were selected because many observations of the diffuse x-ray background have been carried out from rockets launched at these ranges by x-ray astronomers. The cal culated x-тау fluxes are comparable to the observed diffuse fluxes at energies below a few keV. The apparent absence of an interstellar absorption effect in the observed diffuse x- ray spectrum has previously caused consternation (Grewing and Walmsley, 1974, Hill and Silk, 1975). As a result, various theories have been formulated to explain the intense soft energy flux (Bunner, 1974). The foregoing analysis suggests that much of the observed flux below a few keV may have been locally generated by magnetospheric electrons in the residual atmosphere above the rocket-borne detectors. Accurate corrections for the atmospheric contribution discussed here are difficult to make because the magnetospheric electron population is highly variable both spatially and temporally (Seward, 1974; Paulikas, 1975). In fact, the variability in the measurements of the low energy diffuse spectrum as noted by Henry et ai. (1971) and Shukla and Wilson (1971) can be attributed to the variability of the low altitude magnetospheric electron population. The results reported here indicate that although rockets may provide a suitable pldtform for discrete source observations at x-ray energies £ 1 keV, diffuse component observations should probably be carried out from satellites which escape atmospheric effects entirely. ACKNOWLEDGEMENTS The authors wish to thank H. H. Hilton and L. M. Friesen for their assistance in the computational effort and A. L. Vampola, J. H. Underwood and H. R. Rugge for helpful discussions. This work was supported by The Aerospace Corporation Company-Financed Research Program. 84 REFERENCES Berger, M. J., and Seltzer, S. M., 197», J. Atmoa. Terr. Phys., 34, 85. Beuermann, K. P. 1971, J. Oeophys. Res., 76, 4291. Bunner, A. N. 1974, in Proceedings of the International Conference on X-rays in Space, Calgary, Alberta, Canada. Cash, W., Malina, R., and Stern, R., 1976, Ap. J. (Letters), 204, L7. CIRA 1965, COSPAR International Reference Atmosphere, North Holland, Amsterdam. Gorenstein, P., Kellogg, E. M. and Gursky, H., 1969, Ap. \, 156, 315. Grewing, C. M. and Walmsley, C. M., 1974, Astr. and Ap., 30, 281. Henry, R. C, Fritz, G., Meekings, J. F., Chubb, T., and Firedman, H., 1971, Ap. J. (Letters), 163, L73. Hill, J. K., and Silk, J., 1975, Ap. J., 198, 299. Hoffman, R. A., Burch, J. L., Janetzke, R. W., McChesney, J. F., Way, H. S., 1973, Radio Science 8, 393. Horstman, H., 1973, in Gamma Ray Astrophysics, NASA SP-339, National Aeronautics and Space Administration, Washington, D. C. Imhof, W. L., Nakano, G. H., and Reagan, J. B., 1976, J. Geophys. Res., 81, 2835. Kraushaar, W., 1974, in Proceedings of the Workshop on Electron Contamination in X-ray Astronomy Experiments, Goddard Space Flight Center, Greetibelt, Maryland. Luhmann, J. G., 1977, J. Atmos. Terr. Phys., in press. Luhmann, J. G. and Blake, 1977, J. Atmos. Terr. Phys., in press. Lyons, L. R. and Williams, D. J., 1975, J. Geophys. Res. 80_, 943. Paulikas, G. A., 1975, Rev. Geophys. Space Phys., 13, 709. Peterson, L. E., 1975, Ann. Rev. Astr. and Ap., 13, 423. Roederer, J. G., 1970, Dynamics of Geomagnetically Trapped Radiation, Springer-Verlag, New York. Seward, F. D., 1974, in Proceedings of the Workshop on Electron Contamination in X-ray Astronomy Experiments, ed. S. S. Holt, Goddard Space Flight Center, Greenbelt, Maryland. Shukla, P. G. and Wilson, B. G., 1971, Ap. J., 164, 265. '6С? 8S ££Ш OBSERVATIONAL CONSTRAINTS ON THE POSSIBLE EXISTENCE OF COSMO LOGICAL COSMIC RAYS T. Montmerle Service d'EIectronique Physique, Centre d'Etudea Nucllaires de Saclay France The possibility that cosmological cosmic rays ("CCR" : protons and or particles) may have existed in the post recombination era of the early universe (iM00) is examined. In this context, the CCR interact with the ambient gaseous medium. High energy collisions (£ 1 GeV/n) give rise to diffuse background y-rays via n* decay, and low energy collisions (M0-100 MeV/n) give rise to light nuclei : "Li, 'Li and 'Be (via the or + a reaction), D and He (via p + or reactions). Taking expansion and ionization losses into account, a system of coupled time-dependent transport equations is solved in the case of a CCR burst. The 1-100 MeV Y-ray background spectrum and the light element abundances are then taker, as observational constraints on the CCR hypothesis. It is found that, in this framework, it is possible to account simultaneously for the y-ray back ground spectrum and for the otherwise unexplained Li/H ratio, but there are some difficulties with the 7 Li/6 Li ratio. To avoid these, it is possible, because of the spread in the у-гаУ data, to lower the CCR flux, so that the ССЯ hypothesis cannot be ruled out on this basis at present. I. THE COSMO LOGICAL COSMIC-RAY HYPOTHESIS a) The astrophvsical context . It has long been known that, there exists a "casmological window" (Stecker 1973 and refs therein) in the Y_rfy r»nge (~1-100 MeV,hereafter "у-гаув") of the diffuse background. This 'window" allows us potentially to see (with Y-ray "eyes") up to redshifts as high as "~100-300, due to the low opacity of the universe in this energy range. Our у_гаУ "eyes" (i. e., our detectors), are perhaps not quite suited to look that far at present (see Montmerle 1975, 1977a .hereafter Ma), but we retain here the possibility that у-гаув may be produced in high-energy events, in the post-recombination era of the universe. This idea was in fact first put forward by Stecker (1969), who suggested that the then discovered "bump" in the 1-10 MeV range of the у-г*У background spectrum (Vette et al. 1969) originated in red- shifted neutral pion decay. Although the original "bump" has been substantially reduced, essentially because of instrumental background problems (e. g. ,Trombka et al. 1977). the theoretical interpretation of this feature in terms of тт* decay remains the most likely at present (e. g., Stecker 1977). 86 although alternative explanation*hive been proposed (Rocchla et_al. 1976). In our context neutral pions can be produced either by matter-antimatter annihilation (NFJ hypothesis, hereafter) or by inter actions between high-energy particle* and the ambient gas (cosmologic- al cosmic ray, or CCR, hypothesis, hereafter). В е cause of strong difficulties brought to light recently (e.g., Ramani and Puget 1976, Combes et al.1976, Steigman 1976), the NN hypothesis seems much less likely at present, although some questions remain controversial (see remarks in Stecker 1977 ; also Aldrovandi and d'Olival 19.76). This situation provides a strong motivation to study the alternative CCR hypothesis, which, because of the early success** of the NTT hypothesis (e.g. Stecker lud Puget 1972, Omnes 1972), has not been looked into in detail up to now. ' b) Interactions of CCR with the ambient gas. Although some kind of CCR sources have been proposed ("protarr". Stecker 1971 ), we shall concentrate here only on observational constraint» on the possible existence of CCR. One of such constraints is the already mentioned у_г*-У back ground spectrum. The v-rays are produced essentially by ~ 3 CeV protons (Stecker 1973) colliding with the ambient gaa, via rr* decay. 3ut at the low-energy end of the CCR spectrum, other interactions take place, leading to as many observational constraints. Since, at the epochs we consider, the air.bient gas is made up almost entirely of hydrogen and helium-4 (with 4He/H *• 0.1), the postulated (unknown) mechanism that brings proton* to cosmic-ray energies must also similarly accelerate *He nuclei. As a consequence, in the ~10-100 MeV/n energy range, light elements are formed : D and 3He nuclei via pp and per reactions, and 'Be nuclei via crcr reactions. These nuclei are slowed down and thermalized in the ambient gas, out of which galaxies subsequently form. The abundances (absolute and relative) observed now of these,light elements thus provide further constraints. Note that the light nuclei may be produced via excited states which decay through nuclear v-ray emission, thereby providing a contribution to the X-and у-ray diffuse background. Light element production, in relation to the у~г*У background, has been studied in forthcoming papers (Montmerle 1977b, c, d, here after Mb, Mc, Md ) ; the framework will be summarized in what follow* but details on deuterium production will be given in a companion paper (these proceedings, OG-130, hereafter Ml). Nuclear ч-тьу production has been examined (Montmerle 1977 e, hereafter Me) and will be detailed in another companion paper (these proceedings, OG-45, hereafter M2). Details on the CCR model and the transport equation* will be given in the next section ( \ II). A comparison with observation* of the results, applied to Y-ray and light element production, and conclud ing remarks,will be presented in the final section ( «till). 87 CI. THE CCR MODEL AND PARTICLE TRANSPORT The ambient gas density it governed чjsentially by two parameters : the Hubble constant HQ and the deceleration parameter _1 1 q0. We take here H0 •' 55 km s Mpc" and q0 » 0.1 (for a discussion see Mb). The simplest CCR model involves a burst of protons and ar- particles (with o/p = 0.1) at some redshift zg. The injection spectrum r at ze is proportional to (E + EB)~ , where Е is the kinetic energy per nucleon cf the CCR. Eg will be taken either as the proton rest i mass E0 ("total energy" CCR spectrum) or as zero ("kinetic energy" CCR spectrum). Also, we take Г = 2.6 (see discussion in Ma). For i. given value of q0, the results depend only on the parameter zs. Now the behaviour of aU particles (CCR protons and a- particles, CCR-produced y-raya and light nuclei) is governed by transport equations that can be cast into the form : 3N. lit1 - ?1<ьА.н> + "кЛк"^ • Ч,н«'«> (2"1} In eq. (2-1), Nk Hs Nk| H(E, z) is (in |peV/nJ " ) the number density of moving particles k, per unit energy interval, of kinetic energy per nucleon Е at redshift z, with respect to the ambient gaseous hydrogen in a oomoving volume, i.e. Nj^/njj 0(l+z)*3. Quantities of the form 3£^ stand for X^ (dt/dz), thereby relating the standard time dependence of eq. (2-1) to its redshift dependence. The transport function bk= (ЪЕ/dt)^. represents the energy losses per nucleon per second (in CeV/n 8-1) of particle k (ionization and/or expansion') ; the destruction of particles k is accounted for by the destruction lifetime Tj> ^ (see Mb for details). Qy. н(Е, Z) *• numerically defined like N^ н(Е» z) in a unit redshift interval ; it is the source-function for the production of particles k. Eq. (2-1) is solved semi-analytically in Mb (see also Montmerle 1977f, hereafter Mf ) in the present context, for any kind of particles, relativistic or not. This allows ,\o compute at the same time the evolution of the CCR flux,' the у~гаУ background, the abundance of light nuclei after thermalization in tfie ambient gas (Mb) or the nuclear v-ray flux (Me) To fiud the source-function appearing in eq. (2-1), let us consider the collision between a particle i of energy E1 (moving with a velocity eg") and a particle J at rest (of number density n») giving birth to a particle of energy E. (Note : a given particle will be denoted by capitals if at rest, by the corresponding minuscule or symbol if moving.) Let ст... . (Е1, Е) be the cross-section for this process. The source-function for the production of particles k is then Q (E,z) , , , k,H •JaijH.kCE ,E)nJce NijH(E ,z)dE» (2-2) 88 or, as the case may be, a turn of equations of the form (2-2). This equation shows how the transport equations for particles i and k, of the form (2-2), are coupled. Now, we shall make the simplifying assumption that, notwithstanding the kinematical details, Е is proportion al to E'. Depending on the reaction (especially in the case of D- producing reactions), this may be more or less true and will be discussed in Ml. In. other words, we write the cross-sections as : a (E E S(E ,4 2 3 ij-* '- > " °А(Ю(Е,) " h£ > • < - > describing the kinematics of the reaction (for instance, *Q4H °*0.25). k е y-r*.y source-function is taken from Stecker (1971b ; see also Mb). A more general calculation is given by Bonnardeau (1977). Ш. APPLICATION TO Y-RAY AND LIGHT ELEMENT PRODUCTION a) The v-ray background. In order to find the absolute value of the CCR flux, we normalize the results of the calculations of the v-ray background to an intensity Iv at a given energy Ev. The choice (I,,, Ey) has been discussed in Mi taking into account the spread in the Y-ray data, especially in the ~10 MeV region. Hereafter, as in Md, we take 1 1 ly = 6.10" cm" e" sr" MeV" at Ey = 4 MeV. Indeed, it is found (see also Mf) that the corresponding theoretical spectra (which include a contribution I =10"^(EY/l MeV)-^.3 cm"' a-1 sr-1 MeV_1 [see Fichtel et_al. 1975 , Mb] extrapolated from the hard X-ray range)fit the most recent data, obtained by the Apollo (Trombka et al. 1977) and SAS-2 {Fichtel et al. 1975) experiments, for 60S se£150 (with q0 =0.1). Other normalisations have been used to estimate the uncertain ties on the results (light.element abundances etc.) due to the uncertain ties on the у~гаУ background (Mf ; see fig. 1). Note that the theoretical spectra are almost independent of q0 (Mb, c). It is then clear that it will always be possible to find a value oi z, leading to a good fit between theory ana observation. However, the use of other normalizations shows that, provided the bump still exists, the values of ze always lie in roughly the same range as quoted above. In passing, note that the ratio between the CCR flux intensity now (in intergalactic space) and the galactic cosmic-ray (GCR) flux intensity is in the i-ange 10_4-10-6 (Mb), i CCR do form by far the bulk of the cosmic rays observed in the vicinity of the Earth, and as a consequence, CCR have nothing to do with the so-called "universal" cosmic rays put forward by some authors(e. g., Brecher and Burbidgel972jL b) The lithium abundance. The cross-sections o/ and a-j (appearing in eq. 2-3) for the production of °Li and 7 Li by е е reactions (taking the decay of 7Be into 'Li into account) have been discussed in detail in Mb ; (on the 4He(cr, )°Li reaction, see also Me) ; their values above a few tens of MeV/n are highly uncertain. While this uncertainty 89 has been shown (Mc) to affect the final results on the ^Li and Li abundances by a factor of 2 at moat, it become* quite important when comparing the calculated 7Li/"Li ratio with observation» (see below). The observed . abundance 7Li/H is ~10"9~ (e.g., Reeves 1974, Boesgaard 1976), and remains essentially unexplained up to now (Reeves 1974 ; see also a discussion in Mc). Fig. 2 shows the \ calculated 7Li abundances for a total energy CCR injection spectrum (EB = E„) normalized as explained in the preceding section, as a function of z , for qo=0.1. It can be seen* that the CCR hypothesis is then able to account for the Y-ray background spectrum and the 'Li abundance simul taneously. The uncertainties on this result, due to the spread in the у_гаУ data, are also shown. On the other hand, with the same normali- ._!_ A compilation of observational data zation and for a kinetic energy on the y-ray backgrour 1 (see Mb for refe- CCR injection spectrum (Eg=0), rences), along with theoretical curves 7 Li is overproduced by a factor coa'>uted in Ле CCR hypothesis, normalised ~103. Thus, in practice, one *" shown' Tbe curves depend on z , but es- has to discard any CCR ««rally not on 4Q. injection spectrum much steeper (at low energies) than a total energy spectrum. Consider now the Li/ Li ratio. This ratio has been, measured only j.n meteorites and only lower limits exist for stars (~10, Rettvts 1974 and refs therein). It is shown in Mc, d that, in the CCR hypothesis, the computed ratio falls short by a factor ~3 - 5 to account for the observed value with reasonable values of the o6 and e? cross-sections. The dependence of the 7Li/6Li ratio with o Jo, is shown on TableЛ\Ш\ Therefore, it is necessary to find a mechanism able to destroy the CCR-produced bLi (the observed 6Li would then be produced only by GCR, see Meneguzzi et_al. 197П wb4V leaving 7U intact. Such a mecha nism exists and is sketched in Mc, i, out n appears that most-likely it is impossible to moetify the too large CCR-proaiicea 7Li/6Li ratio.' 90 Ла a consequence, to Avoid conflict with 7 Li/6 Li I i observations, it ia neceaaary to lower the CCR flux ; at pre a eat thia ia permitted by the apread of the у-**У da** •ю- •"- У A \1 but then, the CCR hypotheaia A doea not explain the ?LA A A abundance any more. Thua, 4 the CCR hypotheaia can be 4 ^Ap 10- — у at leaat compatible with the У y-ray obaervationa and the lithium iaotopic abundances. У Aa eh own in companion papers (Ml, M2) thia conclusion ia not modified by the consider *r ation of the abundance of deuterium or of the nuclear га а y-ray flux. Further у- У Чо 0.1 \ 2 observations in the 1-100 10"' MeV range, especially around ~10 MeV, are clearly ne.eded to accept or reject the CCR hypothe sis. 1000 100 ю 1+2, TABLE I The calculated 7 6 Li abundance as a W Li/ Li function of z . The curves are label*Л ac V. cording to the normalizations indicate! in Fdg.1. For each curve, the upper limit on 1.37 z is determined by the "CCR lifetime cons Ю 5.42 traint" (Mb,c); the lower limit by ihi fit 100 7.69 to the observed if-ray background spectrum (+) Above 100 MeV/n, with (Mc). Observations lie between the thin horizontal lines. Mncaerl* T. 1977*. pr*prist (И*) 1977*. Ixoe. 12th MLAB lyep. (КП АМготашЦ t, 4'Olinl J.I.I. 1*76. Ap. 4 tp. Set. 44, 471 *- 1972, Fkya. taper» 3C, I lontutd A.M. I»», rub. A*tr. loe. !кШс, $1, 393 . A., rueec •»•!•• I97«. A*tr. Ap. 12. «II lecaarduu M. 1977, preprint > I. 1974, Am. lev. Aatr. Ap. 12, 437 •recber K.. lurki*** G. 1972. Ap. J. 174. 233 locchia k., et «1. 1974, Ap. J. 209.T50 Ccrtti Г., «t «1. 1973. Ap. 4 Sp. teiT~3>, 131 Steeker t.V. \Ш, ««tor* 224, «75" Tickcel C.E.. «t «I. 1*73. If. J. ,198, 173 1971, «ASA 11=279 «вцт! И., «t «1. 1*71. Aacr. Ap. IS. 337 1973, SMA 1Г-339, p. 211 HnEMrt* It 1975, Ap. J. 197. 233 — 1977, Ap. J. 212, 60 1977». »roe. T5tB ESUI Sjrap. THE CONTRIBUTION OF DISCRETE SOURCES TO THE FLUX LF EXTRAGALACTIC GAMMA RAYS A.W. Strong, A.W. Wolfendale (.nd D.M. Worrall, Phytic» Department, University of Durham, England. Estinatei have been made of the likely contribution to the diffuse gamma ray flux from a variety of objects. If the flux is indeed due to extragalactic sources then normal galaxies are probably responsible for only a few per cent of the flux but radio galaxies can, in principle, produce much of the radiation in the range 1-10 MeV. Seyfert galaxies and clusters may be responsible for most of the 100 MeV emission.' 1. Introduction. In addition to the flux of y-rays of clearly Galactic origin there appears to be a near-isotopic diffuse component which has been identified by many authors as being of extragalactic origin. Before endeavouring to explain its origin it is necessary to sound г cautionary note. All detectors have their attendant backgrounds - y-rays produced locally and not of cosmic origin - and the identification of the magnitude of the flux, which must be subtracted from the signal, is not such a trivial problem. Perhaps even more important is the difficulty of identifying the fraction coming from the extragalactic sources in the presence of an undoubted component from the Galactic halo. The magnitude of the latter is inevitably the subject of debate in view of the uncertainty in the existence of the halo itself. The contribution .from the halo has been considered recently by some of the present authors (Worrall and Strong, 1977). It is concluded that inverse Compton scattering of electrons in the Galactic halo may well produce a large part of the diffuse flux. Although a large anisotropy of arrival directions should result from the halo component (we are not at the centre of a spherical halo) the measurements reported to date do not-- appear to be sufficiently precise to rule out this hypothesis. In what follows we consider genuine extragalactic components, bearing in mind the remark about the halo. 2. Contribution from Normal Galaxies. In an earlier work (Strong et al., 1976), taking the value for the emissiyity of the Galaxy found by Strong and Worrall (1976), i.e. 1.3 x 1042 Y'« e"1 above 100 MeV and, using a variety of arguments, the contribution to the EG flux was shown to be •ъЬХ of the measured diffuse flux (Т^(> 100 MeV) - 1.9 x 10-5 cm"2 s"l sr_1: Fichtcl et al., 1975). Worrall anu Strong, using a halo model with a particular value for the diffusion mean path of the electronp (X • 20 pc) find a Galactic emissivity of 7 x 1042 y's s-1 above 100 MeV and translating this to all normal galaxies the 4Z is raised to 20Z. 92 Insofar as this value of X giver the bulk of the measured diffusa flux having сом fro» the halo, 20Z is presumably a lower llait to the fraction of the diffuse flux having come from genuine axtragalactic aourcaa (although a smaller fraction can hava coma froa nonaal galaxies). 3. Contribution fro» Radiogalaxies. If thera ia a correspondence between the radio emission of a galaxy and its y-ray emission than radiogalaxies will contribute «ore than normal galaxies. This arises because of the argument of Schmidt (1972) that the contribution from normal galaxies ia only about 251 of the total flux (at 408 №1). Reverting to the AX of §2, the radiogalaxies would then contribute about 16% of the total. There does appear to have been detection of Y-rays from a radiogalaxy: Cen A. Hall et al., (1975) have detected a continuum extending to~10 MeV and likely lines at 1.6 and 4.5 MeV. The 178 Mix flux from Cen A is 4800 Jy and the background intensity Mt the same frequency ia 1.9 x 10~22 U m~* sr~l Hz~l yielding the contribution from similar radiogalaxies shown in Figure 1. 4. Contribution from Seyfert Galaxies. No Y~rays have yet been detected from Seyfert Galaxies so that predictions are very model dependent. Using the model of Bergeron and Salpeter (1971, 1973) the Y-ray flux in the region of 100 MeV comes out to be close to observation but the closeness may be fortuitous. 5. Contribution from Clusters of Galaxies. Rowan-Robinson and Fabian (1975) have examined the possibility of the X-ray background being due in part to clusters. With an evolutionary situation the contribution at 4.1 keV is in the range 22-65% and, if inverse Compton interaction is dominant, the flux will be very significant at Y~ray energies. However, thermal bremsfrahlung may well be the main X-ray production mechanism, in which case the \-eay contribution will be negligible. A large fraction of the energetic y-rays (> 100 MeV) may arise from interactions of trapped cosmic rays with the gas in the cluster. Stapley et al. (1977) have devised a model which can explain the cosmic rays above 1018 eV or so in terms of particles escaping from clusters and this gives (perhaps by chance) the correct order of magnitude of cluster y-ray emission. Insofar as the mechanism is by way of IT0 production it will only be relevant above about 100 MeV. 6. Conclusions. There appears to be no shortage of contenders for . explaining the diffuse background of y-rays. However, before further speculation is made, it is necessary to have more accurate experimental data on the intensity as a function of Galactic longitude and latitude. The observation of a significant anisotropy will indicate a big halo contribution and small non-normal galaxy contributions. The absence of anisotropy will signify larger contributions from non-normal galaxies and, perhaps, Cosmological effects. 93 Norm«l (halo) 4fi* Ь 7 * to9 Figure 1. Energy spectrum of diffuse ^-rays. The beet estimate and limits to the measured spectrum from our earlier work (Strong et al. 1976) are shown. The contributions are as follows: Normal: normal galaxies for.no halo Normal (Halo): normal glaxies with halos as in model of Worrall and Strong (1977). This line is essentially the lower limit to extragalactic contribution in general; if such halos do exist then the bulk of the diffuse flux will have come from our own halo, if not then other E.G. sources are necessary such as Radio galaxies, Seyferts, etc. Radio: scaled from the observations by Hall et al. (1975) of Cen A. 94 References. Bergeron, J. and Salpeter, E.E., 1971. «p. J. t»tt., 9, 121; 1973, Astron. and Astrcphys., 22, 385. Fichttl, C.E. et •!., 1975, Ap. J., 19S, 163. Hall, R.D. tt al., 197S, Proc. 14th Inc. C.R. Conf., Munich, 1, M. Rowan-Robinson, И. and Fabian, A.C., 1975', Hon. Not. K. Aicroo. Soc, 170, 199. Schmidt, M., 1972, Ap. J., 176, 238. Stapley, N.R. Wdovccyk, J. and Wolfendale, A.W., 1977, these proceeding! (OG 192). Strong,. A.W., and Worrell, D.M., 1976, J. Phye. A., 9, 823. Strong, A.W., Wolfendale, A.W. and Worrall, D.M., 1976, J. Phye. A., 9, 1553. Worral, D.M. and Strong, A.W., 1977, Aatron. and Astrophye., (in the press). 95 THE GAMMA-RAY LUMINOSITY OF SPIRAL GALAXIES. ITS EVOLUTION AND ITS CONTRIBUTION TO THE DIFFUSE BACKGROUND ABOVE 100 MeV G.G. Lichti3. G.F. Bignami1, J.A.Paul2 1. Laboratorio di Fisica Coamica е Tecnologie Relative del CNR, MiUno, Italy 2. Service d'Electronique Phyaique, Centre d'Etudes NucKaires do Saday, France 3. Space Science Department of the European Space Agency, ESTEC, Noordwijk, The Netherland» The contribution of normal sp*"l galaxiea to the high galactic latitude gamma-ray background > 100 MeV ia examined in the light of the eatimatea of ita flux from the SAS-2 measurement». The gamma-ray luminosity of each object ia inferred from the known Milky Way value normalised to the corresponding optical quantity. Several possibilities are considered for the responsible physical production mechanism both diffuse and localized ; there are then aet to evolve wit., the galactic age according to 3 well-knowr evolutionary models. A final space- time integration leads to results, expressed as fraction ot the measured background. It is seen that the model presented can play an important role in the region > 100 MeV where the information on the spectral shape of the radiation is still ' very poor. 1. INTRODUCTION. A great wealth of theories has been put forward to account for the reported high galactic latitude gamma-ray flux reported by Fichtel et al. (1975) in terniB ot mechanisms ranging from p-p interactions in the early epoch of the universe to electromagnetic inter galactic cascades. Recently, Schlickeiser and Thielheim (1976) have shown that the contribution from the local gamma emission of the galactic disc to the medium to high latitude measurements cannot be neglected. In the present work, we set forth to show which fraction of the reported gamma-ray flux can be accounted for by summed emission from normal spiral galaxies, when one integrates back into the cosmo- logical past up to the onset time of the galaxies. Each contribution was taken to be proportional to the optical luminosity of the object, using as a reference point* the known emission of our own Milky Way, and was then set to evolve during the galaxy lifetime , That is one step forward with respect to what has been done firstly on the same line by Strong et al. (1976) who did not consider evolutionary effects. In the evolution factor thus arising are contained the physical hypothesis on the origin of the gamma-ray luminosity. In what follows, we consider several possible causes of the high energy gamma-ray luminosity. A discussion on the various possibilities will be given in § 2, leading to a number of starting cases. These will then be treated in parallel, namely, carried through the evolution process(§ 3) and the 96 •pace-time integration (§ 4). Result! will be presented and diicuned in § S, in term of fraction of the high latitude flux explicable with emission from normal galaxiea. 2. THE GAMMA-RAY LUMINOSITY OF A SPIRAL GALAXY. The overall gamma-ray emission from the disc of our own galaxy can bs evaluated above 100 MeV in about 1 to 2 1 (гг photona a'1. Thia value waa firatly derived from the OSO-3 measurement» (Krauaha^r et al. 1972) and subsequently confirmed by the SAS-2 satellite. Also, theore tical evaluations (e.g. Bignami et al. 1975) arrived to values "ell within such range, using as a reaponsible physical process for the emission the interaction of GeV cosmic rays with the total interstellar gas. However, the identification of two pulsars (NP 0S31 + 21 and PSR 0833-45) as sources of gamma radiation (Hartman et al. 1976, Bennett et al. 1977a) and the suggestion that several other may also emit pulsed radiation (Hartman et al. 1976) point to a new component in the luminosity. This is to be expected, to some extent, since it can be shown (Ogelman et al. 1976a) that pulsars radiate the bulk of their energy in the 10" - 10° eV range. Using different assumptions on the pulsar birth-rate and intrinsic luminosity Ogelman et al. (1976b) and Higdon and Lingenfelter (1976) obtain respectively 0.5 1041 and 4 1041 photons (>100 MeV) for the total galactic luminosity due to pulsars. It appears that the contri bution from known radio pulsars to the measured gamma-ray luminosity of the Milky Way can be a fraction of the total even if it falls clearly short of accounting for it all. In addition, some recent COS-B results (e.g. Bennett et al. 1977b) show that at least a part of the emission from the disc may be due to localized objects. Especially with this in mind, we shall explore' the rather extreme possibility (possibility A in this. context) in which the total of the gamma-ray emissivity is due to pulsar-like objects. This will represent the antagonistic situation to that where all the emissivity is assigned to diffuse processes (possibi lities В and C). 3. EVOLUTION OF THE GAMMA-RAY LUMINOSITY. If one adopt the possibility A for the gamma-ray luminosity of a galaxy, i. e. it all be due to pulsars, the evolution of this luminosity can be simply described in term of variation of the pulsar birth -rate with the galactic age. The birth rate is the only parameter considered since the observable pulsar lifetime (about 4 10' years, Gunn and Ostriker, 1970) or any fraction of it which is spent emitting gamma rays is small compared to the galactic age. On the other hand, following several theoretical works we can believe that the diffuse processes are largely dominant, and particular ly the interaction of cosmic-ray electrons and nucleons with the inter stellar matter. The evolution factor of the gamma-ray luminosity can be then described if one knows the evolution of its basic ingredients i. e. the interstellar gas and the cosmic rays. One can simply assume that the cosmic-ray (CR) density is proportional to the supernovae (SN) 97 rate arid that thl* relation can be maintained at any galactic age (possi bility B). Without any assumption on the nature of the CR sources, one can alao aaaume that the CR density la proportional to the gas density as e.g. in Bignami et al. (1975) and Paul at al. (1976) (possibility C). The evolution of the basic ingredients of the gamma-ray lumi nosity, which are then the pulaar rate, the SN rate and the gas content have been deacribed by several author* aa results of evolution model* of the galaxy needed to explain the observed elemental abundance* and the chemical composition of th*> solar system and its neighbourhood. We have derived the dependence of these basic ingredient* on the ' galactic age from 3 different evolutionary model* '. (1) the evolution model of Truran and Cameron (TC, 1971), (ii) the standard model of Quirk and Tinsley (QT, 1973), (Ш) the combined halordisc model of Ostriker aud Thuan (OT, 1975). 4. SPACE AND TIME INTEGRATION. It can be shown (lichti et al. 1977) that an isotropic and homogeneous universe filled with galaxie* which emit gamma ray* isotropically lead* to a diffuse flux at the present time : г в /max QQI+,)B .^ds ,, {i) J o (Eo ) =4TTH -= ' O J (1+л) Vi+*q0» o where H is the Hubble constant, TI0 the present number density of galaxies, C the speed of light, s the redshift, t the onset time of galaxies, q0 the deceleration parameter, Q the mean gamma-ray emis- •ivity of a galaxy. According to recent estimate*, we assume for the emissivity spectrum above 100 MeV a differential power law with a spectral index of 2. Then A f(«) Q Г(1+*)Е0. « 1 = 1 -Г (2) E„ (l+«)2 where f(z) is the evolution factor of the gamma-ray emiaaivity. Equation (1) then read* : A i. C f "** f(«) d« J (E0) = °—Г I „ (3) Z J ° 4rrH0E J (1+«) VJ>2qo« 0 If the gamma-ray emisaivity of a galaxy is related to the optical one ' (through (he mas*) we define the constant A as : ^G 44 _1 A x — x 10 MeV a (4) hfw where L^i* the mean optical luminosity of a galaxy, L.™. that of the Milky Way, the gamma-ray luminosity (> 100 MeV) of the Milky Way being 10*? photon* *~*. If Е is-the mean luminosity of the Univer se, i|_ • C/LQ and then *" equation (3) the fir»t factor i* only function 98 of L . The evolution factor f(s) has been computed for the 3 evolution model» and for the 3 poaaibilitiea of the gamma-ray luminosity produc tion and evolution. The retulting values of the second factor of the equation (3) are given in table 1. Evolution Evolution of the gamma-ray luminosity following posit model bilitv ABC OT 1. 07 2. 5 2.04 TC 1.68 11.4 13.5 QT 2.70 20.8 1.02 Table 1. Values of the integral in the relation (3) These values are obtained for q0 averaged over the rauge of values 0. 04 to 0.24 ; in this range the resulting value of the integral is weakly dependent on the parameter qQ. Finally, we have taken for the constants used in (3) : 1 3 10 H = 60 km s" Mpc" , LMW = 1.22 10 L© (Allen, 1973). An esti mate of the mean luminosity of the universe -can be found in Lichti et al. 1977. The most probable values range between 1 and 3 10^ LQ Mpc . 5. RESULTS AND DISCUSSION. Having estimates for all the parameters in equation (3) we present in the table 2 the values .of the diffuse flux (> 100 MeV) induced by the summed contribution of normal spiral gala xies for the 3 evolution models and the 3 possibilities of the evolution of the gamma-ray luminosity. Evolution Evolution of the gamma-ray luminosity following possi- model bility А В С OT 0.3/1. 0 0. 8/2. 3 0.6/1.9 TC 0.5/1.6 3.5/10.6 4.2/12.5 QT 0. 8/2. 5 6.5/19.6 0.3/1.0 Table 2. Diffuse flux (>100 MeV) in unit of 10_6ст"2в"1вг"1. The two values given in each case are for the two extreme estimates of the mean universe Iiminosity (see § 4). The results presented here, and in particular when compared to the SAS-2 measurement of the high galactic latitude flux above 100 MeV : 19.3 ±2.6 10-6cm"2s"isr-i (Fichtel et al. 1976) lead to some points of discussion. It appears that a significant contribution from normal spiral galaxies to the high galactic latitude flux cannot be excluded, even with respect to the SAS-2 rough measurement. In addition, if one believes that the galactic disc (Schlickeiser and Thielheim, 1976) and possibly the halo (Strong and Worral, 1976) contribute also to the high galactic latitude flux, the present results indicate that the summed contribution from normal galaxies may be a more substantial fraction of the extragalactic gamma-ray flux. In this latter case, the problem arises of the comparison of the spectral shape of the extragalactic gamma-ray flux with that of the summed contributiqn from normal 99 galaxiea. We only remark here that the real *hape of the extregalactic •pact rum above 100 MaV ia vary poorly known. Firstly, the only avai lable differential muiurimmti (Fichtal at al. 1975) only reach 170 MaV, and aacondly it now begin* to appear that the local diac contribution may totally maak the extragalactic flux abova 1 00 MaV. Therefore aenaible comparison* muat wait for better eatimatea of the galactic contribution and also for better measurement* of the energy apectrum above 100 MeV of our reference : the Milky Way. Acknowledgement*. One of ua (GGL) acknowledge* the receipt of an ESA feUowahlp. REFERENCES Allen, C.W. 1973, Л«trophy». Quantitiea, Athlone Preaa. Bennett, K,, Bignami, G. F., Boella, G., Buccheri, R., Hermann, W., Kanbacb, G., Lichti, CO., Maanou, J. L., Mayer -Ha**elwander, H.A., Paul, J. A., Scarai, L., Swanenburg, B.N., Taylor, B. G. , and Will*, R. D. 1977a, to be published in Aatron. and Aatrophya. Bennett, K., Bignami, G. F., Bonnardeau, M., Buccheri, R., Hermieo, W., Kanbach, G., Lichti, G. G., Mayer-Hasaelwander, H.A., Paul, J. A., Scarai, L., Stiglits, R., Swanenburg, B. N., and Will*, R.D. 1977b, Aatron. and Aatrophya. 56, 469. Bignami, G. F., Fichtel, C. E., Kniffen, P. A., and Thompioa, D.J. 1975, Ap.J., 199. 54. Fichtel, C. E., Hartman, R. C., Kniffei>, D. A., Thompaon, D. J., Bignami, G. F., Ogelman, H. B., Ozel, M. E., and Tomer, T. 1975, Ap.J., Jjif, 163. Gunn, J. E., and Oatriker, J. P. 1970, Ap.J., 160, 979. Hartman, R. C., Fichtel, C...E., Kniffen, P, A., Lamb, R. C.. Thompson, D. J., Bignami, G. F., Ogelman, H., Orel, M., and Tamer, T. 1976, Proc. of. the Int. Symp. on Gamma Raya, GSFC preprint X-662-76-154-12. Higdon, J.C., and Lingenfelter, R.D. 1976, Ap. J. (Letter), 208. LI 07. Krauahaar, W. L., Clark, G. W., Garmire, G. P., Borken, R., Higbie, P., Leong, C. , and Thorioi, T. 1972, Ap. J.. 177. 341. Lichti, G.G., Bignami, G.F., and Paul, J. A. 1977, preprint. Ogelman, H., Ayaeli, S., and Hacinlijan, A. 1977a, Proc. of the Int. Syir.p. on Gamma R»y«; OSFC preprint X-662-76-154, U8. Ogelman, H. B., Fichtel, C. E., Kniffen, D.(A., and Thompaon, D. J. 1976b, Ap.J., 209. 584. Oatriker, J. P., and Thuan, T. X. 1975, Ap.J.. 202, 353. Paul, J., Caase*, M., and Ceaaraky, C.J. 1976, Ap.J., 207. 62. Quirk, W., and Tinaley, B. M. 1973, Ap.J.. 121. 559. SchHckeiaer, R., and Thielheim, K.O. 1976, Nature, 261, 478. Strong, A.W., Wolfendale, A.W., and Worral, D. M. 1976, M.N.R.A.S. 175. 23. Strong, A.W., and Worral, D. M. 1976, preprint. Truran, J.W., and Cameron, A. G. W. 1971, Aatrophya. and Space Sci., 14, 179. 100 COSKIC DIFFUSE САЖА-RAYS AT MKDXUM INCtRIES R. Stephen White, Shin H. Moon, James H. Ryan, Robert 1. Wilson, Allen D. 2yc*> Physice Department and Institute of Geophysics and Planetary Physics Unlvaralty of California, Rlvaralda, California 92521. U.S.A. and Bruce Dayton Phyalcs Department California State University, Los Angeles, California 90032, U.S.A. Abstract. The fluxea of cosmic diffuse gamma-rays are reported In 6 energy Intervals from 2 to 25 MeV. They were obtained with the UCR gaaaa-гау tele scope flown on a balloon launched from Palestine, Texas on Hay 13, 1975. The observed fluxes are compatible with the energy distribution of 2.65 x 10"2 E-2'08 (photona/cm2-s-ster-MeV) proposed by Oennla tt al. (1973) at lower energies, steepened at higher energies to meet the slope of Е"2'1* of Flchtel et al. (1975) for energies above 35 MeV. No statistically significant devia tion from ieotropy In direction of the cosmic diffuse gamma-rays were ob served. These observations give upper limits of a few x 10"3 photons/cm2-s for a number of possible sources. 1. INTRODUCTION. The shape of the cosmic diffuse gamma-ray energy dlatrlbu- tion in the energy range of 1 to 30 MeV has important theoretical cosaologi- cal implications. Morrison (1958) early listed mechanisms that could be Im portant for x-ray and gama-ray production. Since then many theorists have , discussed these and additional source mechanisms. It la important, therefore, [ that the shape of the observed cosmic diffuse gamma-ray energy distribution be ' correct and not distorted by measurement. This has been emphasized by a num ber of authors including Golenetsfcli et al. (1971), Pal (1973), Daniel and Lavakare (1975) and Horatman et al. (1975). The cosmic diffuse radiation Is considered isotropic in space and constant in time, so it Is not possible to use variations in direction or in time to isolate the gamma-rays, as is pos sible from a directional source. The observer must subtract every background, measured and calculated, and the remaining flux, unexplained In any other way, is the cosmic diffuse gamma-ray flux. Several unwanted backgrounds have already been identified. Induced ac tivity, that gives delayed counts in Inorganic crystals such as Hal, Is a serious problem In the energy range from 1 to 30 MeV (Dyer and Morfill, 1971; Dyer et al., 1972; Flshman, 1972; Trombka et al., 1977). The induced activi ty is especially annoying with detectors on spacecraft that are bombarded by cosmic ray particles and protons trapped in the .radiation belt. Original measurements of Vette et al. (1970) were found to be much too high because of the Induced activity. The measurements of Trombka et al. (1973), obtained on Apollo 15, although corrected for induced activity, were »i.ill too high. Ths latest Apollo 16 and 17 results at* now limited to 1-10 MeV and have been lowered significantly (Trombka et al., 1977). In a double Compton scatter telescope the Induced activity, with its delayed counts, is no problem. When flown In the earth's atmosphere, however, it is possible for the albedo neutrons to interact with the carbon in the organic scintillator to produce 4.4 MeV gamma-rays that may be a source of background. The albedo neutrons can also be captured by the hydrogen, after being moderated, to produce 2.2 MeV gamma-rays. The earth's albedo neutron flux ia almost constant near the top of the atmosphere (Prasder et al., 1974) so these gamma-rays are not eliminated with the atmospheric growth curve. The carbon and hydrogen reactions have been studied by Whit* and 101 SchOnfelder C1975) and verified by SchOnf .ilder et al. Q975). The correc tion» for the carbon Interaction* In the Max-Flanck-tnstltute, IT!, telescope were found to be about 15X of the flux remaining attar subtracting the atjeee- pherlc flux belov 3 MaV and about SOS from 3-10 HaV. The hydrogen capture reaction was netlitlble. The correctlone for carbon Interactions la Пе Oml- ver~*ty of California, Riverside, UCR, double Coapton scatter telescope gives In nUt 1 are considerably lower, laaa than about 201 of the coamlc diffuse fluxaa for gaaaa-ray energies laaa than 10 NaT. Thla la largely because the* hydrogen ro carbon ratio in the OCR alnaral oil scintillator la 1.82 Instead "f the l.l In the plastic scintillator. In tha liquid scintillator at eater- glea below 100 MeV, the probability of np elastic scstterlng Is larger than the probability of C12Cn,XY)Z* Inelastic scattering. A neutron preferentially scatters from hydrogen, losing half its energy In each acatter on tha average, rapidly degradea in energy below the threshold tar producing a 4.4 MaV gaama-ray from carbon. Also the 1 a x 1 a liquid scintillator of the OCR teleacope la ao large that one part acta aa a neutron shield for other parts» thus, reducing the gaaaa-ray production further. The capture of neutrons by hydrogen to give 2.2 MeV gaaaa-raye was found to be negligible. TABLE 1. 2. OBSERVATIONS. The PERCENTAGE RATIO OF NEUTRON-CARBON INDUCED BACKGROUND comic diffuse gaama- TO THE EXTENDED SPECTRUM OF DENNIS ET AL. C1973), Г** Ш*?"?Г"?"ПД11.,№Г* 2.65 x 10-2 E-2.08 Photons/c»2-e-ster-MeV ?*?,** ** "^ „ double Coapton acatter Energy Zenith Angle (deg.) gaaaa-ray teleacope (KeV) (Herzo et al., 1975; 5* 15* 25* • 35* 45* Zych et al. 1975) 2-3 0.7 1.0 flown on a 0.54 all- 3-5 5.7 2.8 1.4 1.4 lion a3 balloon ac 5-7.5 17. 10. 8.3 6.9 5.9 3.5 g/ca2 residual 7.5-10 20. 23. 11. 7.6 atmosphere launched 10-15 21. 47. 23. 15. froa Palestine, Texaa 15-25 4.0 120. 90. on May 13, 1975. For a description of tha telescope, see paper 0G-16 of this conference. The cosmic diffuse gamma-ray fluxes during ascent and descent were determined by a fit of the relation T - a + bx + c e"*'* (1) to the total downward gamma-ray flux T (photons/cm2-s-ater-MeV). The downward atmospheric flux bx is considered linear from 90 g/cm2 residual depth to the top of the atmosphere. The neutron induced gamma-ray flux, a, is calculated by the method of White and Sch»nfelder (1975). The mean free path X (gm/cm2), as a function of thezenith angle and gamaa-ray energy, is considered known— The quantity x (g/cm2) Is the residual atmospheric depth. The atmospheric gamma-ray fluxes b (photons/ca2-e-ster-MeV~(g/ca2)) end the cosmic diffuse flux c (photons/ca2-s-ster-MeV) are determined from a best fit to the data. The atmospheric fluxes are published elsewhere (Ryan et al., 1977). Since b is Independent of time, equation (1) can also be used to find c at other times and, thus, other directions on the celestial sphere. 3. RESOLTS. A. Energy Distribution. The cosmic diffuse fluxes are obtained from the values of c that result from the best fit of equation (1) to the data. These-are listed in Table 2. These fluxes averaged over angles of 10* to 40* are plotted in the energy distribution of Figure 1 along with fluxes neasured by.several other investigators. 102 The OCR comic diffusa fluxes are in good agreement with the other obser vation* where they overlap fro» 2 to 10 MaV. From 10 to 25 MeV, the UCK upper llalta are consistent with the same alopa. It appears that the line of Dannie at al. (1973) with a slight increase in slope at the higher energies to seat the power law of Fichtel et al. I (1975) of 1.7 x 10-2 (B/MeV)-2-" (photona/ca2-s-ster-HeV) above 35 MeV ! givaa a reasonable fit to all data. 7 i 1 The observations of Hazeta et al. (1975), Trombka et al. (1977) and I Fulcuda et al. (1975) may give a dip I in the energy distribution between, the energies of 0.2 and 1 MrV. If this - • Tim Ei ia real, it could be explained by se им « * IMF» lective absorption of gamma-rays at Пием инЛтт №51 вшш * *. сега their source (Bocchfa et al., 1976). = OF** « * IIW5I - «.Циан И d. (Ifrjl B. Angular Distribution. The - «Stan ft*. tlWI 4 HtMrick tl * lUfJI lsotropy of the directions of inci . o ми* a * U*rai dence of the cosmic diffuse gamma- l{P - • «at « «. I XT}) - • VMiattM « * (IfOI rays was checked over the energy «tMlfRI intervale of 2-3, 3-5 and 5-7.5 MeV на» • «и W0) flcMfl • * (WJI during a given time interval by ob ПсМО • * tWSI serving the fluxes at zenith angles тямм «инотл *f Own ti * ШП1 of 10%-20\ го^-ЗО», 30*.-40' and 40*- ~ • i • "•"! _i_ ' 50*.- Different regions of the celes ai W Ю tial sphere were observed at different GAMMA B" ENERGY ( times'. The times of observation, R.A., 6 and galactic coordinates for' the FIG. 1—The energy distribution zenith direction are given in Table 3. of the cosmic diffuse gamma-rays. A typical set of distributions is Our data from 2 to 25 MeV ere shown in Figure 2 that begins at 15:10 Included along with the data of UT soon after reaching ceiling. The other observers. average for the 4 angle intervals is given by the dashed line at each of the energies. There is no Indication that the diffuse radiation is statistically different from isotropic. We also give in Table 4 for 5 different times the weighted mean, error of the mean and the TABLE 2. COSMIC DIFFUSE GAMMA-RATS IN UNITS OF (photons x 10"'*/cm2-e-eter-MeV) Energy Zenith Angle (deg.) (MeV) 5* 15* 25» 35# 45* 2-3 61 1 14 72 ± 12 3-5 29 ± 9 21 ± 4 13 ± 5 30 ± 10 5-7.5 51 ± 54 9.3 ± 3 8 2.8 ±3.6 4.9 ± 4.7 7.3 ± 3.2 7.5-10 -2.3 ± 3. 0 2.6 ± 3.4 3.6 ± 3.3 -0.2 ± 4.4 10-15 -0.7 ± 1. 5 -0.3 ± 1.4 -0.5 i 2.3 -3.6 ± 3.4 15-25 -0.1 ± 0.6 1.0 ± 0.9 1.7 ± 1.6 103 FIG. 2—The angular distribution of th« comic diffuse gamma-rays during OBSERVATION КЯЮ0 No. 2 ont time period. Flux*» for Incidence energies of 2-3, 3-5 and 5-7.5 MeV arc given. The daahed lines give the fluxes averaged +-t over angle. ! -(»« 10.41 ild* confidence level of the fit to en Isotropic flux for the fluxes at -и eiasi.Kj1 the 4 angles at energies of 3-5 and 5-7.5 MeV and at the 2 angles at energies of 2-3 MeV. From Table 4, we see no evidence for в a non-lsotroplc cosmic diffuse и Z gamma-ray distribution. 0 г-l M«V s O 3-5 C. Upper Limits on Point Л 5-7.5 Sources. Our observations of gamma-rays from the galactic го JO 40 50 anti-center region are ZENITH ANGLE (1M4.I reported In paper 06-7. To avoid Including the anti-center gamma-rays with the cosmic diffuse gamma-rays, the times of observation in Table 3 used to determine the cosmic diffuse fluxes are earlier and later than the times of the aerldlan.passage of the anti-center region. The OCR telescope aperture used here Includes the celestial sphere out to angles of 40* to the zenith. TABLE 3. OBSERVATION PERIOD INFORMATION Interval Elapsed Observation Midpoint Time Depth Zenith Coordinates Period (U.T.) (sec) Ся/cm2) R.A. « I11 b" (- • 14:32 13/5/75 1522 4.6 23b20» 31*51' 102* -27* 2 15:10 13/5/75 2932 3.4 0« 6» 31*51' 112* -30* 3 00:05 14/5/75 2900 4.2 8h49" 32*28' 192* 38* 4 00:41 14/5/75 1398 * 4.7 9*24» 32*30* 193* 46* 5 01:28 14/5/75 3816 5.0 10hl5» 32*27' 195* 56* TABLE 4. MEAN COSMIC DIFFUSE FLUX FOR INDIVIDUAL OBSERVATION PERIODS IN UNITS OF (photons x 10-3/cm2-s-ster-MeV) CONFIDENCE LEVEL OF FIT OF ZENITH ANGLE DISTRIBUTION TO MEAN Observation Energy (MeV) (Zenith Angle Interval; Degrees of Freedom) Period 2-3(20*-40*; v - 1) 3-5(10*-S0*: v - 3) 5-7.5(10* -.<0*; Л 4.8 ± 1.8 1.4 ± 0.5 0.J ± 0.4 532 42Z 27Z 7.5 ± 1. 4 ± 0.4 1.2 ± 0.3 75Z 21Z 71Z 7.6 ± 2. 2 ± 0.8 1.5 ± 0.5 587, 59Z 56Z 8.1 ± 2. 7 ± 0.6 0.6 ± 0.4 41Z 29Z ' 72Z 6.7 ± 1. 7 1 0.5 1к0 ± 0.4 79* 83Z 61Z 104 TABLE S. LOCALIZED 90ГОС1 СГГЖХ LBOTS 95% Ceox idear.e Level On* or mora possible ьоигсса a* In imtta o* (photon* x КГ1/»2-!*-*?) of gamma-rays arc In our aper Flux* in unit* of (pbotona ж 10~'/еиг-»*с) ture moat of the Time. An Isotropic flux during the Thla Experiment (3-25) times of Table 3 sets upper Source a* Flux* limits for localised sources. M31 1.9 5.7 These upper limits are given F*TMUI Cluster 2.1 6.2 in Table 5. It Is noted that Ся A 7.2 21. both our upper limit and the Ся X-l 7.2 21. one from ScMnfelder and Cyg 1-2 2.6 7.5 Lichtl (1974) for Cyg X-2 C« X-3 3.4 10. plus Ся Х-4 are signifi сях-* 2.8 8.1 cantly lower than the fluxes Cjrginu Loop 2.В 8.1 measured by Dean et al. M87 3.4 10. (1973). H82 4.1 12. Cas A 1.9 5.6 4. CONCLUSIONS. From our Tycho SNR 2.7 3.0 obaervatlous of cosmic dif CIA 1 SNR 2.7 8,0 fuse gamma-rays, we draw the Honoceroa Nebula 2.1 6.3 following conclusions: 3C 273 3.0 8.8 1. Calculations show Quiet Sun 2.0 5.8 that the albedo neutron interactiona on carbon give *Flux - /aFT2dE a background gamma-ray flux that is less than auout 20Z of the cosmic diffuse gamma-ray flux at energies less than 10 MeV. These low- values result from the high hydrogen to carbon ratio of 1.82 In the liquid scintillator. 2. The observed cosmic diffuse fluxes from 2 to 25 MeV are compatible with an energy distribution of 2.65 x 10-2 E"2*03 (photons/cm2-s-eter-MeV) proposed by Dennis et al. (1973) at lower energies, steepened at higher energies to meet the slope of Е~2аЧ of Fichtel et al. (1973) at energies above 35 MeV. 3. At 5 different times during the flight for 10 deg intervals up to "0° from the zenith and at celestial coordinates varying from i^ of 102° and b-" H of -27° to tH 0f 195» and D of 56", excluding the region in the galactic anti-center direction, no deviation of the gamma-ray flux with isotropy was observed. 4. From the isotropy of the cosmic diffuse gamma-rays, it is possible to deduce upper limits for fluxes of gamma-rays of 3 to 25 MeV for а number of possible sources, as listed in Table 5. He thank NASA, NSF and 0NR for support and NSBF for balloon launch, flight and recovery. REFERENCES Agrinier, B., Forichon, M., Leray, J.P., Parlier, B., Montmerle, T., Boella, G., Maraschi, L., Sacco, B., Scarsi, L., DaCosta, J.M. and Palmeira, R. 1973, Proc. 13th Int'l. Conf. Cosmic Ray 1, 8. Bratolyubova-Tsulukldze, L.I., Grigorov, N.L., Kalinkin, L.F., Melioransky, A.S., Pryakhin, E.A., Savenho, I.A. and Yufarkin, V.Ya. 1970, Proceedings of the XI -International Cosmic Ray Conference (Budapest, 1969) in Acta Phys. Acad. Sci. Hung., .29, Suppl. 1, 123. Daniel, R.R., Joseph, G. and Lavakare, P.J. 1972, Astrophyв. and Space Sci.. .18, 462. ins Daniel, R.R. and Lavakare. r.J. 1975, ^ сос. 14th Int'l. Con*. Co—le Кат 1. 23. Dean, A., Gerard 1, C, DeMertlnla, C, Monastero, G.F., Ruaao, A. and Scaral, L. 1973, Astro, and An., ?8, 131. Dennla, B.R., Surl, A.M. and Froat, K.J. 1973, Ap. J.. 186. 97. Dyer, C.S., Engcl, A.R. and Quaenby, J.J. 1972. Ap. and Space Scl.. 19. 3S9. Dyer, C.S. and Morflll, 6.E. 1971, Ap. and Space Scl.. 14, 243. Fichtel, C.E., Hartman, R.C., Knlffen, D.A., Thompson, D.J., Blgnaml, G.F., Ogelman, H., Ozel, M.E. and Turner, T. 1975, Ap. J.. 198. 163. Fichtel, C.E., Knlffan, D.A. and Hartman, R.C. 1973, Ap. J.. 186. L99. Fishman, G.J. 1972, Ap. J.. 171. 16?. Fukada, Y., Hayakawa, S.» Kaaahara, I., Makino, F., Tanaka, T. and Sreekantaa, B.V. 1975, Mature. 25ft. 398. Golenetskii, S.V., Mazets, Е.Р., 11'inokll, V.N., Apteka*. R.L., Bredo», M.H., Gufyan, Yu.A., and Panov, V.N. 1971, Aatrophyaical Latter». .9,-69. Herzo, D., Koga, R., Millard, W.A., Moon, S., Ryan, J., Wllaon, R., Zych, A.D. and White, R.S. 1975, Nuclear Instruments and Methods, 123, 583. Heturich, W., Finkau, K., Rothermel, R. and Soamer, M. 1973, Proc. 13th Int'l. Conf. Coaalc Ray. 1, 8. Hopper, V.C., Mace, O'.B., Thomas, J.A., Albats, P., Frye, Jr., G.M., and Thomson, G.B. 1973, Ap. J. (Letters), 186, L55. Horstman, H.M., Cavallo, G. and Morettl-Horstman, E. 1975, Rlviata Dal Huoro Cimento. 5_, 255. Krauehaar, W.L., Clark, G.W., Garmire, G.P., Borken, R., Hlgbie, P., Laong, C. and Thorsos, T. 1972, Ap. J., 177, 341. Kuo, F., Frye, Jr., G.M. and Zych, A.D. 1973, Ap. J. (Letters), 186. L51. Mazets, E.P., Golenetskll, S.V., Il'lnskli, V.N., Gufyan, Yu.A. and Kharitonova, T.V. 1975, Astrophysics and Space Science. 33. 347. Morrison, P. 1958, II Huovo Clmento. ]_, 858. Pal, Y. 1973, X and Gamma-ray Astronomy, edited by H. Bradt and R, Glacconl (Dordrecht, 1973) 279. Preszler, A.M., Slmnett, G.M. and White, R.S. 1974, J. Geophye. Res., 79. 17. Rocchla, R., Ducros, R. and Gaffet, B. 1976, Ap. J.. 209. 350. Ryan, J.M., Dayton, B., Moon, S.H., Wilson, R.B., Zych, A.D. and White, R.S. 1977, submitted for publication to the Journal of Geophysical Research. SchSafelder, V. and Llchti, G. 1974, Ap. J. (Letters). 192. LI. Schttnfelder, V., Llchti, G., Daugherty, J. and Moyano, C. 1975, Ftoc. 14th Int'l. Conf. Cosmic Ray, 1, 8. Share, G.H., Kinzer, R.L. and Seeman, N. 1974, Ap. J.. 187, 511. Trombka, J.I., Metzger, A.E., Arnold, J.R., Matteson, J.L., Reedy, R.C. and Peterson, L.E. 1973, Ap. J.. 181. 737. Trombka, J.I., Dyer, C.S., Evans, L.G., Bielefeld, M.J., Salter, S.H..and Metzger, A.E., 1977, Ap. J., 212. 925. Valentine, D., Kaplon, M.F. and Badhwar, G. 1970, Proceeding* of the XI Inter national Cosmic Ray Conference (Budapest, 1969) in Acta Phya. Acad. ?ci,. Hung., 29, Suppl. 1, 101. Vette, J.I., Gruber, D., natteeuu, J.L. and Peterson, L.E. 1970. Ap. J. (Letters). 160. L161. White, R.S. and SchBnf elder, V. 1975, As trophy», and Space Set.. 38. 18. «,-, Zych, A.D., Herzo, D., Koga, R., Millard, W.A., Moon» S., Ryan, J., Wiiaoia;» R., White, R.S. and Dayton, B. 1975, IEEE Trans. Nucl. Set.. Vol. 9S-22'.-l 605. £ JOb COS-В OBSERVATIONS ОГ GAHHA-RAY EMISSION ПОМ PULSARS The Caravane Collaborat ion Huyeens Laboratory, Leiden, The Netherlands Universita di Milano, Italy Universita di Palermo, Italy Max Planck Isntitut fiir Extraterrestrische Physik, Garching b«i Kunchen, TRC Centre d'Etudes Nuclfaires de Saclay, Trance European Space Agency, ESTEC, Noordwijk, The Netherlands. Theoretical Q ExperimenUl [x\ Do,h Q The gamma-ray telescope aboard ESA's COS-B satellite has surveyed a major part of .the galactic disc, enabling a study of many of the known radio pulsars. The pulsed gamma-ray emission of PSR0531+21 has an energy spectrum represented by (2.4 + 0.6) x 10"7 E~'-l photon cm" 2 s"1 GeV-1. The spectra of the two peaks are similar. PSR0833-45 was observed in Novmeber 1975 and in August 1976. Comparison of these two measurements and the earlier observation by SAS-2 shows a strong variability with time. A search for pulsed gamma-ray emission from radio pulsars for which the period is known to sufficient precision is in. progress and results will be reported. Coordinates: 0G1.2 (Gamma-Ray Sources) Mailing address: Dr. R.D. Wills COS-B Project Scientist High Energy Astrophysics Division European Space Research and Technology Centre Noordwijk, The Netherlands 107 ат.-в I'Hsi-KVniioN:-, 01 Pisrki.n. -..OVHCI.:: O\ HI -M-I.M K _V JJ^MA KM IAI I :. The Глгдудги? Coll.iboi-j t i on ^ Huygens laboratory, Leiden, The Netherlands University di Milano, Italy Universita di Palermo, Italy Max Planck Institut fur Kxtraterrestrische Physik, Carolling bei Munchen, 1H. Centre d'Etudes Nucleaires de Sa'clay, Trance European Space Agency, ESTEC, Noordwijk, The Netheilands. Theoretical Q Experimental [Fj Both Q The galactic distribution of gamma-rays of energy > 70 MeV detected by the COS-B satellite shows, in addition to the characteristic structure of the'general galactic emissiolo a number of enhancements where spatial distributions are consistent with their being due to localised sources. Apart from those which can be identified with pulsars by their short-term variability, sources have been detected at (ln : 78.ЗО + 0.4°, b11 = tl.3° + O.U°), (l11 = 136.5° + 2.0°, bH = -H.50 + 2.0°) and (lH = 195.5° + 0.6°, b11 = +3.9° +"o.6°)! It has not been possible to identify any of these with objects radiating at other wavelengths, results will be presented on the energy spectrum and temporal characteristics of these three and possibly other sources. Upper limits will be presented for the intensities of gamma-rays from various objects which have been predicted to be emitters of high energy radiation. Coordinates: 0G1.2 (Gamma-Ray Sources) Mailing address: Dr. R.D. Wills COS-B Project Scientist High Energy Astrophysics Division European Space Research and Technology Centre Noordwijk, The Netherlands £$№К&3 Elcctrooagnctic Processea In Pulaara under Strong Electric and Magnetic Field Condition» S. Ay «ell, A.llacinliyan, H.B.Ogelaaa Department of Physics, Middle East Technical Univeraity, Локах*, Turkey J.K.Daugherty Universitat Hamburg, I.Inatitut fur Theoretlaehe Phyaik, Hamburg 36 Weat Germany It is believed that pulsara poaeaa hug* electric and Magnetic fielda. However, the electric field ia coaaonly neglected in calculation* of the rate of pair production, a proceaa which ia thought to be greatly important in the radiation aechaniaaa of pulaars. To aaa the effect of the electrib field, the pair production ia calculated for arbitrary electric aad Bagmatio field configurations. The formulae thua obtained arc than applied to palm are. It ia shown that the correction to the "polar gap" height calonlated in the Ruderman and Sutherland model is negligible, although it might be important for the spectrum of emerging photons» I.Introduction All theories of pulsars seem to agree on the importance of the pair production process in the radiation mechanism of pulaara» Shay make uaa of the pair production rate formula calculated by Irber for atatic magnetic fielda. A pulsar, however is believed to poaesa a strong electric •* •* 12 field Е in addition to the strong magnetic field В of the order 10 gauss. It is this Е field that is responsible for the acceleration of particles (2 Ъ) to relativistic energies. Recently Daugherty and lerche ,y pointed out that it might be misleading to neglect the 2 field in considerations of the pair production processes in pulsars. They calculated the photon attenuation coefficient in strong E* and B* fields for the special cases of'" Е . S = 0 and £ x 3 x 0. In the corotating nagnetosphere of pulsars the E* field is perpendicular to the В field so that particles can not be accelerated along the field lines. However, if as in the polar gap model of Ruderman and Sutherland , a magnetospheric gap is produced oomewhere, E.B does not vanish there, allowing particle acceleration parallel to B. In such a situation, where neither E.B. nor ExB vanishes, the results of Daugherty and Lerche can not be used directly. It is possible, however, to make a Lorents transformation 109 to a fraae where tht traanforaed S and В field* ere parallel, perfora tha calculationa thara, and tranafora tha reeulta back to tha original fraae. In aection II pair production coefficient for a general JC and В field configuration ia calculated by aeana of aucceaaiTe Lorenta tranaforaatioa. The reault ia applied to pulaara in Section III. In «action IV tha, calcula tion* of Kuderaan and Sutherland for tha polar gap height are repeated including tha affect of the ? field. The reaulta are diacuaaed iin Section V. II.Pair production rata in arbitrary electric and aagnetic field configuration» Suppose we obaerre, in a reference fraae S, a unifora aagnetic field В in the z-direction and a unifora electric field в in the. y-s plane, aaking and angle d with the s-exia. Consider noe another reference fraaa S' which is Boring along the х-axis with a Telocity be relative to S. An observer sitting on S' will observe the .fields I and В which are related to the original fields' by the Lorents transformation formulae: (1) The condition t x Й * 0 ia satisfied if we have: E'V в; Solving equation (2) for 6, we obtain: ь 7bi;lL~[ ***** j (3) Consider now a photon with a wave vector k in S', with directional cosines n, n , a and energy ш*Щ<- • The observer in S' will observe the aaae photon with directional cosines: and frequency, ^ . ш'~ t(ys-n.f)= ait(i-pn„) , (5) propagating in parallel electrio and aagnettc fields. J ' Xvea in S' it ia diffioult to oalculate the attenuation coefficient for "....•• • • • ' photona in arbitrary directions, but it is possible with a seoond Lorents traaJsforaatioe to go to a fraae S", where the new photon direction Is perpendicular to the оовво* direction of the fields, for aiaplicity we will talee*.?'and B*in the ^'-direction and'the photon ware vector lc'. in the y'-s' 110 plane, мЬЧщ an angle 6 with the a'-axla. Suppose that S" ia moving along the s'-axls with a Telocity f>c relative to S*. Tht requlreaent that th* maw wave vector haa BO component along th* fields yield* iaaediately, р'ж. П) •» Ссав' , (6) Th* frequency of the photon in S" la, then, Со"»«о'<С1-р'с*в')тсо*&.л.в' , (7) and the new field are i--s* , *"--*' (8) low, we can ua* the pair production coefficient given in reference 3, *nd dbtain, in 6"» #J\ , _L -. < ^^ll.Cl%T.fX1+i-0 •* W?5*i (9) for parallel polarization, and, . «-i . <- ^^itft^'X^^^li^F1^J (10) for perpendicular polarisation. Here, - - г -£?• т. > (ID and '-'o ~ "7Г " k.kxXO gauss* o(. and X. are the fine structure cons- tant and the electron compton wavelength reepectirely. The unpolarired > coefficient is: . » ii K" ^» -*• ^ *• ш ~~A (12) Hote -that the above formulas are valid only for X, 4СД» The pair production rate on S* is found by utilising the transformation formula which ia given in reference 2; .ti-^Wf'&=¥,? *>" аз) Subatituting the value* of K and Ъ' w* geti To expreaa Й* in terms of th* quantities in S' we make use of the equation (8) and the fact that fields are parallel in both S» and S". Hence, we can writs 1L •=. — -ж. £> . and obtain th* attenuation coefficient im s» Ъ" Ъ' &^ wltb X.'«J- tiU.' *i . Hers, V«b'&«e ' «4 E.'.E'Si»©'x , г me* ьл *• finally, with a second inverse transformation w* obtain rats of pair prodnctiom la th* original frame 8| Ill Лл* IA->T^. 1 *• (i6) Palm* eqaatloa (1), (3) •** (t) w* oaa express K.' la terse of the original fields I , В and th* dlr*ctiomal cosines a.a.a . The exact result la too loag aad difficult to analyse. fortunately, It la poaalbl* to reduce it to a simple fora for th* eases that are of Inter»at. II.T.applicatiomb>**HPolar dap" model of pulsars la th* polar gap modal of pulaara, proposed by Kudermaa and Sutherland the electric field component, parallel to the magnetic field, in the gap la estimated aai ьп* ~ —z— (17) «here h Is the gap height. The perpendicular component for the electric field can be taken equal to the ralue in the corotating plasma; ^„uSp* , (*.*.*), (18) where Cp is the angle between the rotational and the dipole axes of the pulsar. The maximum value of E..S, obtained for K—,as, is: (EMgW*=*err~ ^«•0'Cb,aSr";-hA«- (19) Е * is maximum for (p =90°; , (A*U,-^~ZL%l0'^ AP"} ""^(20) Thus it will be a good approximation to neglect terms of second and higher orders in _ in the following calculations. Furthermore, in the polar gap of pulsars, photons are radiated by relativistic particles through curvature radiation, hence, propagate almost parallel to the field lines initially. As they proceed the angle between the photon direction and the magnetic field line increases, but the photon travels all the time in a plane defined by the Ваов! Б fields at the point of its emission, so, we can take & »0. To express K.' in equation (16) in terms of unprimed quantities we need to evaluate B^aodl £± . from equations (1) and (40 it can be easily shown that: •*• ba'=^_ and "BI Ь' В Vi+p*i!-T*P-£a / (22) here ^чт^х? *nd ^Чг ^nS • Taki,iE only the first order termo in п: together with the acoumption, n »0, we outa-in: Substituting these into equation (16) w» get: in the first order of J_ . In the limit a, -»0 thin reducea to the reault of Daugherty and Lerche calculated for Б.В • 0 case. IV.Calculation of the polar gap height in the presence of electric field. Euderman and Sutherland, in reference k, calculated a minimum gap height that is necessary for а тасшш breakdown. In their calculation* theymade use of the pair production rate; г K, = 0.aA- Ь* exf ^^tJ J (26) calculated for pure ^ field case, which vanishes for photons propagating parallel to 1?. This way, a curvature radiation photon oust travel a dietanc of at least 105 cm (reference k) before its pair production probability reaches a significant value. Equation (25) however, shows that, if the electric field is taken into account even for those photons propagating parallel to the field lines the probability of pair production does not vanish. For Ь =0, equation (25) reduces to: which is the same as the one obtained by Erber for pure В field with mere replacement of Zg with bx. At this point one can ask whether the presence of the f* field leads to a significant reduction of the polar gap height. To see this we-will use the same criteria that Huderman and Sutherland used tp calculate the minimum height necessary for a gup break down. It is simply to require that the exponent in equation (25) be no larger than kO: "&№+*£?* *« ^ ^° (28) nd he We have, from reference *t, BA-« 1LB * * naximum photon energy •Ixid, s i (1^Л ^£- , Here, p ie the radius of curvature of Лм Лл Vtnc* 'J» * J • 1, the field lines and ' AV» 'ЬД.Ьc ie the accelerating potantial across 7. >.e gap. Substituting these into equation (28) we obtain: 1IJ 4-йЧ.П"'^ (г9) where Jl- is the •inioun gap height calculated for pure В field, and h_ is the corresponding value In combined £ and b fiolda. For £.* « 10 volte/ en, Pi 10 ca and )\ « 5x10 ca, the correction comes out to be: • кв which is not significant at all* V.Discussion and conclusion It is clear from equation (30) that the existence of the electric field in pulsars does not affect the polar gap heirht, though it leads to a considerable decrease in the Bean free path of photons. This is due to the fact that equation (2B) is governed equally by the energy of the photon (+i<»>) and the effective field perpendicular to the photon direction (££). Since W depends on h more strongly than B' , it is the main factor that determines the gap height» In other words, although for high energy photons the mean, free path gets smaller with the addition of the electric field, the gap height necessary for the acceleration of particles that can radiate these photons, comes out almost the same as the one obtained fcr pure В field case. On the other hand the decrease in the mean free path of photons might be important for a spectral analysis since it leads to an increase in the number of pairs in the gap which may result in a significant enhencement of the intensity of the radiation. HHFERKJCKS 1. Brber, T., 1966, Hev.Mod.Phye. ^8, 626. Z. Daujjherty, J.K., and lerche, I., 1975t Astrophysics and Space Science IS, h37 3. Daugherty, J.K., and Lerche, I. 1976, Phys.Rev.D, JA, 3^0 k. Buderman, H.A., and Sutherland, P.3., 1971, Astrophys.J. 16^. 529. 114 PULSED НЮН ENERGY GAMMA RAYS FROM PULSARS S.K. Gupta, P.V. Ramana Murthy. B.V. Sreekantan and S. C. Toawar Tata Institute of Fundamental Research, Colaba, Bombay 400 006, India Tbwwbetl • BxpwimmUl g) •*• Q We have juet completed setting up a larga twelve mirror array to detect pulaed gamma rays with energies greater than 100 GeV from tbe various known pulaars. Electron* In the Ell cascades initiated by bign eaerg r gamma rays emit In the terrestrial atmosphere Сегепкот Radiation which Is focuesed by each 90 cm diameter parabolic mirror on to a fast photomulti- plier. All tbe mirrors are mounted equatorially and driven synchronously. Night Sky Background is eliminated by demanding a last coincidence between the mirrors leaving only the particulate air dhowers as the background. Time of occurrence of each event is recorded to an accuracy of 100 microseconds derived from an accurate timing system backed by a Rubidium atomic clock. Result» obtained upto date will be reported at the Conference. Coortlnite»: OG 1.2.1 (Gamma Rays from pulsars) Mailingiddrm: profeeeor P. V.Ramana Murthy Tata Institute of Fundamental Research Colaba Bombay 400 005 India Mmcesy CONTRIBUTION OF PULSARS TO THE FLUX OK GAUCTIC САММЛ RAYS A. W. Strong and A.W. Wolfrndale, Phytic» Departnent, University of Durham, England. Conclusions about the distribution of cosmic ray nuclei in the Galaxy from studies of Galactic gamma rays rely heavily on thi- assumption that the contribution from unresolved pulsars is small. Thr relation ship between gamma ray emission and radio emission for detected pulsars is used to estimate the likely contribution from unresolved pulsars. It is shown that the best estimate of the pulsar - induced fraction is 5-10Z 1. Introduction. Two pulsars have been identified as у~>"ау sources with certainty (Vela and the Crab) and two others (1747-46 and 1818-04) are very likely (see the work of Fichtel et al., 1975, Thompson et al., 1976 and Ogeluian et al., 1976). Insofar as the majority of pulsars in the Galaxy have not been identified in the radio band it is conceivable that these pulsars contribute significantly to the y-ray flux from the Galactic disc. This contribution could, in principle, be large and its neglect could invalidate the conclusions commonly made about the source of the Galactic flux, i.e. that it arises largely from H°- mesons produced in cosmic ray nucleon - I-S.M. nucleus interactions. Earlier versions of this work have been eiven by the nresent authors, Dahanayake et al. (1976. 19771: a similar studv. but with different conclusions, has been made bv Higdon and Lingenfelter (1976). 2. Relation between Radio and y~ray yields. Both forms of radiation originate, presumably, from the movement of pulsar-accelerated electrons in the immediate vicinity of the object and, a1though the energies of the electrons are no doubt very different in the two cases,there is reason to expect a rough and ready correlation of intensities. Figure 1 shows the results for !„(> 35 MeV)and IR (400 MHz). The Y-ray data are from Ogelman et al. (1976) and, in one case, the Caravene collaboration (1976) and the radio intensities are those given by Taylor and Manchester (1975). As will be seen, in all but four of the cases the y-ray intensities are upper limits so that the demonstration of linearity between ly and IR is difficult. It appears that, although for a fixed IR there is a variability of ly from pulsar to pulsar, the average trend is close to the line indicated. The variability could well arise from geometrical beaming effects and is not of great consequence here. I It» —i r—1 » ' I 4- 5 - У 4 id v T • ' /; 1ob- "' ,"''^ v/ . I A _. I Ю Гл(4ос/ИНх) (*J>) Figure 1. Fluxes of Y~rays above 35 MeV and 2o upper limits (denoted T) for pulsars given by Ogelman et al. (1976), The diamond is derived from data given by the Caravene collaboration (1976). Accepting the validity of the argument, the mean ratio of -y-ray to radio intensity is IY ( > 35 MeV) — - 2.4 x 10"6 cm-2 s"1 Jy"1 2.1 IR ( 400 MHz) ' The ratio of the Y-flux above 100 MeV to that above 35 MeV-is ^0.34 (Fichtel et al., 1976) so that I ( > 100 MeV) -L- — - 0.8 x 10"6 см"2 s"1 Jy_1 2.2 IR ( 400 MHz) 3. The integrated pulsar-induced y-ray flux. Seiradakis (1976) has used the Jodrell Bank survey of 51 pulsars to obtain distributions in luminosity, period and spatial position. Using these data and equation 1.2 gives a -y—ray emissivity near the Sun of q ( > 100 MeV) - 0.3 x 10"26 cm-3 s" 3.1 This value is 5% of the total local emissivity of 6 x 10~26 cm"3 s-1 (Strong and Worrall, 1976). 1 11 to3 т—r -4 T ->]Q0) 't.' ••• • • • - л • -5 : (WVf -6 Xs from Pui's-irs -7 J i 1 i J L ,.j L i L # 4 ISO 120 6C r0 -60° -/20 -180° Figure 2. Distribution of чг-rays above 100 MeV it: longitude (for |b| < 10°). The data are from the SAS II experiment: Thompson et al. (1976) and the pulsar contribution is the smoothed predictions of the present work for a local emissivity of 0.3 x 10"26 cm-3 s"1. The longitude distribution of y~rays over the range |b| < 10° has been calculated using Seiradakis' results on the z-dependence of the pulsar density,with the result shown in Figure 2 where the SAS II data are also shown. The fraction of the measured flux above 100 MeV varies from 7% towards th° Galactic anti-centre to ^ 4% for longitudes below ±30°. An independent analysis has been made using the pulsar survey of Hulse and Taylor (1975). These authors examined the pulsed radio flux (periods > 0.033 s) from the region |b| < 4° and 43° < 1 < 58° under conditions of high sensitivity (lower limit 1 m Jy). The total radio emission from the region is 'ь 800 m Jy of which 902 comes from pulsars of flux density above 10 m Jy so that the contribution from very faint pulsars is negligible. Using these data the radio yield per radian of longitude is dip 4.7 Jy rad-1 for |b| < 10° 3.2 dl (a correction was applied for 5°<|b| < 10°, based on the known z-dependence of the pulsars in the direction in question). Ite Gomfrinatioa with equation 2.2 gives dl ( > 100 MeV) 1 _I - 0.37 x 10"5 cm-2 %-l red-1 3.3 dl valid for the rang* 43° < 1 < 58° * lb* measured у~г*У ^\UK In thi» region from the SAS II experiment 2 1 it (3±1) x 10-5 cm" i" rad'l i0 that the pulsar contribution ii (10±3)X. Inapaetion of figure 1 show» that, using the previous method, the average coatribution is i> 61, in good agreeaent, for this longitude range. 4. Conclusions. It has been demon»trated that pulsars probably contribute lass than 10Z of the Galactic v-ray flux; it seeas unlikely that the contribution is as high as the 30Z of Higdon and tingenfelter. Having said that, the possibility of surprises remains; thus, if the spread in the IY/IR ratio is large and if there are pulsars of rather high y-ray emissivity but low radio amission then the picture could change. Acknowledgements. A.W.S. thanks the Commissioners for the Exhibition of 1851 for the award" of a Fellowship. References. Caravene Collaboration, 1976, Symp. "The structure...y-rays", Goddard Space Flight Centre, Greenbelt, Maryland, X-662-76-154, p. 52. Eahanayake, C, Dodda, D. and Wolfendale, A.W., 1976, Symp. "The structure' Y-rays", Goddard Space Flight Centre, Greenbelt, Maryland, X-662-72-154, p. 126j 1977, Mon. Hot. Roy. Aatr. Soc. (in the press)'. Fichtel, C.E., Hartman, R.C., Kniffen, D.A., Thompson, D.J., Bignami, G.F., Ogelmsn, H., Ozel, M.F. and Turner, T., 1975, Ap. J., 198, 163. Higdon, J.C. and Lingenfelter, R.E., 1976, Ap. J. Lett., 208, L107. Ogelman, H., Fichtel, C.E. Kniffen, D.A. and Thompson, D.J., 1976, Ap. J., 209, 584. Seiradakis, J.H., 1976, Syrnp. "The Structure and Content of the Galaxy and Galactic Gamma-Rays", Goddard Space Flight Centre, Greenbelt, Maryland, X-662-76-154, p. 299. Strong, A.W. and Worrall, D.M., 1976, J. Phys. A., 9, 823. Taylor, J.H. and Manchester, R.H., 1975, Astron. J., 80, 794. Thompson, D.J., Fichtel, C.E. Kniffen, D.A., Lamb, R.C. and Ogelman, H.B., •1976- preprint, Goddard Space Flight Centre, X-662-76-94. 119 САЖА RAYS ГЮМ THE DIRECTION OF THE GALACTIC CEFTU A.W. Uolfendala and D.M. Worrall, Physics Department, Univcraity of Durham, England. The SAS II gamma ray re*' Its of Thoapaon at al., (1976) «how a distinct peak toward» the Galactic centre. We have examined the possibility that there are aourcea in the molecular hydrogen ring round the G.C. from which protona and other nuclei, and their aaaociated aecondariea, interact with the local gas to produce pioni and hence Y-ray». The relationahip of the generation rate of cosmic rays to the local gaa density is discussed. 1. Introduction. The SAS II measurements of y-rays above 100 HeV as a function of Galactic longitude (Thompson et al., 1976) show a distinct peak towards the Galactic centre. A number of possibilities spring to mind; here we consider a likely one - that the тг-rays arise from cosmic rays interacting with the gas in the ring of dense molecular hydrogen clouds some 300 pc from the G.C. In the analysis we assume that the cosmic rays are produced and trapped in the ring and the factor by which the cosmic ray density is higher than the local value is calculated. The. par tide generation rate : also derived and its significance is assessed. \ 2. Conditions near the Galactic Centre. The. presumption of a ring of dense H2 clouds comes from the.analysis by Kaifu et al., (1972), Scovilli (1972) and others of molecular line observations by McGee (1970) and Solomon and co-workers (e.g. Scoville et al., 1974). The mass of the rii is estimated as *ь 5 x 10' Щ and the volume % 10? pc3, giving an average density of ^ 220 H atoms cm--'. In addition there is a small contributiot from neutral hydrogen. 3. y-Ray Emission. In earlier publications (Wolfendale and Worrall, 1976;1977) we calculated the emissivity of y-rays above 100 MeV from cosmic ray proton and nucleus interactions to be a (TT) - 1*9 x 10"24 f cm-3 s"l where f is the factor by which the injection spectrum at the G.C. is higt than the local value. -. Ia..£hi.8 .must be added the contributions from the interactions of primary and secondary electrons. If the primary.electron injection spectrum is also higher at the G.C. by the same factor f the contributit are: bremsstrahlung 0.85 x 10~2* f cm~3 s"1 Inverse Compton 0-03 x 10~24 f cm"3 s~* 120 With the eaeumed |M«ttT of th* production regie* at the C.C as indicated earlier the expected Intensity of т-raya above 100 Ha» la _i 2 -1 ly - 6-7 ж 10 f em" a . Analysis of tha SAS II data by ua indicates I • «.7 ж 10"* са-2 а-1 ao that a valua of f ч lOOlis require* "for cootiatancy. in fact, this ia probably an underestimate ay a factor of 4 ainca the calculation aeeumad a lccal lifetime againat eacepe of * 2.S ж 10*y and racant astimetea give values higher by a factor of % 4 (aaa tha summary by Webber at al.. 1977). 4. tadio Emission. Bvidanca for energetic alactrona and tlgnifleant •agnatic fialda COMBS from tha obaervatioa of a peak in tha longitude diatribution of synchrotron amission. For example, the summary of Cream (1974) at 40S № suggests a brightneaa temperature of 100* from the central region over a solid angle of * 4 x 10"*sr. Tha work of lirabayashi (1974) indicates that * 75Z is non-thermal, leading to a 40B MBs flux af 1.47 x 10"11 cm"2 »"laV~V and, with tha geometrical factors considered earlier the corresponding emissivity ia 6.1 x 10~6 cm"* s"1 eV~l. Calculations have been made for the 408 № emissivity expected for various valuea of f with the object of deriving the effective mean magnetic field at the G.C.: В . With f - l0Ot Bc «Up gauss ia required for consistency with observation. 5. Discussion. It Jta_a been shown that the injection rate required at the G.C. ia at least 100 tines the local value and probably as much aa 300 times. This figure is not vary different from the factor by which the average gas density is higher than the local value. Insofar as. on*, might expect the denaity of cosmic ray sources to be vary roughly proportional to local gas density (assuming implicitly a Galactic origin for the particles in question, i.e.those in the range 1-10 GaV) tha agreement gives some measure of support to the Galactic origin hypothesis. If the gas and magnetic field were in equilibrium and if tha cloud velocities were similar to those locally then Nn * B^ would be axpactad. In fact we see that B2 is increased by * 14 cf 220 for Ng. Likely explanations He in terms of a lack of equilibrium (the lifetime of th* ring is probably only * 10°y) and differences in cloud velocities in the two locations. References. Green, A.J;, 1974, Astr. Astrophy». Suppl., 18, 267. Hirabayashi, H,,,1974, Publ. Astr. Soc. Jap., 26. 263. Kaifu, N. et el., 1972, Mature, 238, 105. McGee, R.X.* 1970, Austral. J. Fhya., 23, 541. Scoville, N.Z., 1972, Astcc-phys. J. Lett., 175, L127. 121 Scovillc, N.Z., Soloaon, P.N. and Jaffarta. К.1., 1974, Ap. J., 1*7, LbJ. Thoapaon, D.J. «t •!., 1976, Proc. Гим Kay Syapoaiua, Goddard S.r.C, 1. Webber, V.*. at al., 1977, Univ. of Haw Haapahire. Durhaa, U.S.A. (preprint). Wolfendala, A.W. and Worrell, D.M., 1976, Nature, 263, «82; 1977. Aatr. aad Aatrophy*, (in tha praaa). WS>Q£55 SUKUOVA SUKLL аЖКЛШ ТОЧНО FULSUtS A3 A UAJOU, жштмю in попа S&UKK. Y.SJBerealneky and O.r.Prllutaky Institute for Nuclear Baeearoh, Academy of Solano*» of the U33X. The generation of Jf -radiation in a dans* supernova, •hall around young pulaar la oonaldarad. 1Ъа deteotl- on of suoh aouroa in. У -rays with S j £. 100 RaV la poeaihle duxdng~ 5 aontha after a supernova, azpioai- on at a diatanoa up to 10-15 Mpo* V* suggest tha pro- fiu of observation of supernova exploeiom In our Oa- laxyi It la to' Ъа triggered by dataotlon of neutrinos from gravitational oollapaa and inoluda tha aaaroh 1c the same dlraotion for high energy nautrlnoa ( a day 18 after explosion), for neutrana with I>10 aV (a fortnight after explosion) and for 100 NaT and 100 OeV V -raya ( a fortnight and a month after ex - ploaion, respectively). ТЪеае observationa will pro - vide ua tha data on tha rata of o.r. production by a young pulaar. I. Expanding supernova (SB") shall filled with oosaitt rays (o.r.). Tha< SH explosion is usually supposed to leave behind the expanding shell and a rotating magnetio neutron -tar (pulsar), whose nsgnetio- dipole ra diation can aooelerate the partLoles up to the highest energies observed in e.r. (Ounn and Ostrilcer 1969, drawing andiHeinismann 1973 etc). The expanding shall filled, with o.r. can be characterised, by four spealfic moments (ages) (Beresinsky and Prilutsky 1976, Beresinsky 1976). I) The moment t-w- sinoe which on the decay time of charged piona with Lorents-faetor Г beoomes. loss than tha time between two nuoleax colli sions t ^--(•Чй^Гг^М^ 2) Th* moment t sine* whioh on the shell become» transparent for neutrons: 3) Ям momant t у sine* which on tha shall be acmes, transparent f on X-quanta* 123 *,--(5Й5=Г=»*'(Н/Н.У"« 4) Ле- aoaent t». •loo* vhioh on adiabatlo energy, loeaaa begin to do- •.iaate over tboaa due to nuoleex oolliaionat In the Iq.. a (I)-(4) K ia the asaa of tha ahalX («hiob ia aaauaed. to Ъа hydrogen and. equal to Dl© ), u-I.ICroa/a ia tha velocity of the expan- sion, At t < tT ( Г ) the charged piona with Lorants-faotora. higher than. Г do not daeay on the nuclear path length and the hadronlo: oaaoada deve lops in the shell.. It terminates in oharged pione with LorantE-faotos leas than Г whioh decays result in the soft >) -radiation. As the shell expands the free path-length inoreases, pione with higher energies decay and therefore the neutrinos with higher energies, up to the maximum at t^t.j (f^), are produced» f At t %, t neutrons generated; ia pp-oollisions with Lorentz-faotors m 3 a f >ГС *Jtn -2.6.I0 ( Tn -0.93.I0 s) escape from the shell while those with Г^ fc decay within .the shell into proton» and are confined by magnetic: fields, (the mechanism of o.r. exit from dense SH-shells, Beresinsky and Priluteky 1977b). 0 ** "t >. "*» - У -quanta with Ev < 100 OeV, produoed. through If - de cays, freely escape from the shell and it becomes a volume fc -radiator. But unless the temperature T of the shell is less than 4.10% the shell 124 remains opacue fcr Jf -quanta with Ij >, I00-IOU0 OeV due to Vbe pels production 1 * •>"***•«" on t)i* thermal photons. TV» oondltloa of noo- transparenoy la Лу<. ut «bar* the free path-lenfth X> of a photon la given at «Го*/в *Т» I by the formula (Oould and Sohradar 1966) i A^(lVr/etx>(Vm«c)(E,/m,cfc)VV (5) whara -^ •m2©?*/! W, dv -I/I37 end. • la tha maaa of electron, shea, the photosphere shrinks and tha taaparatura of periphery falla down te» T^. A.ioh. tha Jf -radiation, with В% £, 100 OaV originated la. tha outer part* of the shell oan eaoape fro* it because tha oolliaiona with, ther- aal photons take plaoa at saall ancles between the diraotiona of both phot one. But if tha option! radiation of tha shell is produoed, due ta the nuclear oolliaiona of o.r., the tharaml photons are distributed uni formly within the shall, and therefore it oan remain opaque for the 8 -quanta with E, £, 100 OeV till t K tt. At t> to. the adiabatia energy loses* of o.r. dominate over those due to nuclear collisions and. the generation of all the radiations drastically daoreaaes. 3. Flux of У -radiation with By £ 70 MeY. We assume the power-law Bpeotrua of protons injeoted in the shell with low energy out off at E»E . Ibe total number of aooelerated protons' with dimensionlesB energy E £ -V 0 In the shell at the moment t « to, ia (Beresinsky 1976)< N^tHfrU/6)(i*tfty1'tVti-1 w*°, (6) where L is the initial luminosity of the pulsar, Л la the fraction of the energy transferred to the aooelerated part ides, *£ ia the time of the pulsar braking and n 'ia the braking index (n-2 for braking due to a magnetic dipole radiation and n«I for gravitational one). The number of 2f -quanta emitted by the shell per lseo at t <.< t^is <5wO»XCn»-CdEW,CE,l)'fBi(E)(^© , (7) 5 5 where n(t)-3Jv(4Tu' t- nH) is the atomic: density of the shell, ($^ (E) and Yif (E) ale the oroesi-seotion of 1Г -production, and. their- mul tiplicity and E,, is the threshold of pioni production» We assume that EQ^ Е and use for "ущ^ Steoker formula (Steoker 1973)« i.aT^V'^ce2 at 0.4 i В S 0.7 8 I -.4 10-21*Р^<шг at Е > 0.7 I2S V.OTO :•- in Go v. ln»»rlia£ (6) »nd (S) into (7) «re derive for li -K and t= 1.6 a>ftx- fcU-OMc xUp» <&t _.Mi0'*»U < sec-' (9) E-(S X) (B) (B) 1,3 IO 26 2 where °k " & * * ** ' **• " * " " where L la ia ezga/a. The у -apeotrun hea maximum at В v St 70 NaV and -If has a form £ at higher energiea. 4.Detection. If the thraahold flux for detection of ~g -quanta witfc 1 44 Ky 5^ 100 MeV ie vf^~ IO^oa^a" and >Lo~ I0 exga/a, the maximum distance at which the considered souroe oan he detected ia 2 r- ]_*PV (** )/4Wti "2 ^ ~ 15 Mpo. Ihximum flux ia to he dateoted >t t « t| after SN explosion, than the flux diminishes aooording to. (9) up to t s tfc and later on, it begins to diminish drastically. If pulaar brake a due to magnetic dipole radiation, then n»2, X ^ I.IO'e and. If -luminosity of the shell practically does, not change during the in terval to, -t j • In the case of braking due. to gravitational radiatioa "C «To can he small. ( "Cq1^ Ю a, Ounn and Oatrikar 1969) and at the moment t -g pulear luminosity is t j /X-a, times less, than the initial one. Till now we discussed the deteotion. of extragalaotio sources. But much more information oan he obtained in the oase of SN explosion let our Galaxy. The ")f -observation in this case oan reveal the minimum quantity of c.r. in the dense 3N shell compatible with the pulear production, of the Oalaotio o.r. ( ЛЬ~ 5.10 'ergs/е for SN explosion near the Oalaotic Center). The "У -observation: with E^ %, *&> G»V will provide ua with in formation oit the high energy o.r. in the shell, while the fluxes, of neutrons; (E £. 10 eV) and neutrinos (E > 10 eV) carry the information on то 20 c.r. of тгету high energies up to 10-10 «V. We shall discus» briefly the possibility of deteotion of high energy neutrons from the dense shell. (Sexezinsky and Prilutsky 1977a). The total number of neutron» with the energy higher than £. emitted by the shell during the time; tfc - t , in the oaae of a power-law spectrum of accelerated protons is* 126 At *g -1.6 (K -0.12) and >L ~ 10 ^Tgu/ш the nucb«r or allied neutron» with K>10 eY ( Е > IO5) IB *T^~ l.IOiT and Vhe flui at Uie Birth ia «v 10 кш~ , while tha e.r. baokgrouad '.n tho aolld лг-^ie 52.~5.IO »r fop tha мм tiae t« - t la л.1.5 kn-2. If oulaar du- ^- n ring Nti, aooeleratea partiolee aalaly to tba higheat energy t-K (China and Ostrikar 1969, Graving and Helntiaann 1973) the number of «ait- B tad nautrona draatioally enoreasesi ty ~Л1>0*а. / • The high епеяку neutrino radiation of tho ahall ia disouesed in the pa par by Bereainaky and Prilutsky (1976) and Baresinsky (1476). Tba high energy nautrinoa oan ba dateoted due to nuolear-eleotroeagnetio oaaoadaa ia. the eointilla- tion tanka for detecting nautrinoa fro*: gravitational oollapaa and by generation undeground muona. Tba high energy neutrinoa can be dataotad alao by Davis.' installation. The high energy neutrino observation, is va ry important beoauae it is generated during the earliest ages of a pulsar as a o.r. aourea and in the ease of a vary rtpid braking of a pulsar it is the only radiation which survives. All the proposed program o? observations must ba triggered by detecti on, of 10-20 №7 neutrinos from gravitational oollapas (Doaogataky and Zatsepin 1965). In the nex-fc two years several detectors, of this type will operate in different countries. The search of neutrino, neutron, and ~6 -radiations is to be started after the interval» of time t^, , t and t у respectively in the direction of the discovered SV-ezplosiosu The observations wiH. give the ooapleto information on the pulsars as the main o.r. bouvoee in our Oalaxy. References. Beresinsky V.S. and Prilutaky O.r. 1976 Proa.Int.Conf ."Heutrino-76" Aaoher Beresinsky V.S.1976 Prooj.euramer workshop "O0]UUn>-76H (Bd.A.Roberts FOAL, Batavia), 246 Beresinsky V.S. and Prilutsky O.P. 1977a Plana Astron.Journal 3, H4.I52 Berezinsky V.S. and. Prilutsky 0.7. 1977b Pisma Astron.Journal 3, H6 Domogatsky 0.7. and. Zateepin G.T. 1965 Proo^9Int.Conf .Coan.Ray,London 2,1030 Ferrari A. and Trussoni E. 1974 Phya. Lett. 42A, 345 Gould R.J. and Sohredsr 0. 1966 Phys. Rev. Lett. 16, 252 Growing H. and Heintemana- H. 1973 Phya. Lett, 42A,345 Ошш JJ.K. and Ostriker JJ.P. 1969 PhyB.Rev/.Lett. 22,728 127 A POSSIBLE COratlMTHW 07 DISCOT1 30ШСЖЗ TO THE DIFHJSI QAfeHA-lAY DISTtTBUTIOI 0ТП GALACTIC LOVOITOOtV A.A.Stepanlen Crimean Astrophysloal Observatory, Crimea, 33*413. OTSE The possible contribution of discrete sources la the galaotie Centre region to the distribution of /-rays with energy >108 ev over galaotio longitudes is conside red. It Is noted that the i-r*y emission of the sources la variable according to botn the data of high-energy ( >10в ev) and the data of very high energy (1011-1013 от) r-amission. The conformation of high-energy eleotron distribution отег the Galaxy to the nuclear component dis- tribution deriTOd from the data of /-ray oaiaaion amy be attained if we take into account the discrete aourcea contribution* It waa ahown that the discrepancy between the bal loon data and satellite data amy be caused by the exis tence of point sources and the variability of these sour ces. As it is known the measurements by 030-3 hare shown that a the galactic disk emits y-rays with energy ~10 ет. Calcula tion of the /-ray flux produced by interaction of cosmic rays with interstellar gas results in good agreement with measure- aenta [1] * Thus it is one more evidence of the existence of a nuclear component of cosmic rays in the Galaxy* The density of cosmic rays necessary for providing /-ray flux from the Ga laxy ia the same as local density for all directions exoept the Centre region of the Galaxy: -30°< ^J<+40°. The flux from this region is 3 to 5 times larger than expected. It is clear that flux excess may be provided either by increasing the cosmic ray density or by increasing the density of interstellar gas or electromagnetic radiation. The existence of a number of point sources in the Centre region cannot also be excluded. Since up to recent time there was no data about point sources except the Crab nebula and Tela X, this possibi- ):л lity waa not ooneldered extensively. As it m found in [2], the high-eaergy electro* eemalty . In the galaotic Centre region la no more than twioe mm high ee that in the anticentre direction. If we una* that the ratio of nuclear to eleotron eoaalo rays la ooaatant throagh the Ga laxy the oonoluaion aade for eleotrona alao helda trae far the nuclear ooaponent. The density of neutral hldregen la knew* from redlomeaaureaent data. Therefore It la not poaaible te explain 030-3 data without speculative auggeetion. However recent radioaeaaureaente aake it poaaible to la fer the amount of the molecular hidrogen In deaae olooda that are plaoed la the 5 -kpc apiral am. Uaing theee data Stacker et al. [з] пате calculated the dependence of jf-гау flax en ga lactic longitudes and obtained good agreaaent with aoaavreaaats [1,4] . But it way be ahown that the cosmic-ray denaity depen dence on the diatance froai the Centre of the Galaxy ueed la [3] ia in contradiction with raault of paper [2] concerning the abo ve aentloned ratio of the average deaaity of electrcna la the direction of the Centre to that of the anticentre of the Qalaxy. For these reasons it ia worth conaidering the auggeatioa that discrete /-ray sources can contribute to flux from the Centre region of the Galaxy. Some arguments in favour of thia suggestion are presented below. The first reports that claimed to reveal the 10 ev and 12 10 ev point /-ray sources were published a* far back aa 1965 [5,6]. However the statistical significance of theee mea surements made by balloon and on the ground by the Cerenkov tech nique was not very high. The variability of the sources aade these results uncertain. Moreover the data of 0S0-3 and SAS-2 have not confirmed the existence of the point aourcea revealed earlier by balloon measurements. Fluxes of very high-energy y-rays from point aourcea -with high statistical significance were detected using the Cerenkov technique in 1971-72 [7,8]. Recently some of the point sources were revealed by satellite measurements. The /-ray emission in в 12 " the range from 10 ev to 10 ev was revealed for three discre te sources: the Crab nebula, pulsar PSR 0833-45 and X-ray sour ce Cyg 1-3 [9] • It is seen (see [10] and references therein) i:« that the variability of theae eourcea ia complex. Tbt pulsare NP 0532 and PSR 0833-45 bava pulaad »-emission. Tha total flux of both aourcaa varlaa with tljia. Cyg 1-3 haa a periodic component. Tha /-emission laata 15 minutes aaoh pariod (4.8 hour). In addition to periodic emission, aometiaea aporadio r- eaiasion may ba observed. As was ahown (aaa detail in [to] ) tha diacrepanoy between tha observations made by balloon and satellite equipments might be explained if the variability of /-emission and the diffe rences of the effective areas of the detectors are taken into account. low it is obvious that the J-ray flux excesses in the direction of the Crab nebula Vela X and Cyonua regions are connected with discrete sources. A rough estimation of the re lative contribution of these sources to t-emission from the region 55°< £*<• 290° gives the value 25*. It ia natural to sup pose that the contribution of discrete sources to У-ray flux from the Centre region is no less. The comparison of the у-ray data with radlomeasurement data led us to conclusion that about 40% of the flux from the Centre region is connected with discre te sources. Finding of these sources in the galactic Centre direction is hampered by the next reasons: a) compbxity of the r -flux variation, b).faintnees of the sources because of the large distance (note that both PSR 0833-45 and NP 0532 are relatively near and they would not be detected if they were > 5 kpc distant), c) low angular resolution of detectors, which results in su perposition of the sources. It is Interesting to note that it is difficult to explaia the discrepancy in data of SAS-2 and 0S0-3 for the у-flux di stribution from the galactic Centre region statistical fluctua tions. So the conclusion is made that part (~4056) of the /-ray flux from galactio Centre region is emitted by discrete sources. This suggestion has the next advantages: 1. It is easier to explain the redueed to 6*10"' quanta. —2 —1—1 cm е rad diffuse emission from galactic Centre region by .//'"-decay through interaction of cosmic-ray nuclei with ambient 1М 2. It olarifiea the oauae of the dieerepanoy between aate- llite meaeurementa end balloon aeeaureMnta, ooMoernimg Alaerete eouroea. SOM hundreds of dieorete eourcee auoh as the Crab aebmla are neoeaaary to give aa anion м 40Х of tha flux «a la produced by th* galactic Centre region. It auat ba noted that total eaia- aion of X'1** with energy >10 ет la equal to: for the Crab nebula - 2Ч03' erg a"1, for Vela I - 2«1034 erg a"1 end for Cyg X-3 - 2«1037 erg a-1. Referencea 1. V.L.Kraushaar, G.V.Clark, G.F.Gemlre, R.Borken, P.Higbie, V.Leong, T.Thoraoa, Aatrophya.J., 1972, 177. 341. 2. R.B.Daniel, S.A.Stephens, Space Scl.Rer., 1970, JO., 599* 3* F.V.Stecker, P.M.Solomon, V.Z.Scoville, C.B.Ryter, Aatrophya. J., 1975, 201,- 90. 4. D.A.Kniffen, R.E.Rartman, O.J.Thompson, C.E.Piohtel, Astrop- hys.J., 1973, 186» 405. 5. CD.bong, B.McBreen, H.A.Porter, T.C.Veekes, Proo. 9 Intern. Conf.on Cosmic Rays, т 1, London, 1965, p.318. 6. J.G.Duthie, R.Cobb, J.Stewart, Phys.Rer.Lett., 1966, ^J.263. 7. H.F.Helmken, G.G.fasio, B.O'Mongain, T.C.Veekes, Astrophya.. J., 1973, B, 245. ^^^ th 8. B.H.Vladialrsky, A.A.Stepanian, V.P.Fomin, Proc.of-the 13 Intern.Conf.on Cosmic Rays, т 1. OG session, p 456, Denver, Colorado, USA. 9. A.M.Galper, V.G.Kirillov-Ugriumov, A.V.Kurochkin, B.I.Luch- kov, Yu.T.Yurkin, Plana v ZKTP, 1973, J§, 217. 10. A.A.Stepanian. iKv.Krimsk.astrophys.obe., 1977, 5Z» 99. 131 ПВЮШОАЬ GAsaU-BADIAKO" nm m Х-EAT аошюж ого>*. Gainer A.M., KirllloT-Ugryomov Y.Q., KuroohKt» A.Y., Lelkov I.Q., Lunbkov Gamma-radiation of Сук Х-Ъ with the energy above 40 MeV «as detected la three balloon flight* (Oaloer et el, 197?! 1973) and waa found to be pulsing «1th a period corresponding to the X-ray and Infrared radiation. A wldeapertore gamma-telescope «1th spark chambers (Galper •t el, 1974) was launched three tlasa on sltltydo balloons 1 10/12/72, 07/10/74 and О7/ОЗ/76. The altitude height waa 6+ 10 g/cm of residual atmosphere. Gamma-telescope registered gamma-quanta with energy sbore 40 MeV, their detection area being 250 cm , s total aperture 80», whloh allowed to obeerre the source each time during 5 hours. Averaged on spectrum the angular resolution of th* Instrument was about 5*. In accordance with ansular resolution the number of detected gamma-quanta with coordinates deviating from th* direction to Cyg 1-3 not more than 3*5* war* determined. Be sides the expected number of such gamma-quanta due to stmosphe- rio background has been estimated (Kurochkin et al, 1974). Statistically essential excess in number of gamma-quanta was obtained only 1л the first flight (October 1972). Excess amount was four standard deviations from bscJ«*— nd that corresponds to the probability of its accidental appearance 5«10**^. In other two flights statistically essential excesses in this ares has not been observed. To determine a possible radiation -periodicity for any parti cular ganma-quanta the phase of X-rsy period (Paraignault et al, 1976) corresponding to its detection time was determined. Tigure shows the phase distribution of deviation of detected events from expected ones. The probability that the observed excesses ere sceidentsl is deposited In the axis of ordlnates. There were observed statistically essential excesses in phases 0.1* 0.3 in all three flights. vThe probability that the detec ted peaks have an accidental origin, in accordance with Poisson distribution, is 1.3*10"*5 for the first flight, 3.Ю"3 for the second and 1*10 for the third. The observations of gamma-radiation from Oyg X-3 with the energy above 2*10 eV were carried out during four years (1972 -197З) using the registration of Cerenkov's radiation from MJ П 'I * •>'! SAB with the detectors of the Crimean Astrophysical Observatory (Stepanjan et •1. 1971). The time of the source observation was about -150 hours. The treat ment of the observed data allowed to discover a pe riodical component with a high degree of reliability and to improve the accuracy of the period T=ofl 99684+ 5 7 I1H +С&ХХЮ02. It should be n 2ik is' noted that the presence of periodical component is observed every year and the value of. its flux does not depend on the time (Stepan- yan et al, 1976). The posi- •'» 0 « '* « * tion of periodical gamma-radiation pulse with the energy above 12 10 eV marked with an arrow on the figure. Thus the experimental data unambigiously pointed out that a gamma-radiation from Cyg 1-3 in two energy intervals (£4*10 oV 12 and*2«10 eV) has aperiodical character. This radiation pe riod agrees with X-ray one..Because the statistically essential excesses deal with phases 0.1» 0.2 It is possible to say that the radiation is of pulse character, the pulse length during one period being about 1 hour for B*10 eV and 15 min for В » 12 10 eV. The fluxes of the pulsing component of gamma-radiation with ((he energy 40-lleV_were (11+3).l0""5cm~2sec~1, (6±3>10~5 cm ee© and (8+6)»10~:,cm""2eec"^ for three flights correspon dingly. The gamma-quanta flux with the energy»2•1012eV avera ged on four years of observations (1972-1976) was LJ^IO"11 ca sec . The pulsing character of gamma-radiation with the period 4*6 hour was also discovered by the satellite SAS-2 (Lamb et al, 1977). \U Reference*, iller H., Rode* p». 1973* Nature Rhya. Sol., £*£, 40. Beoklin Х.Ж., Havfcina P.J., Maaon 1.0., Uatthewa X., •eugebauer G., Paokaan D., 8anford P.«., Scnupler B., Stark A., Vynn-*illiaaa G.G., 197*, Ap. J. Lett., 1«g, L119. Galper A.M., Kuroohkla A.V.( Leikor N.G., Luohkov B.I., Turkln Tu.T., 197*• Pribori i Teoanika Bxperlaente, 1, 50. Galper A.M., Urillov-UgryuaoT V.G., Xuroohkin A.V., LuohkoY B.I.» Turkln Tu.T., 1973* Proceed, of 13th Inters. Confer, os Ooaalo. Bays, Dearer, 1^, 509. Galper A.M., lirilloY-UgryuaoY V.G.,.lurochkin A.V., Leikov K.G., LuchkoT B.I., Turkin Tu.T., 1975, Proceed. fell •* of 14w Intern. Confer, on Coaalc Rays, Munich, ^, 9$. Gregory P.O., KrSnberg.P.P., Seaquist S.B., Hughes V.A., Woodsvcrth A., Viner M.R., Betallack D., Hjellalag B.M., Balick B.I., 1972, Hature Pays. Sci., 23_2, 114» CaniEarea C.S., McOlinbock J.B., Clark G.W., Levin W.H.G., Schnopper H.W., Sprott G.F.» 1973, Nature Phys. Sci., 241, 28. Kurochkin A.7., Luchkov B.I., Turkin Tu.T., 1974, Koamichea- kie Isaledovaniya, 12, 486. Leach S.W,, Hurray S.S., Sohreier E.J., Tananbaua H.p., Ulaer M.O., Paraignault D.R., 1975, Ap. J., 133, 184. 1ляЪ B.C.,.Fichtel C.B., Hartnan B.C., Sniffen D.A., 1'hompeon D.J., 1977* Ap. J. Lett., 212, L63. Paraignault D.B., 1977» preprint SAO B"#545. Stepanjan A.A., Foain V.P., Vladialraky B.M., 1971» I*v* XRAO, 45, 42. Stepanyan A.A., Vladimiraky B.H., Heahpor Tu.I., Foain v.P., 1976, Izv KBAO, 55, 157. GAMMA-RAY EMISSIOI SPECTRUM OP CYG X-3 SOURCE AMD ITS POSSIBLE ORIGII. A.A.Stepanian, B.M.Vladimirsky, Yu.I.leshpor, V.P.Pomin Crimean Astrophyslcal Observatory, Crimea, 334413 USSR Four-yeara„of observations of ~f -ray flux fro* Cyg 1-3 uaing the Cerenkor technique are analysed. It ia 1? shown that f -eaiaaion of quanta with the energy >10 ev coneiete of two componenta: a periodic comppnenx with the period T»4.8 houre and a sporadic component. The periodic component hae no significant time variation. It is found to be on the average 1.3'Ю"11 quanta cm~zsec~'. The spe cified value of the period is equal to 0.199684+2*10~bd. The sporadic component has a steeper spectrum that the periodic one. Comparison of the f-emission spectrum of Cyg X-3 with the spectra of other galactic sources was made. Very high energy t" -quanta may be generated either via *3l° -decay processes or by bremsstrahlung of high-energy electrons. Observations of very high energy f-emission from X-ray source Cyg 2-3 were begun at the Crimean Astrophysical Observato ry in September 1972 [1]. Briefly, the installation for detec tion of EAS Cerenkov flashes consists of four parabolic mir rors 1,5 meter in diameter with photomultipliers in their foci. The pulses from two detectors were supplied to a coincidence circuit with a resolving time of 5 nsec, so there are two inde pendent sections. The threshold energy of EAS registered is 12 about 2*10 ev. The results of the observations of Cyg 1-3 obtained from September 1972 to November 1975 are presented here and discussed. Total time of Cyg X-3 observations was about 130 hours. The mean values of the jf-quanta flux expressed in units —11 —2 —1 10 quanta cm sec for each month are presented in Table 1* The flux averaged over four years of observations is —11 —2 —1 equal to (3.0+0.8)*10 ' quanta cm sec . Intensity variati ons with X-ray period 4. 8 [2] have been revealed in the data obtained. 'Analysis of all the observational data allows more IJ6 Table 1. aonth VII VIII IX XI year 1972 +19,4+7.8 -9-8+10.0 +33.2+28.1 73 -9.8+8,6 +15.3+5.6 +5.1+8.1 74 -3.8+3.2 +8.4+2,0 +0.2+1,8 +3.9^.2.4 +0.5+4,9 75 +2.0+2,2 +1,7+2.7 -0.7+3.2 +1.0+3.5 preetse definition of this period. Folding all the data as a function of trial period T results in a very good maximum of the X -test for the values T«0.199684 d (see Fig.1). The pha se diagram for this period is shown in Fig.2. A % is the mean amplitude in the interval of the phases, the zero point correspo nds to JD«2441581.874 [2j. Maximum "f -radiation i« observed for the phases 0.130-0.185. The second maximum is not statis- # ticaly significant.HUOllli., fCi'^giilPeriodiUc XJ.UJfluxL IiDs C4UCIequaXl tIrUo (1.6+0.35)V liOTUiJ^y 10\\i' —2 —1 — Foquantr anothea cm r sepeac k o(phasn thee averagintervael fo0.74r th0 e• phase 0.795s )0.13 the0 flu r x0.185 is . (8.3+2.8)*10~11 quanta cm sec The mean value of Y*-ray flux averaged over the period is —11 _0 _1 1 fi£l.3*10 quanta cm sec . As one may see in Fig.2 the 30 light curve of у -radiation I differs from that of X-ray emission, which has a near Sins го ly sinusoidal pattern. The 0 analysis of monthly mean I I flux values of periodic ra 4. ю • diation does not reveal any ' essential, change over all the time of observations. •tf n S^S^-j/ ^v^i^ _, In all probability these 0.193660 aW690 month to month variations Treat period day are statistical fluctuations (P(>X2) - 0-25). 137 Th» sporadic component of f-radiation may be de tf 4 • •I fined as the «mission out- aid* of above mentioned phaae interrala. The aean 3 oJUrLarrUlf u flux of thie component for TJ all time* of observation 2--. ia « 1.7*10 1 quanta en <ц ач as at ю <• sec" . Sporadic flux chan The phase diagram ge a drastically from on* ob fig. t. servational time interval to another. For example, in August 1974, the flux value was (6.2+2.1)*10~'%-11' quanta cm sec < f- test shows that such changes can not be interpreted as statistical ones because the probability of such deviations occuring by chance is only Р(>Я« 3-1СГ3. A suggestion was made [3] that the Cyg 1-3 •/* -radiation spectrum is steeper than that of cosmic ray charged particles in the same energy range. If this is so, the f -radiation flux may be expected to fall off more rapidly than cosmic ray back ground flux, for increasing zenith angle. This relation was investigated separately for the perio dic component and the sporadic one. Analysis of all the data led to the conclusion that the sporadic ~f -ray flux of Cyg X-3 was observable only when the zenith angle was Q& 20° . If we assume that the dependence is caused solely by the diffe rence between the spectra of the source and cosmic rays, a source spectrum slope may be estimated. It was found that for the integral power low spectrum the value of the power is >3-0. It is very interesting that the amplitude of the perio dic component does not depend on zenith distance essentially, so we may conclude that the spectrum of this component appa rently is the same as cosmic ray spectrum. The spectra of Y* -ray emission of the periodic oomponent of galactic sources FSR 0833, NP0532 and Cyg X-3 are compared is 7ig.3. One can see that spectra of all the sources are simi- . 138 . lar. It may be noted that the duty cycle of radiation pulaee of PSB 0633 and MP 0532 ahortena «fall* the energy of tb* quan ta increaeea. Thie duty cycle la equal to ••* 0.1 for the energy arЮ av and It diminiahee up to * 0.05 for 1012 ev quanta. The pulaa of 10' ат ^-emission of Cyg 1-3 la approximately aa aaall aa of 0.05 of the period alao, but it aay be about 0.2 for the anargiea & 4*10' ev [5}. The analogy in thia feature and aimilarity of the apeotra euggeata the identity of the ae- Fig.3. The apectra of )f - eaiaaion of the periodic compo nent of Cyg X-3, «P 0532 PSH 0833. 0 - Stepanian at al [4] Д - Vladiairaky et al [5] 0 - Kniffen at al [б] o - McBreen at al [7J • - Grindlay at al [в] •>. Thompson et al [9] 4- Bennett et al [10] Ю7 f *' «? «*' «? — toi»dl.y at al [11] energy, e/ chaniama of ]f-emission of all theae objects. The preaence of pulaed radiation teatifiea that particle acceleration takes place near the objects. It is possible Qunn-Ostriker'a mecha nism [12] operated. The broadening of radiation pulaea for jf- quanta with the energy st 10 ev may be explained if we assume that these quanta are generated via 3f -decay. Angular half- width of J*-quanta generation with this energy ia about 30°, i.e. the duty cycle ie equal 0.1+0.2, in accordance with the measurements. Another approach to understanding the |f -ray generation proceae is a speculation that the spectrum of the electrons emitting radio synchrotron radiation continuoa up to the ener- gy 10Ie ev. |f -radiation of Cyg X-3 may be explained for thie case as bremsstrahlung of such electrons in the vicinity of a IS9 young pulsar la • binary system. A plasma number density of 1012 om~3 is netted at a dlatanoe at5*1011 om from the aulsar to provide the flux measured. The Boat important oonolualona are: 1. Vary high-energy aaiaaion of Cyg 1-3 oonalata of two oea- ponentet periodic and aporadlo. 2. The beat determined period of if-radiation of Cyg 1-3 is 0.199684+2»10~6 d. 3. The apectrum of f -quanta in the Cyg 1-3 aporadio oompo- 12 nent in 10 ат range ia steeper that of oommio raya, ao the power of the apectrum may be aa large aa 3.0. 4. The apectrum and light curre for ]f-emission of Cyg X-3 are aiailar to thoae of two other galactic f -ray sources> pulsars IT 0532 and PSR 0833-45. References 1. Stepanian A.A., Vladimireky B.M., Parlor I.V., Pomin V.P., 1971, lBT.Kryms.Astrofis.Obs., £3., 42. 2. Canizares C.H., McClintock J.B., Clark O.W., Lerin W.H.Q., Schnopper H.W., Sprott G.P., 1973* lature Phys.Sci., 241. 28. 3. Vladimireky B.M., Stepanian A.A., Pomin V.P., 1974» 1ST. Krymek.Astrofis.Obs., SI, 3. 4. Stepanian A.A., Vladimireky B.M., Xashpor Tu.I., Pomin V.P., 1976, IsT.Krims.Astrofis.Obs., £5., 157. 5. Vladimireky B.M., Oalper A.M., KiriloT-UgrumoT V.O., Kuroeh- kin A.V., Luehkor B.Z., Beshpor Yu.I., Stepanian A.A., Pomin V.P., Turkin Yu.T., 1975, Pis'ma т Aetron.J., J,, 25. 6. Kniffen D.A., Hartman R.C., Thompson D.J., Bighami O.P., Pichtel C.B., 1974, Mature, S51, 397. 7. McBreen B., Ball S.E.Jr., Campbell M., Oreisen K., Keen D., 1973, Ap.J., J8±, 571. 8. Grindlay J.E., Helmken H.P., Weekes T.C., 1976, Ap.J., 209. 592. 140 . 9. Thoopaoa fi.J., Plcbtal C.I., Knlffaa D.A., 6(«laaa M.B., 1975, Ap.J.latt £2S> L75« 10. Baaxatt K., Bifnaal 0.»., Boalla 0., Buoeharl 1., Gorl- ••• H., Hamaan I., Kanbaoh 0., Llchtl CO., Mayar- Haaaataandar H.A.., Paul J.A., Soaral I.., Saaaanburg B.I., Taylor B.C., till* K.O., 1976, Aatroa. ana Aatropbyaiea, 52, 157. 11. Qrindlay J.B., Halakan B.P., Broaa Huabury 1., Dayla J., Allan L.8., 1975, Ap.J., £21, 82. 12. Ошш J.I., Oatrlkar J.P., 1971, Ap.J., J^, 523. 141 TITLE: GAMMA RAYS FHON CYC X-3 AND CYC X-l AUTHORi Krishna M.V. Apperao* AFFILIATION< Center tor Space Research Massachusetts Institute of Technology Cambridge, Massachusetts 02139 ABSTRACT Theoretical Q Experimental Q •<>«» Q The X-ray objects Cyg X-3 and Cyg X-l emit radio radiation. The scattering of х-radiation from trie radio electrons results in данта radiation. The gamma ray flux from these objects are calculated using such a model. The observed gamma radiation from Cyg X-3 may be explained by the above model. The calcula ted flux for Cyg X-l im just under the limit set by observations Correlated radio and gamma ray observations are suggested for these objects. In the case of Cyg X-3, the X-ray object is surrounded by a thin gas shell, as indicated by infrared observations as well as burst radio flux decay. The passage of radio electrons through the shell results in a bremsstrahlung gamma ray 'burst*. The calculated flux will be presented. TOn Sabbatical leave from Tata Institute of Fundamental Research, Bombay, India Coordinates: OG 1.2 (Gamma Ray Sources) Majlintaddress: pi0fessor Krishna M.V. Apparao Center for Space Research M.I.T. 37-535 Cambridge, Massachusetts 02139 ПИ MOUNT NOPXIMS «KY ЪОНУЮГ ОР SOUJCU 0Г Ю11 to 10l2*V GAMMA «AYi T.C. Ники, И.Г. Belsfcen. j.l. Grlndlay, ». tor in* Theoretical Q E«ylHrtil Q »*• Q The 10m reflector on Mount Hopkins in southern Ariiona has been used to scan the celestial sphere front declination -10* to +70* for point sources of 10^ to 10 eV gamma rays. Ten independent channels each with a one degree beast were used at two thresholds. In addition, 4 1.5m reflectors were used with somewhat larger fields of view. The results of this two year survey will be reported. Particular emphasis has been devoted to the region* where point sources of 100 MeV gassu rays have been reported by the CASAVANE collabration; results of these observations will also be reported. Coordinate!.- И-*"»-*"—1 Dr. T.C. Weeke. Mt. Hopkins Observatory Box 97 AMADO, ARIZONA 85640 U.S.A. 143 COSWtC Y RAY OISEKVATpSS П» THE 0. I - 4 H»V ЕЗТСТИГ КМКХ Г. MANDROU, M. НШ-, A. ГЛГГШТ. C. VEDREKKE Centre d'Etude Spatial* dee Rayonneaenta - Toulouse Theoretical Q Experimental {x\ Both Q Energy spectra of Che Crab Nebula and Cygnua X-l in the 0.1-4 MeV energy range, obtained with a balloon borne у ray telescope, are preaented . In addition, prе1ininary reaults obtained for the galactic center region during a February 1977 Brazilian balloon campaign will be discussed. Coordinates: OG-1.2 Mailing address: CENTRE D'ETUDE SPATIALE DES RAYONNEMENTS 9, avenue Colonel Roche 31400 TOULOUSE ч France 144 DIFFUSIVE SOFT X-RAY BACKGROUND FROM THE INTERACTION OF MAGNETOSPHERIC ELECTRONS WITH THE UPPER ATMOSPHERE J. P.. Inl Theoretical Q Experimental Q Bo«« 0 X-rayc produced by magnetoapharic electron» that interact with the resid ual atmosphere above typical rocket altitudes contribute a aignlflcant diffuse background to soft x-ray astronomy experiments launched from low and mid- latitude sites Including White Sands. New Mexico and Kauai, Hawaii. This at mospherlc emission ia not necessarily accompanied by an electron Пих at the rocket because the electrona that generate the x-rays mirror above the rocket at geomagnetic "quiet times. An approximate calculation of this x-ray back ground based on extensive satellite measurements of low-altitude electrons suggests that the reported diffuse galactic x-ray flux below a few keV ia actually largely atmospheric in origin. This calculation emphasizes the need for measurements of soft cosmic x-ray component from a substantially higher altitude (~500 km) in order to escape atmospheric contamination. Coordinates: ©G 1.3 (Х-Ray Astronomy) 'Mailing tddress: J. G. Luhmann Space Sciences Laboratory A6/2447 The Aerospace Corporation P.O. Box 92957 Los Angeles, CA 90009 f у/ ARIEL S NtASUftENENTS OF SOUftCCS WITH MVEA LAW SPECTM AND FAST TINE VARIATIONS H. J. Co*. A. K. Engel and J. J. Quenby •lackatt Laboratory, Imperial CoIIeg», London SW7 212 Several hard X-ray source» appear to show both power law spectra and exhibit fast, aperiodic tloe variations. One example of this class, Ser X-l, Is studied In detail. It has a mean spectral shape тг - 2-9-Ю"2-3 ph cm*2 ke\rl sM, shows no periodic variations above the St level at low energies In the 6 1. Introduction Fast fluctuations in the intensity of X-ray sources are indicative of compact objects and, in particular, attention has been drawn to the associa tion of the millisecond temporal variations of Cyg X-1 and the periodic motion of material in stable orbits close to the supposed black hole. Cyg X-l may be a source by virtue of the repeated action of the Inverse Compton effect, resulting In a power law photon spectrum (e.g. Carpenter et a). 1976, Shapiro 'et al. 1976). In this paper we investigate whether or not there Is a class of objects which combine the characteristics of fast, non-periodic time variations and power law spectra. They would contrast with binary pulsar sources where the fast time variations »re dominated by magnetospherlc rotation effects and where the spectra seem to be modifications of thermal emission with typical temperatures, kT, of up to~13 keV (e.g. Her X-l, Cen X-3). 2. Observations Forman et al. (1976) have studied 17 strong galactic sources, many of which are known to be binaries, but which do not lie in the class of estab lished binary pulsars. A Table Is provided ranking the percentage power in the variable component, measured in 0.096s sampling bins and for between 500 and 3000 samples. The Forman et al. results are reproduced below, together with the exponent of the power law spectral fit obtained by Ariel 5 where possible (e.g. Carpenter et al. 1976) and In some other cases using the_ Guo and Webber (1975) balloon data, although these last authors cannot, in general, differentiate between a power law and an exponential form. Excluding the. Crab, up to 9 of the 16 sources In the Table may have power law spectra and there is little doubt about this association In the cases of the three sources at the top of the list. Low hard X-ray intensities pre clude definite statements abbut the other seven sources. In the following, we will confine our attentions to a detailed study of one member of the suggested class of objects, namely Ser X-l. 146 TMLE Spectral exponent • In dTaE Percent power In Source 0.096t lln Ariel 5 Cuo and 1Мйг 3U1956*35 (Cyg X-l) 18.4 i 0.5 2.2 t 0.3 3U1837+04 (Ser X-l) 13.1 ± 1.8 2.3 t 0.2 3U1516-56 (Clr X-1) 12.5 ± 0.6 2.1 ± 0.3 3U1744-26 (CX3+D 11.6 ± 1.1 ) 3.0 ± 0.5 ) 3U1735-28 (CC tran, 11.5 ± 1.3 ) compos 1te GCX7) ) 3U1636-53 (Nor X-1?) 11.2 ± '.7 [3.5 ± 0.2 7 [3.0 ± 0.3 7 3U1820-30 (NGC6624) 11.0 ±_ 1.3 3U1811-17 (GX13+1) .10.8 ± 1.2 3U1758-20 (GX9+1) 10.7 ± 0.8 Too weak to see 3U1702-36 (GX349+2) 10.0 ± 0.9 3.4 ± 0.4 3U2030+40 {Cyg X-3) 9.2 ± 1.1 1.6 ± 0.3 3U1705-44 9.1 ± 1.3 3U1642-45 (Ага Х-1) 8.8 ± 1.1 3.5 ± 0.2 3U2142+38 (Cyg X-2) 8.7 ± 1.0 Too weak to see 3U0531+21 (Crab) 8.3 ± 0.9 2.25 ± 0.05 3U1758-25 (GX5-D 6.7 ± 0.8 Confused with GCX 3U1813-14 (GX17+2) 6.1 ± 0.5 Joint study of Ser X-1 (Coe et al. 1977) with the Ariel 1 scintillation' telescope, col 11 ma ted proportional counter and rot fit I on modulation collimator has established a number of characteristics. The spectrum in 1976, days 284- HM > 289, may be represented by a power law -JF ~ 2.9 E"2-3 ph cm"2 keV"1 s"1 wjth a * low energy cutoff arising from a hydrogen column density of (1.1 x 0.4)x 1022 atoms cm"2 (Figure 1). No periodic variation of amplitude greater than 5* of ^ the steady flux has been seen at low energies (2-8 keV) within the period - range 64 - 24000s. In 1975 and 1976, Intensity fluctuations tend to lie \ 147 - . ,- - I - »— - within « factor 2 of th# 1 лмап value, but the frac tional power spectrum In \ SER *-' the aperiodic varlatlont follows roughly a ml nut one power law in I0"s - Ю - 10° s"1 range (Figure 2). SAS III and UHURU obser у vations were used for the highest frequency point (Li et al. 1977. Forman et al. 1976). LI et «I. have also detected up to 1<П- 22 X-ray bursts from Ser X-l, the largest of which и reached flux levels about I 10 times the normal intensity with a 3s rise ОС " \/"^ time. J Д/ 3. Discussion of Е em!ssI on mechanisms gо for Ser i-\ о Using the Seward distance Indicator with I СРС \ 4 1976 \ the observed low energy DN 2ве-6-2£3-6 V cutoff and taking the spectrum down to 0.3 keV ю - . ST —^ yields a luminosity TON 26Д-6-239-6 estimate of L "6.7 x 1037 erg s"1 for Ser X-l. If it is a neutron star, black body radiation from an object of this size falls to explain the 10,- 5 i I i.i 1 . emission because the 10" 10 10' 10 required temperature, kT, is only 1.38 keV. ENERGY (ktV) Thermal bremsstrafelung Fig. 1 Col I(mated proportional counter and (TB) with kT = 20 keV may scintillator telescope data obtained by Ariel. 5 fit the higher energy spectral points and hence should be considered. From an accretion disk, thickness 2h and radius'r2, the TB luminosity for electron density n Is 27 2 2 2 L - 1Л25 x n" T*n 2uhr2 erg s"4Equation )) with R typically M0" r2. Gravitational energy as source of emission means L * —— for accretion rate A, 2* r2 2h • compact object radius rj, while fluid continuity implies птруг — H (Equation 2) for^dfffusion velocity vr. Heat loss considerations limit the and value of r2 to < 100ri using this.value, parameters for a neutron star and (1) and (2) yields an estimate vr ъ 6 x 107 cm s"l and n <\. 1019 cm"3 on the TB model. v. is suitably less than the Kepplerlan velocity in the disk and the value of n yields a skin depth at 20 keV much greater than r2, thus confirming the assumption of optical thinness. 14* Invert* Complon cooling According to Pc - Vj oT cpn, where p it photon energy den»Ity, r % ) end oT it th« Thomson crost- section, It con- pctltive with the TB effect. Pho- tont ere attuned to dlffute through the tldet of the ditk with flux F-DgradP*^^ 2 and L — 2i4-2 F of Compton cooling dominates. Hence we derive 2 2oT'n hl ^ and irr2z find the ratio Compton cooling TB cooling -k.Z x 1011-Ь- r22 (Equation 3). For 8 r2 - Ю cm - lOOh, FREQUENCY!S1 this ratio is 42:1*. This demonstration Fig. 2 Ariel 5 rotation modulation collimator data of the dominance of (vertical bars) and SAS 1,11 plus UHURU data (horizon the inverse Compton tal line) for fractional power density. effect is supported by the power law nature of the spectrum. Several workers suggest that such a shape arises from the repeated collisions of Initially soft, sub kaV photons with a near relativistlc electron gas (Shapiro et al. 1976. Carpenter et al. 1976). Such a gas is predicted in the two temperature model of Eardley et al.(1975) where a cool, geometrically thin outer disk feeds a hot, optically thin but geometrically thick Inner region with electron temperature -И09 °K. An instability Inherent in the mechanism and located at the boundary between the two regions where gas and radiation pressures equalise may explain the observed temporal fluctuations. Suppose that diffusion resulting from hydrodynaraic turbulence Is tbe mode of matter transfer across the inner region and that increased X-ray luminosity corresponds to an upward fluctuation In diffusion rate. In turn, this Increased diffusion can be due to the appearance of a favourable scale size and eddy velocity in the spectrum of turbulence produced by the instability. The lifetime of this particular scale of eddy is supposed to depend on the life against viscous decay. 149 C*pr**»in9 these principles In equation» yield*: L ~ 7— • (Equation *), «Л»г» d 1% inner region width (3.2 * 107 ся taking the Shapiro •( «I. 1J/74 para meters for Cyg K-l, • source of roughly similar lunlnos!ty). 4t — diffusion t in*, c — constant. R ~ VA* (Equation 5) where R •*• KQ ~ Reynold* nutnber for onset of turbulence. v > is eddy scale size, v^ eddy turbulent velocity, u —coefficient of viscosity and m the mass density. 2 2 d ~ 2nic (Equation 6), n —number of collisions eddy make», *c— collision length determined by average properties of turbulence. n - — (Equation 7) t< "~ col I Is Ion time for eel I *. vjt» - Xc (Equation 8) F(k) — Kk" I1 (Equation 9) The turbulence is assumed to have a tolmogoroff spectrum of eddies, wave number k — £* . дТ — 6Av,*3 (Equation 10), &T — time of intensity fluctuation (decay time) for cell size i, conventional hydrodynamic theory for decay time of eel) Л being assumed. Equations (It), (5), (6). (7), (8) yield L - l^Sii if. (Equation 11) dm X Defining a frequency spectrum G(v) of intensity fluctuations where u — (дТ)~1 by F(k)dk -G(v)dv yields from Equations (5). (9). (10) G(v) а \Г /6 (Equation 12) The observed power spectrum of fluctuations is LJG(v) and Equations (11) and 2 (12) give L2G(v) a\ z v" 1г (Equation 13) J.r f This predicted v /3)aw is in reasonable agreement with the high frequency points in Figure 2. X-ray bursts from Ser X-1 as observed by Li et al.(1977) can be understood as extreme fluctuations in the above model. A factor 10 increase in X-ray intensity requires a factor 10 decrease in At in Equation CO. It is sufficient to suppose that Xc becomes 3XC due to some change in the overall character of the turbulence to achieve this required At change. Then, using At — 3 sec,- the required rise time, d — 3-2 x 107 cm and estimating 5 Xc ъ 2h ъ 8 x 10 cm In Equations (6), (7) and (8) provides an estimate Уд i" 2 x 108 cm s"1, still suitably le'.s than the near relativtstlc Kepplerian velocity of the inner region. Hence it is plausible to believe that X-ray bursts from an Inverse Compton emitter are just a further manifestation of the instability which drives the source. References Ccte, M.J., Engel, A.R., Quenby, J.J., Carpenter, G.F., Be)1-Burnel1, S.J. and Davison, P.J.N.,1977, Submitted for publication Carpenter, G.F., Coe, M.J., Engel, A.R. and Quenby, J.J., 1976, Proc.R. Soc.J-ond.A., 350. 52 Eardley, O.M., Ltghtman, A.P. and Shapiro, S.L., 1975, Ap.J.Lett. 199, Ц53 Forman, W., Jones, C. and Tananbaum, H., 1976, Ap.J., 208, 849 Guo, 0.0.S. and Webber, W.R., 1975» *roc.1l>th Int.C.R.Conf. Munich, 1, 162 ISO If. Г.К., l«*.ln. W.H.G.. Clerk. G.W., Ctoty. J., «off»*". J.A. *n4 **pp*pori. S.A.. \»77. "IT Preprint CSH-»-77-OH Shapiro. S.L.. Llghiman, A.I», and f»rdl*y. D.H., IJ/4. Ap.J . . m. 1$; 1*J ARIEL V HE*$UHEHENTS Of SOURCES WIT* *OVER LAW WC1AA WD ГА5Т Т1ИС VARIATIONS N.J. Co*. 'A.R. Engel and J.J. Quenby Inpvrl*l College, London Theoretic»! • Experiment»! Q °<>th (jj Previously Ariel V results have suggested that the Black Hole candidates Cyg X-l and A0620-00 «xhlblt а common power law type spectrum, shine by an Inverse Compton mechanism and one, at least, exhibits fast time variations Indicative of a compact source with a scale of a few SchwarzschiId radii. Several other galactic sources are known to show fluctuations on a time scale <0.1 second and one at least, SER X-l, Is now shown by us to have a power law spectrum. Further studies of a number of other candidates for this class of source are being carried out and the results will be reported. Coordinates: OG 1.3 (X-ray Astronomy) Mailing address: Dr. J. J. Quenby The Blackett Laboratory Imperial College London SW7 2BZ, UK *C-JJ 152 ARIEL- 5 OiSCRVATIOHS OF Е XT RA&ALACT IC Х-RAT SOURCE* AND IMPLICATIONS OF HARD X-RAY GALACTIC SOURCt MEASUREMENTS FOR THE GALACTIC DISC y-RAV EMISSION C. S. Oyer, N. J. Co». A. R. Engal and J. J. Quenbv tlackett Laboratory, Imperial Collet*, London SW7 217 The hard X-ray telescope carried on the Ariel S satellite hat detected power law X-ray spectra at enerfle* greater than 26 keV from the ext regal act ic objects ten taunt* A and NGC *i 15 •. and from eight known sources In the galactic plane. When compared with other data, these results llluetrate the varia bility of Cen-A at all wavelengths, but euffast that herd X-ray enission from NGC *J 151 may be constant. The existence of a large subset of power law galactic sources (which are determined to at least 200 keV In the data) suggests a possibly significant contribution to the galactic plana y-ray emission. 1. Introduction The hard X-ray telescope (experiment 'F') carried on Ariel 5 has operated successfully since hunch in October 1971*. The telescope comprises an 8 ся2 x k cm central detector of Csl (Na) covering the range 26 keV to 1.2 HeV. Active col Iimation is employed to give an opening angle of 8° and the telescope is off-set by 3° to the satellite spin axis in order to separate sources from background modulation. Two off-sets are used on each source in order to remove background modulation. So far about twenty sources have been identified by this instrument. in this paper we concentrate on those which have been measured up to at least 200 keV and are best described by power law spectra. 2. Extragalactic sources (a) NGC **151 This Seyfert galaxy was identified in the UHURU data (3111207+39) and was found to have changed in intensity at low X-ray energies from data taken in November 197^ with, the coll inflated proportional counter, experiment C, on Ariel 5 (Ives et al., 1975). Experiment F viewed this source from 3rd to 7th January 1976 and again from 17th to 23rd December 1976. The totality of data from these two observations gives the three significant points plotted ir Figure 1. Experiment C also viewed the source at these times and for January 1976 obtained a spectrum identical to that of November \ST*. The December 1976 experiment C data also agrees at high energies (Barr, private communication). The experiment F points fall on a power law fit, E"1'6, to the low energy data and confirm the power law nature of the source up to 200 keV. Also shown are hard X-ray data from a balloon flight by the UCSD group (Paciesas et al. 1977) during June 1975- Their data also agreed with 0S0-7 observations taken in October 1972 and April'to May 1973 over the range 10 to 100 keV. Thus it would seem that over a four-year period the hard X-ray emission has remained very constant, suggesting that intensity variations at low energies may be due to variable absorption of a constant intrinsic source. (b) Centaurus A (NGC 5128) This source, among the brightest and nearest of radio sources, has been observed In nearly every photon energy domain, although unfortunately simultaneous observations have not occurred. The X-ray emission 153 woo Id teen to art** In the nucleus of the galaxy and I» found to vary on tin* scales of 1 1 nun—1 1 1 мин—1—1 11 мм months and pottlbly days (e.g. Grlndlay Л«\ «—•r Z •t al. 1975). - + + Jul On Iff» Variability ha* alto been observed at m and m» wave- a.c.1. langtht. + f Experiment F viewed this source from 23rd to 26th Decem ber 1975 although unfortunately an off-set far from optimum was em ployed and only limited sensitivity nr* - was available. Results of this measurement are presented in Figure ! 2 together with a-. results from experi ment C obtained in January 1975 (Stark et al. 1976). Also shown are results from Lanpton et al. (1972) obtained in 1969-70 and from Hall et al. (1976) obtained In April 1974. These data are shown as fitted by the authors and J—1 1 1 nm 1—1 1 1 inn - 1 1 error bars are •• 1*0 indicated. EMiir k«V Although the sta tistics are limited, the hard X-ray results show a considerable enhancement over the 1969*70 measure ment and are not inconsistent with the April 197*1 data. Stark et al. note that Copernicus observations showed a k0% reduction in the 2-6 keV flux by July 1975 and the hard X-ray observations are not Inconsistent with a reduction of this order. (c) Discussion Although our results confirm the variability of the hard X-ray flux from Cen A on the timescale of several years, the flux from NGC 4151 would seem to be much more constant. . The model most commonly Invoked to ex plain hard X-ray emission from such compact extragalactic objects Is the Compton scattering of radio and Infra-red photons by the same energetic electron popu lation which produces the low energy photons by magnetic bremsstrahlung on the J et-J? IS4 «i*lent aaqnetlc fields. Such • n«.r> a nod* I has •*•» app lied to NCC 4151 by Nushotiky (1977) and to C*n A by Grlndley (1975)- In tn» latter cat* a two-component nod* I would seen to b* demanded by the radio and high energy Y-ray (>3 x 10u *V) data. In both cases, the hard X-ray flux results from Coup- ton scattering of 15 HeV to 1.5 GeV electrons on a syn chrotron spectrum produced on a ^2 i gauss field and ex tending from cm or mm to micron wave Е lengths. It Is interesting to note that the sizes of the emitting regions demanded by the radio synchrotron self-absorption are 1.5 Iight days for NGC 4151 and 11.5 1ight days for Cen A. These values .nay be compared with the 7-day llfe- timeagainst synch J—I. I I I llll » IXI i I ill rotron energy loss for the electrons responsible fori 10 Eiwrj/ luV keV photons. Thus the electrons res ponsible for hard X-ray emission have lifetimes which are longer than the source transit time for NGC 4151, but shorter for Cen A. This could Imply that a large number of par ticle sources are required ю fill the emitting region of Cen A, whereas a single constant source could suffice for NGC 4151, thus explaining the contrast in variability. However if a magnetic field of 1 gauss is assumed for Cen A, as favoured by Price and Stull (1975) from circular polarization measurements at 10.7 GHz, the lifetime of electrons is increased at least fourfold and this argu ment no longer holds. It is also interesting to note that Price and Stull observed a factor 1.8 increase in the radio flux between a period from June to August 1973 and later observations In Hay 1974 and January 1976, while X-ray observations made close to these times showed no such variation (<1.3). OG-37 » IS5 Clearly, ettael Ishment of the validity of the «oo»l for Cen A «nil ecler-Mint- netion of the source parameter» require* timulteneou» observation» ecrott the radio. Infra-red and X-ray bands. It U hoped to perform such observation* during Jon» Sth to Hth, 1)77. when the Ariel $ experiment» 'C and 'Г< «III view the source. Use of optimum off-sets will greatly I «prove the sensitivity over the result* presented here. 3- ttalectlc plene sources (•) Results From observations of known galactic plane source», experiment F has measured hard X-ray spectra, which are best fit by power laws, fro* a sub* set which totals eight to date. These observations, for which significant data points exist up to at least 150 keV, are presented In Table I, and Include cases for which a smgle power law fit suffices (e.g. The Crab) and ones for which a thermal component plus hard X-ray tall Is required (e.g. Cyg X-3). The spectral exponents (photon number) relevant to the hard X-ray data are given, and the fits »r* extrapolated to give the Integral fluxes of демм-rays of energy greater than 100 MeV which would be observed If the hard X-ray spectral form extended into the gamma-ray region. Also given are gamma-ray observations and upper limits ob tained with the SAS-2 experiment (e.g. Hartman et al. 1976). Errors in the spectral fits themselves lead to typical errors of about a factor 60 in the extrapolated gamma-ray fluxes. However it can be seen that the fluxes obtained are in the realm of current frsy sensitivity (*»10~* cm*2 s"1) and in many cases well above, so that the spectra cannot maintain their form ail the way to 100 MeV. Only in the case of the Crab, for which experiment F data extends to 1 HeV, does the hard X-ray data power law fit smoothly onto the gamma- ray spectrum. It is possible, however, that a large number of X-ray sources have hard X-ray and gamma-ray tails arising from Inverse Compton scattering (see below), which can contribute to the galactic disc y-ray emission. One source of strength M0'6 cm'2 s"1 in each 5° of galactic longitude would supply about half of the galactic disc emission outside of the central regions (-30" < I11 < 30°). Table 1 Power law X-ray sources Galactic Spectral Extrapolated Source Name Coordinates Exponent y-ray flux 1 7-ray observation Integral flux cm"2*"- >100 HeV GCX + 0.0 ) (GX3+D 3.+1 ) 3.0,± 0.5 3.7 x 10-9 (GX1+4) !,+-> ) Ser X-1 36,+5 2.3 ± 0.2 7.3 x 10-7 Cyg X-1 71 ,+3 2.2 ± 0.3 1.7 x 10-5 < 2.7 x 10-6 Cyg X-3 80,+1 1.6 ± 0.3 3.5 x 10-1» 4.4 x 10"6 Crab 185,-6 2.25 ± 0.5 4.7 x 10-6 3.7 x »0"6 3U1223-62 301,-2 2.0 ± 0.3 2.2 x 10"6 3U1254-69 303.-6 1.6 ± 0.1 8.7 x 10"5 Clr X-f 322,0 2.1 ±0.3 2.5 x 10"5 156 (t>) Discussion of tourer mechanltmt Sonc theoretical ju» 11f1 tji ion •» required for ertrapolal Ing »ucS * wide variety of objecti to Y-r«y enerijiet. In the case of the Crab, synchrotron radiation of a power law electron tpectru» Is thought to produce both the X- and gamma radiation. For the other caiet. X-rays are thought to arise from accretion of material onto a collapsed object from a massive early-type companion. In such cases, about 100 hfeV per proton Is readily available in gravitational energy but the spectral form of the emitted radiation depends on the details of the slowing down mechanism Involved. Host of the radiation would seem to arise thermally with temperatures of I to 10 keV but there exists some efficiency for producing hard radiation tails up to lP'j MeV by Inverse Compton scattering of an energetic electron distribution on the soft X-ray photons (see for example, Shapiro and Salpeter, 1975)- An efficiency for gamma-ray production cf 1% of the hard X-ray fluxes would be sufficient to make significant y-ray sources out of the objects presented here. An alternative method for production of an energetic (100 MeV) electron distribution occurs if the collapsed object is a rapidly rotating neutron star (pulsar?) as suggested for Cyg X-3 by Fabian et a). (1977). In this case, the hard X-ray tail arises from inverse Compton scattering of the stellar ultra violet photons, and the 100 MeV gamma rays from Inverse Compton scattering of the soft X-rays. As noted by Fabian et al., produrtion of the gamma rays nearer to the compact object than 1010 cm would be difficult because of pair production by interaction with X-rays. The young pulsar model of Lamb et al. (1977) might allow the relativistic electrons to get out further from the object, but then some evidence of ъ 30 msec periodicities should be evident in the sources con sidered. We also note that Cir X-l is suggested by Fabian et al. as being another example of the Cyg X-3 type of source. Clearly all these objects are worthy of further study across the entire range from 50 keV to 100 MeV in order to elucidate these source mechanisms and clarify the possible contribute . of discrete sources to the galactic disc gamma- ray emission. k. References Fabian, A.C., Blandford, R.O. and Hati-^ett, S.P., 1977, Nature 266, 512 Grindlay, J.E., 1975, Astrophys.J., 229, Ь9 Grindlay, J.E., Schnopper, H., Schreier, E.J., Gursky, H. and Parsignault, O.R., 1975, Astrophys.Ji, 2p_1_, LI33 Hall, R.D., Heegan, C.A., Walraven, G.D., Djuth, F.T. and Haymes, R.C., 1976, Astrophys.J., 2_l£, 631 Hartman, R.C., Fichtel, C.E., Kniffen, D.A., Lamb, R.C., Thompson, O.J., Blgnanl, G.F., Ogelman, H., Ozel , M. and Turner, T., 1976, NASA X-662-76- 154, p12 Ives, J.C., Sanford, P.W. and Periston, M.V., 1976, Astophys.J., 207, L159 Lamb, R.C., Fichtel, C.E. Hartman, R.C., Kniffen, D.A. and Thompson, D.J., 1977, Astrophys.J., 2J1> L63 Lampton, M., Margon, B., Bowyer, S., Mahoney, W. and Anderson, K., 1972, Astrophys.J., JjM., H5 Mushotzky, R.F., 1977, Nature, 265, 225 Paciesas, W.S., Mushotzky, R.F. and Pel ling, R.M., 1977, M.N.R.A.S., 178, 23p Price, K.M. and Stull, M.A., 1975, Nature, 255, 467 Shapiro, S.L. and Salpeter, E.E., 1975, Astrophys.J., 298, 671 Stark, J.P. Davison, P.J.N, and Culhane, J.L., 1976, M.N.R.A.S., _^7jt, 35p 157 RECENT ARIEL V MEASUREMENTS ON HARD X-RAY »URSTS AKD THE IMPLICATIONS FOR т-IURST ORIGIN J.J. Qu»nby. N.J. Co» and A.R. Engal Imparl*l Col lag». London Theoretic»! • Experimental Q] Both Q Further Ariel V observations of X-ray bursts In the Galactic Centre region era scheduled In Spring 1977. The anticipated results on the relation between hard and soft X-ray measurements of the same events will probably show similar time of arrival differences to those seen by Ariel V In Spring 1976. The implications of this and the relationship of these events to y-ray bursts will be discussed. Coordinates: OG 1.4 (Gamma-ftay Bursts) Mailing address: Dr. J. J. Quenby Blackett Laboratory Imperial College London SW7 2BZ UK IJI OBSM.VATIO* OF ТИ HARD HUDIATIC* FROM ТШ( ОШЛГ I.L.Aptekar*, S.Y.Qolenetekli, YU.A.Our'yan, T.I.Il^Mkll, 11 I.P.maaet*, T.I.Faaov, X.A.Sokolov A.F.Ioff* Fhyeioal-Teohnloal Institute, Academy of Solenaee, USSR Theoretical II Experimental At th* study of the degree of ieotropicity of the diffuse backgro und lm the herd X- end soft gem—-rays on Kossos 461 satellite, a weak anisotropic oomponent of the radiation was revealed, .coning predominantly fron the galactic equator region with 315°>1 *45°. She spectrua of this radiation traced up to energy of 600 ke? differs sharply fron the diffuse background spectrum. The flux of this radiation in the 30 - 600 keV range varied from the value of ~ 8x10"^ em"2*"1 to 2*10"' om"2s~1, what was only small part, about 2-8Jt, of th* full flux** of th* diffuse radiation in the corresponding spectral range*. Further investigation* are specify ing th* radiation intensity distributions along the galactic lat- titudes and longitude*. Coordinates: 00 1.1. (Diffuse Comic and Oalaotio Gamma-Ray*) Hailing addresst A.F.Ioffe Physical-Technical Institute, Academy of Sciences of the USSR, 194021 Leningrad, USSR 159 HAW) X RAT OaSCRVATIOMS OP №1820-») by Mvin MUtLiy Centra d'Btuda Spatial* daa llayonneaanta - Toulouaa - Franca Thaontiul Q ЕхраНииНа! g) **k Q Tha X Ray buratar 3U1820-30 waa obe*rv*d tvica with balloon borne hard X R*;y detectors, one* on 1 March 1973, and once on 30 Nov** ber 1976. The high energy X ray spectrum la given, and the evidence for variability la diacuaaed. Coordinates: 0G1.3 Mailing address: Kevin HURLEY C.E.S.R. 9,. avenue Colonel Roche 31400 TOULOUSE 160 Observations of Co«mic Х-Ray Souroi Above 20 keV from 0S0-8. B. R. DENNIS, C. J. CRANMELI,. J. F. DOLAN*. K. J. FROST, L. E. ORWIG, NASA-Goddard Space Flight Center, J. H. BEALL, U. Md., G. F. Maurer, Cath. U. —The high-energy X-ray detector on OSO-8 has observed nost of the known X-ray sources detectable above 20 keV. The energy spectrum of the Crab Nebula has been determined with a precision not previously possible up to 500 keV. In March 1976 the spectrum of the total emission is well fitted by a power law with an index of 2.00 + 0.08. The pulsed fraction of NF0532 was also measured over this energy range with a time resolution of 0.3 ms. Cyg XR-1 was observed in Nov. 1975 and again in Nov. 1976. The intensity above 20 keV in 1975 varied by 40% during the 8 days of the observation in snticorrelation with the in tensity below 10 keV, clearly showing the pivoting effect of the power-law spectrum. The spectrum of Cen A measured over 8 days in July/August 1975 and again for 11 days one year later shows no significant variation in the 20-200 keV range, contrary to the decrease in intensity observed over the same period at lower energies. The eclipsing binary source, Vela XR-1 (3U0900-40) is observed to be eclipsed above 20 keV in phase with the observation at lower energies, ruling out certain models of this system. *NAS/NRC Resident Research Associate Submitted by Dr. Brian R. Dennis Code 682 NASA-Goddard Space Flight Center Greenbelt, Maryland 20771 161 TIME VARIATION OF HARD X-RAY FROM CYG X-l M.ttakagawa, H.Sakurai, and M.Uchid* Department of Physics, Osaka City Univereity, Osaka, Japan Flaring activity from Cyg X-l with a timing resolution of 1 msec have been obtained during a balloon flight in 1972. For the Flare, millisecond variation of the intensity is studied with a particular method through the medium of Poisson statistics. The hard X-ray at Flare is concentrated about 95 msec periods intermittently. K Introduction. Cyg X-l is one of the most extensively studied X-ray sources. We shall be concerned here with varia tion of short time scale and long time scale in the intensity of this source in the energy range 30-100 keV through a discussion of preliminary results of a balloon experiment intended to the observation of the emission features of this object. The long time scale variation of hard X-ray, especially a few minutes variation called Flare event, was detected by four groups (Agrawal et al., 1972, Fuligni et al., 1973, Jain et al., 1973, and Nakagawa et al., 1973). The other hand, several "bursts" of millisecond duration in Cyg X-l have been reported from the soft X-ray observation by rocket borne experiments (AS&E group, GSFC group, NRL group, and MIT group). So, it is attractive and important to study the time variation, particular millisecond variation, of the inten sity of the hard X-ray from Cyg X-l (Nakagawa et al. , 1975). In this paper, therefore, we present a particular data analysis method of using Poisson statistics in order to find the millisecond variation from the balloon borne experiment data. And the method is applied to the data of Flare event obtained from the observation of the hard X-ray from Cyg X-l with the balloon in 1972 and the result is presented. 2j Experimental. The experiment was launched from Sanriku Balloon Center (JAPAN) on the 7th of October 1972. The balloon floated at a constant ceiling altitude of about 5.6 g/cm from 7n20mU.T. of Oct. 7 until the cutdown at 1ГзЛ.Т. of Oct. 7. The balloon payload comprised two scintillation counters consisting of Nal(Tl) crystal of 3.0 mm thickness x 76 mm diam eter and a plastic scintillator of 30 mm thickness x 76 mm diam eter sandwiched between the crystal and a photomultiplier (in detail mentioned at the 13th jI.C.R. conference by Nakagawa et al A cellular collimeter with 4°(north to south) x 20°(east to west) FWHM field of view was set on each counter. An Am^^l radioactive source which has two prominent lines at 16.8 keV and 59.5 keV was provided as in-flight calibration with the comand 162 signal from the earth when the detector systce chafed the №>J*> The aspect of the two set of thii detector system, with each axis inclined at 4* with respect to the lonlth so m to obtained the most favourable conditions for the observation of Cyg X-l, is controlled to point alternately to north (OKF-node) and south (ON-mode) with the comand signal from the earth. At ON-mode the X-ray source at which is taken aim, which is Cyg X-l, lies in the field of view of the detector and at OKF-node Cyg X-l can not seen. The charged particle is removed using the difference uo in the pulse rise Ocl T OT time of the Nal(Tl) crystal and the plas tic scintillator. So X-ray signals from each detector which are satisfied each circuit condition are 1 transmitted in an analogous way in two separate channels with a timing acuracy -- + + + of 1 msec. JL. JfL A U.T. J 3. Result and Discussion. Fig. 1 Tyne variation in the X-ray flux Fig. 1 shows the at the top of atmosphere from Cyg X-l. source intensities The source spectrum is assumed a power plotted versus U.T. law form with index -1.7. Taking into account attenuation of atmosperic thickness along the line of sight the intensities represent the photon flux at the top of atmosphere averaged the energy range 30-70 keV assum ing the source spectrum on a power law form with index -1.7. In this figure, error bars are one standard deviation in the count ing rate over the presented period. At 10n05mU.T. the X-ray intensity of Cyg X-l suddenly increases by factor 4 and this is Fhoton bwgjfcrt» Fig. 2 Time variation of the spectral form. The entire data train is classified into four except the Flare section, (a),(b),(c) and (d) is в^З!1"-, 8n47m-, 9n16rt-, and 9h33mU.T.-, respectively. IftJ namenod "Flare". A» shown this fttjuto. the Flare event be«)ipe neair 9пn50S0^J.Tл1и.Т. and goos down to th« ordinary level near 10 lb U.T,, so the Flare event keep» collectively on about IS minute» In order to 9et the varia FLARE tion of spectral fori», the entire data train are divided in five section! including the Flare event, and five number spectrum» are shown in Fig. 2, 3. Assuming the source spectrum on the power law form, and taking account of the depth of atmosphere and win dow thickness along the line of sight, counter efficiency including escape effect, and counter resolution, the number spectrum is gotten. And the line shown for the number spec trum is the best fit to the data. Fig. 2 shows the number spec trum of four ordinary states except of the Flare event. The averaged index of four spectrums is-1.8±0.2. Fig. 3 shows the number spectrum of the Flare event, and the index is-2.230.2. 30 50 70 Compared the index of the Photon Energy (keV) average and the Flare, the spec tral shape is not difference Fig. 3 The number spectrum of between indices within a stand the Flare section. The dotted ard deviation. line is the best fit to the _ .. „, . .. data with index -2.2i0.2. „ For the Flare event, the exsistence of the millisecond variation is important to consider the mechanism of Cyg X-l. Therefore, at first, the power spec trum method using the Cooley-Tukey fast Fourier-transform al gorithm is applied to the OFF-mode above mentioned, the ON-roode in theordinary state, and the Flare event. These results are shown in Fig, 4. In generally, if it is that X-rays fall on a time series at random, the counting rate should obey a Poisson distribution and the power density versus frequency should be constant. Then, the OFF-mode in-Fig.4c is likely to occure in a white-noise sample. As shown in-Fig.4a, the power density spec trum for the Flare event happened periodicity at 95 msec. The other hand Fig.4b shows that the ordinary state is not so sim ple periodicity. As a further check on the sensitivity of our experiment to pulsated х-radiation, we exami.ie a particular method through the medium of Poisson statistics for the Flare event. In Fig. 5 the histogram descrived with solid line, the lower shows.the number of X-rays events recorded for each time spacing and the upper the dispersion of the number of events from the expo nential distribution drawn with solid line curve on the lower 164 part of figure. to 20 so *o SO 100 ISO 200 FREQUENCY (W * 10 Hi) TIME SPACING (msec) Fig. 4 (a) and (b) shows the Fig. 5 For the Flare event, the power density spectrum of the distribution of tha time dura Cyg X-l observation for the tion and the dispersion (upper Flare event and the ordinary portion of figure) of that from state, respectively, (c)Power the exponential distribution. density spectrum for the OFF- mode. £; Conclusion. We point out explicitly tha* the Flare-like event detected at high energy range 30-70 keV with the balloon borne detector-. The millisecond variation of the intensity of the hard X-ray from Cyg X-l, in spite of unfavourable S/N value comparing with the data obtained from a rocket borne experiment, is likely to occure in the Flare- event with the periodicity of 95 msec. But becouse of poor statistics, it is hard to say clearly of the existence of periodicity in the Flare state. So, to the Flare event obtained by other groups a particular method should be applied. Then, in order to say clearly of the milli second periodicity, it is necessary to observe the intensity of the hard X-ray from Cyg X-l with the larger area detector (over 1000 cm2). 5_: Acknowlegement. The authors are very grateful to profes- sor J.Nishimura and members of the balloon flight group of Institute of Space and Aeronautical Science, University of T"4yo 165 for the balloon flight. 6. Reference». Agrawal, P.C., Gokhale, O.S., Iyengar, V.S., Kunte, P.R., Manchanda, R.K., and Sreekantan, ».V. 1972, Nature 2)1,22. Friedman, H., Frits, G., Henry, B.C., Hoilinger, J.P., Meek in», J.P., and Sadeh, D. 1969, Science 22_1, 345. Foligni, P., and Prontere, F. 1973, Astro. ( A»trophya., 2§_, 173. Holt, S.S., Boldt, K.A., Schwarti, D.A., SerleaUtaoa, P.J., and Bleach, R.D. 1971, Ap. J.(Latter»), 141, L119. Nakagawa, M., Sakurai, H., and Uchida, M. \VTS, 13th Int. Conf. Cosadc Rays 1, 73. Nakagawa, H., Sakurai, H., and Uchida, H. 1975, ISth Int. Conf. CoMklc Rays 1, 139. Oda, H., Gorcnstein, P., Gursky, H., Kellogg, B., Schraiar.B., Tananbau», H., and Giacconi, R. 1971, Ap. J.(Latter*),16* LI. Rapparport, S., Doxsey, R., and Zaumen, W. 1971, Ap. J.(Latter»), 168, 1.43. 166 ааскст мкашна* т moria a-aar MWKU И МГТ B-MV CKMK* я.1. МЦ •.». »MMk »•*.«. Nnt! ».»• a. Мм raaulta •* taa aaakat 'llahta mat Ш aaft И« ма »iNiiU«Ml IWMWI MMIUN «-муа in tha MMfy миаа 0*i to MA keV* Ifcf fit** fitaht aatoatot «ita • total мм #f lit aa* aw a I IUM alaf al tMlatUn wMw. lie* HUIIIUH a f lai* ef »«м iri'iir rmm fa* tnu Ik* мм* tti*ht aaplayaai gwMm havlM • total ana a» 4N «P Mi M I dmn aalyataaylaM гмИШи vintav* The MUiMto* fa* ihia aaylaai MftMi a Malt #f via» af •• a It* ПМЯ. Им aaaalta ГАМ ма to* ttiaht pattal* to ttea мП *-rey aaiaalana faaa ИС «-1, IMHMHM JM. Thaaa aeaulta eta eaaainad in the light af aattMt fat thaM aauraaa* 0C 1.3 (Я-му tataanaay) "taf. U."., **a Diractat 18Я0 iatailita Cmln Рмпуа М2140 Bangalore* Inaia 167 raLIMTNJlRT WSULTS ITtOH ТЖ *tC« i SATEU-ITT M. MEL. A.R. lAZER-lAOn. C. VEDUXXE Centre d'Etude Spatial* da* Rayonneawnt• Toulouar Theoretical Q Experimental (*] B • One of two experiment» aboard the French Satellite SICNE ^ ia an X-ray detector. The uae of the double counter technique, with 2 Hal scintillator е (1^2 inchei in diameter) in a Cal shield allows good directivity over the whole energy range, 20 KeV to a few MeV. \ \ The satellite is to be launched on June 15 by a soviet rocket The orbit will be circular at 500 km, and 51* inclination. The experi aent, aiming 10' free the antisolar direction, will scan a 30* wide band around the ecliptic, at the «pin rate of the satellite, 1 turn in It minutes. Real time analysis can be performed on part of the data and result* available at the time of the onference will be presented Coordinates: 0G -1.3 Х-Ray Astronomy Maffin» address: CENTRE D'ETUDE SPATTALE DBS RAY0NNEMENTS 9, avenue_Colonel Roche 31400 TOULOUSE France IM SOME ASPECTS CMCEWIIW3 ЛССЖГПОИ J* ИШТ STSTCRX C«Wt Helr«lles Fllho Institute AstronoaUco • Geoffsico da U.S.P., «reail Tbaontical Ц] Exsarimeatal Q »»«» Q A gas flow In a coaipact binary system owing to non apharlcal accretion la studied. Aa uaually, thla flow stay ba analysed considering a dlac, the Alfven Surface, and tha region Inner to the Alfven Surface. In the outer region, one doea a detail ed treatment of the turbulence; the formalism deployed here allows the treatment of non stationary discs. Fro* the Alfven surface one assumes the flow parallel to the magnetic field and infinite conductivity for the gas. All the relevant Physi cal parameters are obtained. Coordinates: OG 1.3{X-Ray Astronomy) Mailing address: Professor Cesar Meirelles Filho, Instituto AstronSmi co е Geoflsico, Universidade de Sao Paulo Caixa Postal 30.627 - 01000 - SAO PAULO - BRAZIL 169 COSMO LOGICAL COSMIC RAYS AND NUCLEAR >-RAY EMISSION T. Montmerle Service d'Bectronlque Physique, Centra d* Etudes NucKalres d« Saday, Trance In the hypothesis that coamologtcal cosmic rays ("CCR : protons and or-particlea) have existed in the early untveVse (a~100), low energy Interactions take place between CCR and the ambient gas. Light nuclei are formed, in part via excited states. The states which decay through the emission of nuclear у-г*У lines lead to a contribution of these protons to the X-ray diffuse background. The time-dependent transfer equations for 4He*. *U* and 7Be* nuclei and y-ray lines are solved in the case of a CCR burst. The resulting nuclear v-ray spectrum now is shown to be at least 6 orders of magnitude below the observed X-ray background. I. COSMO LOGICAL COSMIC-RAYS AND OBSERVATION The cosmologies! cosmic ray (CCR) hypothesis has. been pro posed in its original form by Stecker (1969, 1973) to explain the shape of the 1-100 MeV v-ray background spectrum. Gamma-rays appear as a result of tt* decay, following high energy (~3 GeV) inter actions at high redshifts (r~l 00) between the CCR particlt-в (p and e) and the ambient gas (H and *He). The context of the CCR hypothesis and the framework for particle transport have been discussed in a companion paper (these proceedings, OG-17, hereafter Ml) and in Montmerle (1977a, here after Ma), to which the reader is referred. Noting that light elements, essentially "Li and 7Li, are produced by low-energy (M0-100 MeV/n) aor reactions, the allowing results were obtained. First, the CCR flux is normalized so as to account for the у*гаУ background spectrum derived from the recent observations of the Apollo (Trombka et al., 1977) and SAS-2 (Fichtel et al.. 1975) spacecrafts. Then the 7Li abundance 7Ld/H, calculated with a " total energy" CCR spectrum, agrees with observations, but °Li is overproduced by a factor of ^-5. In view of the wide spread of the Y-ray data as a whole, it is however'~poesible to lower the CCR flux so that the calculated "Li abundance remains compatible with observations. Of course, this is at the cost of relaxing the simultaneous explanation of both the \-гжу spectrum and the (otherwise unexplained) 7Li abundance. Also, the CCR flux now in intergalactic space is smaller than the galactic 170 cosmic-ray (GCR) Пиж by several order* of magnitude. (Sea alao Montmerle 1977b, hereafter Mb.) la view of this noanegative result, it it of interest to study other cooitqutocti or obaervational teats of tha CCR hypothaela. Among them, are deuterium production, by low-energy pp and pa reactiona (theae proceedings, OG-130, her «altar M2), and the produc tion of nuclear y-гауе, on which w« focua in the preaent paper. Now the nuclear y-raya produced by Interactions of CCR with the ambient gas have reat energiea in the range ~500 keV • 30 MeV {§ II). There fore, becauae of the redehifta involved, the corresponding photons contribute to the photon background down to the aoft X-ray range (§Ш). However, this contribution will be shown to be very a mail (§ IV). II. NUCLEAR y-RAYS a) Energy levels. Low energy pp, par and arar reactions (here after par - pa reactions) give rise to nuclear у-г*У emission via the decay of excited states of the CCR-produced light nuclei ; y-rays can also arise from inelastic scattering. The light nuclei involved in par-pa reactions and the references where their energy levels can be found are D (e. g., Blatt and Weisskopf 1952), ^He (FUrman and Hanna 1975), *He (Fiarman and Meyerhof 1973), and ^Li, 7Li and 7Be (Lauritsen and Ajzenberg-Selove 1966). As discussed in Montmerle (1977c, hereafter Mc), and taking selection rules into account, we are interested in the *He*(27.4 MeV) and 4He*(30.5 MeV), 6Li*(3.562 MeV), 7Li*(0.478 MeV) and 7Be* (0.431 MeV) states. The other states decay mainly or entirely via particle emission. We will not consider the D and ^He nuclei, because of their low yield in the CCR hypothesis (see M2). b) Cross-sections. The cross-sections cv j for producing nuclear Y-ray lines via the decay of nucleus C* is related to the cross- section a p* for producing the nucleus t in the relevant excited state by ov j * a j* (Гу/Г), where Гу is the width with respect to Y-гау decay of the' level of total width Г (e.g., Blatt and Weisskopf 1952, Evans 1955). The o g* cross-section is known experimentally, or may be roughly estimated when the total cross-section a( for production of the nucleus t is known. A summary of the relevant available data and estimates is presented in Table 1. Note that the 4He*(30.5 MeV) line has the same characteristics as the 4He*(27.4 MeV) line (both have J = 1", T"l, in particular, see Fiarman and Meyerhof 1973), hence their cross- sections should be very similar. Ш. THE NUCLEAR Y-RAY FLUX As has been shown (Ml, Mb), the CCR hypothesis involves maximum redshiits (z ) in. th* range 60-150. Since the rest energy 'of the lines (Table 1) range from "400 keV to ~30 MeV, tUe nuclear 171 Сг—-aactl—a fat ш1мr r-ter lUa —iaaiaa U - -»a raatti—i. cltetaS4 " taaritMrtyr Г./ГДГ laflaf.. •« aaa leval QWP) T (ikm tbrsabalO TU 27.4 (1) < 10 a* ••t.. fraa (1) *U 3-5*2 1 C3) 1/2 a, (-1 a*) ••t. (4) 7 U 0.47* 1 <3> 1/4 •, i «at. (5), fro* 7 ' («)j alaa (7) >« 0.431 1 (») 1/4», Kaferancea to Tabla 1 (1) toslovaky I., laaaty Ж. 1*74, Aatr. 4». 34, 477 (2) liabart Ж. 19», Fay*, tor. J02, 1104 O) La»rlt»«n I., AjsaaberarSalove I. 19M. Kiel. РЬуш. 78, 1 (4) HsntaarU T. 1977, preprint (5) Koalovslr/ 1., hatt; X. 1974, др. J. (Latter*) _[?!. L43 (6) БигсЪаа V.E., at «1. 1958, Kiel. Phya. £, 1*1 (7) Kins C.H., et al. 1975, Phyt. tar. Letter* 35, 988 v-ray flux now extendi from a few MeV down to a few keV, i. a. lies eaaentially in the X-ray range of the coamic photon background. Becauae the relevant energy level» have a very abort lifetime (~10-21«), their decay may be taken as instantaneous ; then the trans port equation corresponding to eq. (2-1) of Ml for excited t* becomes : oN*jH(E.*)/ot = Q*e)H<£,«). (3.!) (The source-function Qee( jj(E, z) involves the ce* cross-sections mentioned in- § lib ; see Mc.) Neglecting Doppler broadening and Compton scattering (see Mc for discussion), it can be shown that the contribution to the observed X-ray background flux due to nuclear v-ray emission from nuclei £* at a given energy Ev is . Il(E ) - F .•,C aE»° Г-Г7У Го*. „(E,zl)(iE v v 1 2 J el H ^4*H0E5(1 +2%гЬ ' o ' (3-2) e e for E v/(l+*,)SEv:£Ev , and I*,(EV) • 0 otherwise. In eq.(3-2),Et is the rest energy of the line, s' is given by 1+z = E-$/Ev. Other Constants are the normalization factor Fv to the 1-100 MeV у-г*У background spectrum, and the number density of the universe now Пд Q. 172 Fig.I The contribu tion of CCK-produeed nuclear Y-r«ye it shown in heavy lines •у for the nuclei indi cated. Tha observa *:: tional data li« wit- bin tha hatched areas (see text for refe rences). Characteri stic energies are shown along the abs cissa. -ць .«и. щд1 turn 'WE? "n? IV. CONTRIBUTION TO THE X-RAY BACKGROUND The nuclear v-ray flux Iv resulting from the decay of the nuclei indicated on Table 1, and calculated according to eq. (3-2) is shown on fig. 1 (from Mc), for the typical caae z, = 100 and qo=0.1. Also shown for comparison is a summary Jx of X-ray observations (Hortsman et al. 1975) along with the Apollo (Trombka et al. 1977) and SAS-2 (Fichtel et al. 1975) \-ray daU. It can be seen that there is at least six orders of magnitude between Jx and Iv. Thi* justifies a posteriori the approximations made in the computations (§Ш), but makes the prospects of detecting Iv (and especially the low-energy edges of the contributions to Iv of each nucleus) quite pessimistic indeed. Note that, because of Compton scattering, the edges should be less sharp in reality than they appear on fig. 1. Thus, nuclear y~r*y emission is seen to be by far compatible with the CCR hypothesis, but it is practically hopeless to detect CCRi through their nuclear v-ray line*. UwriciM T., «jnakm-Selovt Г. KM, Huel. Мис J.H., Watitkepf ».T. I»», Tb«or.ttc.l Huclw fhwtc» pfcy». 71, I OKlqri «w York) NntMrl* T. 1*77». As. J„ is >пм Ota) torn* *.». НЦ.ЛЪ» Ata-lc qcl«». (MeOrw-IiUi м, totk) l»77». A». J., U »r*H (ab> Папи S„ Шии f.t. I»7S. Duel. Tkr». A2M, I — H77e, >M»rLc (He) Пагам I., «qmkof U.S. 1*73. lacl. ПвТПаМ. I Itteker I.V. IK», I*uc* 22*. (TO Tlektd CI.. «Sj\. I»7I. A,. J. Iff. 1« , |»7S, la шПР-З». ,. 211 •теме CM., «t J. I»75. li». ЛГ». Cla. t. 2» tnakk* J.I., «t «1. I»77, A». J. V& MS SPECTRUM, TIME STRUCTURE AND DIRECTION OF INCIOENCE OF THE AUGUST 16, 1976 GAMMA RAY ВиПгИ Michael Sommer, Max-PIanck-InstItut fUr Physlk und Astrophyslk, Institut far extraterrestrische Physlk, D-8046 Garching, FRG Dietrich Mtiller, Enrico.Fermi Institute, U. of Chicago, U.S.A. Henry Horstman and Loredana Bassanl, Istltuto dl Astronomla, U. of Bologna, Italy Two major bursts of energetic photons have been re corded with a new balloon-borne Instrument during the second transatlantic flight In 1976: One In coin cidence with a type III solar radio burst on August ' and a very energetic gamma ray burst of non-solar origin starting at 16:15.5 UT of August 16. Spectra! information of the gamma ray burst has been obtained up to 2 MeV. A crude position of the burst source has been derived from data of a directional detector array after correcting for absorption and scattering in the earth's atmosphere. 1. Introduction. Although more than 70 gamma ray bursts of non- solar origin have been observed, no burst source has been identi fied and even the question whether gamma ray bursts are of galactic or extragaI act Ic origin has not yet been answered. This is partly due to the fact that at present three satellites at appropriate locations are necessary to derive directional information from •"ival time measurements. This paper presents data from a new •rgle instrument capable of determining unambiguously source po sitions within a few degrees. This Instrument has been tested In two balloon flights, one short flight In November 1975 observing Cygnus-XI during a transition from' the ''low" to the "high" state (ref. 1) i?- I 174 •:'. Instrumentation. Fig. 1 shows the apparatus designed by *he> Max-P I anc k- I ns t i t u t (ret. 6). In the upper section 5 detector-, з • 100 cm NaltTI) 5 mm thick have been mounted so that tho central detector is pointing to the zenith, the tour others being incl.nuc; by 45 to the vertical direction and rotated by 90 each in a*i- muthal direction. All the housings of the PM-tubes are shielded against photons from the rear by a graded shield. For X-rays belo» 100 keV the detector can be described by a flat totally absorbing area. The measured intensity of a photon burst Is proportional to the projected area of each detector. The direction of incidence and the intensity of a burst can be derived from signals of at least three detectors. The pulses were counted In fjur Integral energy channels ranging from 15-30, 30-60, 60-150. and > 150 keV energy .* The accumulation time of 6 sec was decreased to 0.5 sec when ever a significant increase of the radiation below 60 keV had been detected. This set of de tectors has been completed by 2 a detector of 1300 cm of Cs I (Tl ) 1 r.I . thick provided by the Enrico iermi Institute tor energies up to 2 MeV. 5 3 A 6-10 m Winzen balloon car rying the instruments has been launched from Trapani, Sicily, on August 13, 1976 by NCAR and' the Transatlantic Balloon Fa cility, an international col laboration of various insti tutions and groups based on the original proposal by Fowler et al. (ref . 2). During most of the 107 hours Fig. 1 The dlrectiona'l detector of the fl'ght to Gardner, Mas- array with 5 Nal(TI) crystals, electronics and a camera for data sachusets, the balloon stayed record!ng. 175 .: . .!• .in altitude o1 58 km and ha-, boon suci'osstul I» t roc kod by iho "< . J i у s t em . •'<• gondola was suspended vertically from the balloon, a/lmuthjl i-i'ormation being derived (rom sunsensors and tnagnetotnet er s. All o.it j were recorded on film, only the analog signal being used to speed up the data recording In the case of a burst has been trans mitted on the low bit rate tele.netry link tor test purposes. 3 • Time Structu re. The gamma ray burst of August 16, 1976 starts to exceed our background counting rate at about 16:15:28 UT. The fast readout mode of the camera recording system starts about 1.5 sec later resulting in a 0.5 sec time resolution. The rise of the main burst is structured, the burst source seems to oszi I I ate with increasing amplitude. Three smaller bursts follow the main burst which lasts for about 8 sec. After 31 sec the counting rate returns to the preburst level again (Fig. 2). —i 1 1 r—i i * !*• 2000 V -ч_ Fig. 2 The time profile of the gamma ray burst of August 16, 1976. The upper diagram shows the counting rate of the five directional detectors between 30 and 150 keV energy. The lower curve shows the time profile on a nonlinear scale as derived from the fast readout trigger signal. 176 4. Direction of Incidence. The appearance of the data for the cos mic burst suggested» both because of the hard spectrum and the quasI -1sotropy among the detector signals a highly scattered con tribution and In particular a low elevation angle of the burst source. These facts required a detailed numerical compute i-1 on of the total fluxes (direct and scattered) reaching each detector assuming different elevation angles for the source and reasonable Input spectra at the top of the atmosphere (see also ret. 3). With a best fit procedure we derived a 90% confidence contour In de tector coordinates. The error box In the table below and In Fig. I was obtained by combining these limits with systematic uncertain ties In the frame of reference of the gondola. R.A. 6 l" b" 83.2° -11.0° 214.4° -21.9° 87.4 - 6.7 212.2 -16.3 105.4 -20.6 232.7 - 6.7 101 .7 -25.6 235.8 -11.9 Tab Ie 1 Error box of the gamma ray burst source. Due to the unknown input spectra on top of the atmosphere and the lack of exact data on the shielding of the detectors for low ele vation angles, uncertainties In the position remain. The ele vation angles of the burst source as shown In Fig. 3 may exceed the error box by a few degrees. The data of the X-ray flare of August 13, 1976 (11:49-11:55 UT) have been used to check the Monte Carlo calculation as well as the measurement of azimuthal direction of the gondola. The true position .of the sun (elevation angle 65.3°) coincided within 2 degrees with the position derived from the X- ray data. This fact gives confidence in the method to derive po sitions of impulsive sources and in particular in the position of the cosmic burst reported here. Recently Mandolesi et al. (ref. 4)' reported a possible coincidence of a radio observation with the gamma ray burst of August 1.6, 1976. The radio signal appeared at about 16:16:30 UT and was assumed to arrive from a direction close to the sun but not from the sun. 177 Since the sun (r.a. - 146.2°, 6 - 13.55 at that time) was far away from our burst position we believe that the two observed phenomena are not correlated. The fact that the SAS-3 satellite was in the sun and operative but did not detect any signal from the gamma ray burst gives preference to our direction since this was occulted by the earth for the detectors on this particular spacecraft. Unfortunately the HELIOS-B spacecraft was behind the sun and no data have been obtained In August 1976. Due to the long baseline between the satellite and the earth one could have expected to define a burst position within less\than 0.5 square degrees after combining the Helios-Solrad data with our dIrectlonaI data. 5. Spectrum and Tota! Energy. The high energy data, being less sen sitive against scattering in the earth's atmosphere, suggest a dif- -2 ferentlal power spectrum of Е ranging from 150 keV up to 1.8 MeV. The slope of the spectrum Is in good agreement with the Irap-7 data, (ref. 5). The total energy can only be calculated as a function of 178 the elevation angle under which the burst source appeared In the frame of reference of the balloon-borne Instrument. The results are listed in the following table: elev. ang 1 е 300 - 1500 keV 100 - 1000 keV 8° 0.7...1 • I0-4 1 . . .1 .4 • 10-< erg/ся2 14° ~ 5 ' 10"5 ~ 7 • 10"5 erg/cm2 A reasonable fit to the low energy data has been achieved assuming the source spectrum (entering the atmosphere at about 80 zenith -1 -2 angle) to be a power law Е breaking to Е at 220 keV and fol lowing it through interactions In the atmosphere to get the ex pected horizontal flux. Ac know Iedgements: We like to thank Or, Ramsden and his associates from the U. of Southhampton, the staff of the Appleton Laboratory and of the Italien C.N.R., the NCAR launch and recovery crew. Prof. Scarsi and Mr. H. Neuss for all their work In Enabling a transatlantic baI Ioon fIIght. We like to thank Prof. Plnkau for his encouraging Interest In the project and K.H. Schenkl, S. Reiter, A. FInkenzelIer, B. Fransen, and J. Nagerl for their help In realizing It. References 1 Somnter, M., H. Maurus, R. Urbach, Nature 263, 752 (1976). 2 Fowler, P.H., J.H. Davles, R.E. Baker, 0. Ramsden. , "A Transatlantic Balloon Facility", Proposal of the U. of Bristol and Southhampton (1974). 3 Horsiman, H., L. Bassani, E. Horstman-Morettl, submitted to Astrophys. Space Science (1977), 4 Mandolesl, N., G. Morigi, P. Inzanl, 6. Slronl-, F.S. Del I I Santl, F. Delplno, M. Petessl, A. Abraml, Nature 266, 427 (1977) 5 Cllne, T.L., U.D. Desal, private communications. 6 Sommer, M., to be published. 179 OBSERVATION OF САЖА RAY BURST AT BALLOOK ALTITUDE J.ftlahlmura*. M.Oda», S.Miyamoto*, Y.Ogavara*. M.FuJil*. T.Yamagaml*. T.Tavara*, M.Yoshlmorl**, H.Murakami** M. Nakagava*** and T.Sakural*** * Institute of Spaca and Aeronautical 61cencc, University of Tokyo ** Dep. of Phyalca, fclkkyo University *** Dep. of Phyilca, Osaka City University, Osaka Theoretical Q Experimental Щ Both Q Long duration balloon flights for 150 hr in total were perforated by three balloons at Sanriku Balloon Center to observ Cosmic Gamma Ray burst in 1975. The -detectors used are each having three counters with the cross rotation modulation collimeters proposed by us at the time of Munchen Conference, having wide apeature but also have high resolution to locate the direction of the burst source. During those flights, an event having the feature of gamma ray burst was ovserved on Sep. 23. The energy of the burst is about 10"6erg/cmJ, and by using the modulation patent the location of the source is analysed and found to be near around Cyg X-l. Coordinate*: Mailing address: Jun Nishimura : Institute of Space and Aeronautical Science, University of Tokyo, F.omaba, Tokyo, Japan 180 HIGH ENERGY GAMMA RAT BURSTS S.K. Gupta, P.V. Raman» Murthy. B.V. SrMkantan and S. C. Tonwar Tata Institute of Fundamental Research, Colaba. Bombay 400 005. India Theontietl p Experimental (X| •e,h D We are looking for possible occurrences of rapid sequence of coinci dences between three horizontally separated scintillators each set at a threshold of one singly charged particle level. The electronics system is capable of recording successive eveat-to-event elapsed time Intervals to an accuracy of 1 u.a upto a maximum of 10 intervals. A rapid sequence of coincidences that is beyond the Poissonian statistics could be a candidate for a Gamma Ray Burst event in which the individual gamma rays have energies 13 "> S. 10 eV. The same set-up is also used to record bursts of gamma rays with energies &500 MeV incident on terrestrial atmosphere by looking for rapid sequences of low energy (» 2 MeV) gamma rays in a single scintillator at our level (800 g. cm'2) of observation. Results j"KgJT««H unto date will be reported at th» Г.тЛлтлп^ Coordinate*: OG 1.4 (Gamma Ray Bursts) Maillncaddrtu: Professor P.V. Ramana Murthy Tata Institute of Fundamental Research Colaba, Bombay 400 005, India. 1В1 A SEARCH FOR y-RAY BURSTS FROH THE EXPLOSIVE EVAPORATION OF BLACK HOLES T.C. Veeket,.:Centra for At t rophyt I ci, Cambridge, Иш., U.S.A. N.A. Porter, University College, Dublin, Ireland. ABSTRACT ' Separated «taotpheric Cerenkov detectors have been used to search for bursts of*Y-rays, not essoclated with single air showers, which Have been predict») theoretically from the explosion of small black holes. Characteristic detectable energies are In the region 200 - 1000 MeV, but higher energy photons would also be detected. Detection sensitivities are In the region 10"13 - 10"11 Joules/metre2 for bursts of duration less than 1 microsecond. In a preliminary analysis an upper limit for primordial black hole explosions In the galaxy is set at 0.01» events/po^-year. Some Information can also be obtained on longer duration bursts, which may occur for different nuclear moc'eis of the black hole explosion process. 1. INTRODUCTION It has been suggested recently (Ref 1) that black holes of mass In the region of 1012Kg will evaporate In a time of the order of the age of the universe, and that the final stages of evaporation will be very rapid. Such black holes, If they exist, must be primordial. The details of the evaporation process will depend on unknown nuclear characteristics, but for the Hagedorn statistical bootstrap model (Ref 2), about 1028Joules of energy will be released In the final explosion In a time of the order 10*7sec. About 10* of this energy will be produced In y-rays of energy 200 - 1000 HeV. Upper limits can be deduced from the Intensity of the background Y-radlation for the rate at which such black holes are exploding in the universe. The number of explosions In our galaxy may be higher by a factor of Ю6 because of galactic clustering, but nevertheless their detection by balloon-borne or satellite-borne у-ray detectors would not be feasible with present system., because of the relatively small 182 sensitive «rest available. An Improvement tn the upper limit» nay however be sat by tbe use of the atmospheric Carenkov technique, where the detection area Is of the order 105 metres2. A limit has already been set by the present authors (Ref 3)> using three coincident detectors separated by l.3Km to reduce the contribution from single ilr showers. The present paper reports two further experiments, one using four instead of three detectors, and carried out at the Mount Hopkins Observatory In Southern Arlzo:u. The other used two larger reflectors: the 10 meter reflector at Mount Hopkins, arid the large solar furnace 400Km away at White Sands In New Mexico. . An indirect method of detecting explosions has recently been suggested (Ref 4). The relatlvlstic shock wave expanding Jnto the ambient magnetic field may produce an electromagnetic pulse, whose characteristic frequency will be determined by a parameter Yf = 2x10l3Kg/Mcrit, where Merit is the mass at which the black hole explodes. This depends on the particle physics; for Yf'107 the pulse will be at optical frequencies. The first of the experiments reported here would be sensitive to such optical pulses, as well as tc the Cerenkov light produced by -y-rays in the atmosphere. 2. LOCAL EXPERIMENT This experiment was local In the sense that all four detectors were connected by high-grade co-axial cable to a single recording station, and prompt coincidences were taken between them. The Fields of view and separations were chosen, however, so that the four sensitive cones In the atmosphere did not all overlap up to heights where shower development Is significant. Hence a single extensive shower would only trigger all detectors if it had a very unusual trajectory with considerable emission of light at angles away from the axial direction, or was extremely large; but a short burst of Y-rays extended In space at the top of the atmosphere could produce a master coincidence even for-y-ray energies In the region 200 - 1000 MeV. 183 Each detector hid • reflector with diameter 1.5 metres, and a single 0.125 metre photomuItIpI Ier at the focus (RCA <»522). The full field of view was limited to The experiment was operated for 37 hours, and 2 four-fold coincidences were recorded. These 2 events were considerably In 7 excess of random expectation, which was about 3x10" events. We ascribe them to extensive air showers, excluding y-ray bursts on the basis of results to be described in the next section. The possibility of optical pulses predicted by the Rees model (Ref 4) cannot be excluded on this basis, however. It will be shown In a later section that this result corresponds to an upper limit of 2.1 events/pc3-year for the rate of pbh explosions In our galaxy. 3. LONG BASELINE EXPERIMENT The local experiment was limited by the possibility that infrequent but not negligible large air showers could trigger all detectors.. In order to remove this possibility two detectors were operated over a separation of 400Km, the 10 meter reflector at Mt. Hopkins, and the solar facility, of comparable size and aperture, at White Sands, New Mexico. To reduce the chance coincidence rate from air showers, advantage was taken of the fact that the angular size of the light spots from Y-ray bursts will be considerably greater than that of single air showers. 184 At the focus of th« '0 meter reflector, invert photomul11 pi I«rs were mounted, and a seven-fold coincidence was demanded to produce an evnnt. The photomuI11pI Iers (RCA 4522) covered a field of approximately 3° diameter. At the solar furnace, four photomul11pI Iers slightly separated In angle were used, and a four-fold coincidence was required to produce an event. The field of view covered was again approximately 3° In diameter. The times of events were recorded at each station on analogue tape recorders, with a high precision 5KHz time signal, and standard 1-minute time markers. Master coincidences between the two stations could therefore be established to an accuracy better than 1 millisecond by comparison of the tapes. In addition a direct telephone link, sending signals from White Sands to Ht. Hopkins, was used to operate a coincidence circuit with resolving time 10 milliseconds, so that an immediate record of master coincidences could be obtained. Slow variations in current from the photomultiplIers were also recorded on separate tape recorders, but have not yet been analysed. The systems were operated In coincidence for 22.5 hours under clear skies. In a preliminary analysts, no master coincidences were observed simultaneously at both stations. We will show that these figures Imply an upper limit for primordial a black hole explosions in the galaxy of 0.0l|/pc year, assuming the Hagedorn model (Refs 1,2). k. INTERPRETATION k . 1 Local experiment From the events observed we can calculate an upper limit for the rate of pbh explosions in the galaxy. The Cerenko/ light output can be calculated for yrays In the region 200 - 1000 HeV. The most reliable calculations are probably those of Turver (private communication) who finds that a 1000 HeV у-ray produces 6000 photons within 2° of the primary direction, and we will adept this figure. From air shower studies and from calibrations with a standard light source we estimate a threshold sensitivity 185 for the detectors of 370 photons/m2. From Turver's figures this gives a primary v-ray threshold of IxlO-11 Joule/m2. For a Hagedorn type explosion releasing 1027 Joules of energy we have then е maximum detectable distance given by r2 » 1027x10ll/4*m2 , from which r * 2.8x10l8metres or 93pc. The sensitive volume for detection is V • Пг3/3, where £2 Is the solid angle for detection. We take this conservatively as equal to the geometric field of view, 4xl0~3 sterad. Hence V ". 950pc3. The two events observed correspond to an upper limit at the 99% level of 8.4 events in the 37 hours of observation. Hence we arrive at an upper limit of 2.1 exploslons/pc3-year. 4. 2 Long Baseline experiment The sensitivity of the solar furnace, which limits that of the experiment, was calibrated at 26 photons/m , and using the same analysis as before this gives a primary sensitivity of 5.6x10-13 Joules/m2. For 1027 Joules of energy released the maximum detectable distance r = 390pc. The true field of view is greater than the geometric value but difficult to determine, so that we take the geometric field, giving a sensitive volume of 4.4x10 pc. The upper limit corresponding to zero events at the 99* level is 4.3 events in 22.5 hours, or 1750 events/year. We therefore arrive at an upper limit for the explosion rate of 0.04 explosions/pc3-year. 4. 3 Opt i cal pulses The possibility that the pulses observed in the long base line experiment were due to the optical signals predicted by Rees (Ref 4) cannot be excluded by the long baseline experiment, because the latter system will not respond to point sources. The predicted energy release in the optical Is 1023 Joules for Yf"107. The threshold sensitivity of 370 photons/m2 corresponds to 1.8xl0~16 Joules/m2. This gives a maximum detectable distance of 220pc. With 6 = 3-9xl0"3 sterad the detectable volume is l^xlO^pc3. The 2 events in 37 hours correspond to 473 events/ year, giving an apparent rate of 0.03 events/pc3-year. The upper limit at the 99% level is 0.14 events/pc3-year. 186 5. CONCLUSION We find no evidence for Y-ray bursts from primordial black holes or other sources; If the Hagedorn model Is used the upper limit Is 0.0^ exploslons/pc -year. If the diffuse Y-ray back ground Is due to pbh explosions (Refs 1,5) their clustering In the galaxy must be less than about a factor of 10 . Two events observed In the local experiment fre ascribed to extensive air showers since they would correspond to a rate greater than that set,by the long baseline experiment. The possibility that they were due to optical pulses according to the Rees model cannot be excluded by the lony baseline result. 6. ACKNOWLEDGEMENTS We wish to acknowledge support from the Smithsonian Research Foundation and the National Science Council of Ireland, and to thank Mr. E. Horine for assistance in the experiments.' 7. REFERENCES 1. Page, D.N., and Hawking, S.W., Astrophys J.. 206, 1, 1976 2. Hagedorn, R., Nuovo Cim. Suppl., 3, 147, 1965 3. Porter, N.A., and Weekes, T.C., Astrophys J.. 212, 224, 1977 k. Rees, M.J., Nature, 266, 333, 1976 5. Fichtel, C.E., Knfffen, O.A., and Thompson, D.J. Paper A29. Proc. Frascatl Conf.. In Press 1977 ijO;*rcl) 01" UJO Cosmic Garma Radiation Bursts v-ith the Ьдогдзг (jroatcr than 100 Uev on the Date froa the Satellite "Coeaoe-?^' .v.I.ueljacvslv', V.L.Bokov, V.K.Sochurklu, I.F.BujakOV, i"..:...uOi'odiric;c'» c.A.Duitriov, i;.L;.Kru3lov, ....V.Ujaklnin, ^.A.Pjatijoraj;y, i..I.Gliuj;;in, V.S.Jufcrev, J.A.Shibanov, li.II.Koles-'iikova, B.A.beloborodko. A.l'.Ioffe Phycical-Technical Institute, Leningrad, USSR. Theoretical Q Experinencal [x] Both | | Problems on the possibility of existence and discovery or the cos- c-ic gamma radiation bursts of K*^100fclev i s of the quite interest for astrophysics. The upper liwit for number of events and their intensity on the data froa the gamma telescope of the satellite "COSHIOD-5O1" waa estimated. The identical telosco^o \.ith Oeometric factor **-• 100 сшс. ster and effective area 2bO cm2 was set in 1975 on "COSL:OS-731". Data processing of this experiment permits to decrease the upper limit for number of events and estinafce their characteristics more accuretely. Coordinates: ОС 1.4.(tGamma-Kay Burets) blsuling address; A.l?.Iofi'e Physical-Technical Institute Academy of Sciences of the USSR 19402.1; .beningrad, USSR. 188 A SEARCH FOR BURSTS OF 10*3 TO 10 "< «V GAMMA RAYS USING SPACED COSMIC RAY STATIONS* C. O'Sulllvan Physics Department, University Collage Cork, Cork, Irelend. D.J. Fegan, B. McBreen and 0. O'Brien Physics Department i University College Dublin, Dublin, Ireland. Two ground based Cosmic Ray detection systems, separated by 250 km have been operating In absolute time coincidence since January 1975. An analysis is presented based on 423 days of overlap. Data records for events at each station have been examined to look for possible coincident transient Y-ray bursts over a variety of time scales from 2 Us to 10 s. Appropriate upper limits are presented. Introduction The initial objective of this long baseline (250 km) coincidence y-ray experiment was to undertake a ground based search for transient bursts of high energy photons (ilO eV) which might be emitted during supernovac explosions (1,2). Characteristic durations predicted for such bursts are sub- millisecond. Cosmic ray detection systems were built at two stations to exploit such time scales. The range of time scales was subsequently extended to cover the time domain from milli seconds to tens of seconds( as a result of satellite observations of Cosmic Gamma Ray bursts (l,b). More recently, bursts of y-rays have been predicted from the terminal evaporation of primordial black holes (5,6). Two characteristic time scales are Involved, depending on the nucl-ear models used. The elementary particle model predicts the emission of 1030 ergs In y-rays of energy 5xl012eV, of ^Research sponsored In part by the National Science Council. 189 durjclon «bout 0.1 ». The composite particle model predict» the emission of I031* erge In -y-rays of energy 1 0e— 10' eV , over time scales of 10"' ». An experiment*! upper limit has been reported based on this model (7). 2. The cosmic ray detection systems Each station consists of four scintillation counters (each of area I m2) with associated electronics, details of which have been described elsewhere (8,9). The following combinations of detected events are formed at each station by the coincidence log Ic. (a) Coincidences between all four scintillators (b) All combinations of three out of four (c) All combinations of two out of four (d) Summed counting rate of the four Individual counters Event types (a), (b) or (c) are defined as multiples, while type (d) are defined as singles. The multiple rate is typically 3.0 s'1 while the single rate Is typically 780*820 s"1. Nine independent presettable counters Individually monitor the single and multiple event rates over a range of time scales. If the preset count ts exceeded during any sampling interval, the absolute time of occurrence Is recorded to an accuracy of 10~3s. Details of the preset levels and associated sampling Intervals are shown In Table 1. The code 9 events represent an option which is digitally servoed to remove long term (> 15 minutes) fluctuations In the singles rate arising from variations In- atmospheric pressure. The code k events are multiples which are monitored over the same sampling time (10 s) hut are non-servoed. |4U TABLE 1. (Results of 1 typical 2b hour observation) Simp 11ng Preset count Event source Code Polsson Typlc»l t 1 me level type pred1c ted obs« rved 1nterva 1 r«te(diy-' ) r ate (day* ' 1 lOOus 5 Singles 0 21 .7 18.1 1ms 9 • i 1 l 3. Method of analysis and results The detection systems have operated almost continuously for more than two years. -The efficiency of overlap has been better than 80%. The analysis procedure adopted has Involved dividing the data into sub groups of different characteristic time scales which are suggested by possible transient astrophysfca1 phenomena. For each sub group, the individual event times at each station were compared by computer to generate the number of prompt two fold coincidences (Np) , within a given resolving time. Event times for station I were then incremented by 20 s and a control distribution of random coincidences (Nsl) generated. Similarly by subtracting 20 seconds from the event times at station 2, a second distribution of random coincidences (Ns2) was generated. Four sub groups are defined as follows. A1:- All combinations of code types 0,1,2,3 or 7, at both stations, with a resolving time of Is. A2:- All combinations of a code -5 at either station with codes 0,1,2,3,5 or 7, at the second station. A3:- A code 9 at either station in combination with codes 4,5 or 9 at the second station. AA:- A code 4 at either station in combination with codes A or 5 at the other station. The results of the analysis for all four groups are shown In Table 2. 191 FlG.1 A. В. 0 -1 -2 -3 -4-5-6-7-8 910 FIG.2 10 5 A. h 10 5 В. 10 5 C. 0 -1 -2 -3 4 -5 -6 7 -8 -91-0 Fig (1,2) Time distribution of events for sub groups A1 and A2, out to 1 second 192 TABLE 2. ir——— — Sub Character 1111c Co Inc1 dance No. of Nil Nt2 g roup (1 ma teal• raiolvlng days of Np (•20s) (-20i) t Imt ovar1ap Al ilOOms i Is *«23 5* U SC A2 * It ± Is ьи 8 1 1 8 A3 1-HOs ±IOs 31.8 гь 2b 23 A* 1-И0* ±105 306 5 7 6 Fig l(a,b,c) shows the distribution in time of Np. Nsl and Ns2 for the sub group Al. Fig 2(a,b,c) shows a similar distribution for the events of subgroup A2. U. Oiscussfon and conclusions The number of coincident events obtained in the prompt analysis for all four groups Al to A*! is compatible with random expectat ion. With specific reference to the Al prompt analysis, the four shortest time differences observed were as follows. (a) At = .007s (codes 7,1) (b) At = .012s (codes 7,U (c) At = .019s (codes 7.7) lb) At = .033s (codes 1,7) Although no coincidence has been observed with a time difference shorter than these values the system had the capability of responding to transient.» of duration shorter than 2 Ps for multiple events and shorter than 15 Vs for singles. For singles mode operation 50% of detected showers have energies 13 13 between 1.5x10 eV and 9.0x10 eV while for multiples the 13 1 it corresponding figures are 7x10 eV and 2x10 eV. The minimum energy in primary -y-rays (*2x1013 eV) required at the top of the atmosphere to produce a detectable transient (i1 ms) coincidence is of the order of 10"ц ergs cm"2. Examination of the distribution of time differences observed for particular «ode types in sub groups Al and A2, yields the following upper limits at the 99% confidence limits, for the rate of detection of transient events of different characteristic time scales (т). 193 (1) L _ \0m»; Upper limit. 2 per 423 days • 0.86 yr"1 (2) 10 IT- 100ms; Upper limit, 4 per *i 2 3 days • 1.72 yr"1 (3) Ю0 fT?1000ms; Upper limit, 4 per l| 2 3 days • 1.72 yr"> Eleven Vela y-ray bursts hive occurred over the *• 23 deys for which date has been analysed. No two station coincidences have been found on time scales 1 10s, for any of these events. This negative result does not contradict the observations made previously (10) on earlier Vela bursts. We conclude that these upper limits are applicable to all possible sources of transient y-ray bursts In the energy range I 2x1013 eV and over tlmescales from 2 us to 10t. Acknowledgements. We thank Mary Hyland for her valuer1 contribution to analysing the data and Michael Walsh (Computer Laboratory U.C.D.) for wr.ltlng the coincidence search program. References:- 1.). S.A. Colgate. Can. J. Phys. 46, S476 (1968). 2.) S.A. Colgate. Ap. J. 187. 333 C1 97*i) . 3.) R.W. Klebesadel, I.B. Strong and R.A. Olson. Ap. J. Letters 1B2. L85 (1973). 4.) T.L. Cllne and U.0. Desal. Ap. J. Lett. 196, L43 (1975). 5.) S.W. Hawking. Nature 248, 30 (197*). 6.) B.J. Carr. Ap. J. 201, 1 (1975). 7.) N.A. Porter and T.C. Weekes. Ap. J. 212, 224 (1977). 8.) D.J. Fegan, B. HcBreen, C. O'Sulllvan and V. Ruddy. Nucl. Inst, and Meth. 129, 613 (1975). 9.) D.J. Fegan. Nucl. Inst, and Meth. 131, 541 (.975). 10.) S. O'Brien and N.A. Porter. Astrophys. and Space Sci. 42, 73 (1976). 194 SEARCH FOR COSMIC MUON BUKSTS I-ROM RELATIVISTIC DUST-GRAINS H. Muramoto, S. Yamamoto, T. Takahashi, Y. Tera-noto, S. Higashi and S. Ozaki Department of Physics, Osaka City University Sumiyoshi-ku, Osaka, Japan V'e have searched for cosmic muon bursts, high counting rates of muons within short time at sea level to test relativistic dust-grain modffl for the origin of cosmic gamma-ray bursts. Counting rates of muons within 0.1 s have been measured at two stations. Each station consists of two layers of proportional counters. The average counting rate of about 50 counts/0.1 s has been obtained at each station. The distribution of the counting rates has agreed with the Poisson distribution, and no muon burst has been observed since December 1975. 1. Introduction.. Various possibilities for generating cosmic gamma-ray bursts have been suggested. A relativistic dust-grain model has been proposed by Grindlay & Fazio(1974). When a large relativistic iron grain of radius ~1 mm, Lorentz factor of Y~103 entered the solar system,' it would be ionized and broken into droplets at a distance of —100 AU from the sun due to electrostatic forces. The bursts would arise from scattered X-ray fluorescence excited by blueshifted solar photons on the iron atoms. This model could be tested with the observations of cosmic muon bursts at sea level. The iron atoms enter the atmosphere and interact with the nuclei of the atmosphere. Many muons arrive at sea level. We expect that relativistic iron grains originate muon bursts at sea level as well as gamma bursts. The time spread at sea level, of these muons produced by a relativistic iron grain will be~D/c>2 sec, where D is the distance to the grain breakup point. We would expect muoi: burst duration of~0.1 s for grains of Y—103. Each iron atom suffers —10 i'onizing collisions as it approaches the earth and ^-10 photons are emitted per ionization. An iron atom, therefore yields — 100 photons with energy from 10 Kev to 1 Mev. The photon density and burst duration of the gamma-ray bursts which have been observed with Vela satellites is, ~ 20 photons/cm2 and~l s respectively (Klebesadel et al 1973). About 0.1 muons.therefore must be produced by an iron atom in the atmosphere, in order to detect muon counting rate of —100 counts/0.1 s using the detector with area of~6 m2. We have observed muons w^ith energy of— 10 Gev, then the energy of iron atoms of Y~103 is~56 Tev. This is high enough to produce 0.1 muon in the atmosphere. 195 The experiment consist* of two cosmic muon detection systems at sea level, separated 10 m. If relativiatic iron atoms strike the atmosphere, high counting rate of muon», -100 counts/0.1 s should be detected in coincidence at the two stations. 2. Muon Burst Detection System. The detection and electronic system are shown schematically in Fig.l. Two layers of pr.^ortional counters 19 cm apart, each consisting of 13 rectangular proportional counters(height 5 cm, width 10 cm, length 4.5 m) are used to detect muons for each station. 13 anode wires of each layer are connected, and single particleo are detected with a coi ncidence between outputs of the two layers. In order to reject spurious counts, a counter filled with air instead of PR-gas for each layer is used as a veto counter. The output of each layer is coupled to an amplifier, discriminator and monostable. Two fold coincidences of outputs from each layer with a resolving time of 0.5 MS are generated. These coincidences are vetoed by the output of the veto counters and counted by scaler within the gating period 0.1 s for-each station. The coincidence rates of each station are ~500 counts/s. Digital output of each scaler within 0.1 s converted to analog and memorized on pulse height analyzer(PHA). And the measured distribution is compared with the.Poisson distribution for station 1 and 2 respectively. N0.1 STATION fT^ETO COUNTER I I I tffiB E5 JHJ AMP DISC I OR I I I t:/^3 ГЧнС PR. COUNTER S + H-V. > ' A* чие N0.2) | D-ACONV.| PHA RATEMETER J ^— L VETO "i 0-l«c| > N0.2 >- JMAG.COMP] ~©Y TTT JOR \-«|30s«<:[ 1 N0.2 PRESET DATA I DIGITAL N0.2 (PRINTER N0.2 CLOCK Fig. 1. The detection and electronic system. 196 Coincidence rates of each station are measured with ratemeters with tine consta.it of 0.1 s and recorded bv multi- pen recorder with'high frequency. ',м-' ng rate* of veto signals and the coincidence are also .; jrJed with time constant 0.1 s and 30 s respectively to cncck the detection system. Each of the scalers has a pre-determined count level. Mien both scalers exceed the levels within the gating period 0.1 s in coincidence at the two stations, the time of the event occured and the counting rates of each station are recorded on a digital printer. Each scaler is reset to zero at the end of the period and any time by veto signals. 3. Results. The observation started with one station on December, 1975. Coincidence rates are measured only with ratemeters of time constant 0.1 s, and recorded by the pen recorder. One event of muon burst has been observed on February 2, 1976. The counting rates observed are 132, 142, 107, 99, 112 and 130 counts/0.1 s on 3h26m, 3h35m, 3h40m, 3h43m, 3h53m and 3h56m U.T. respectively. The counting rate of veto counter was almost zero, the coincidence rate within 30 s also normal and the average rate was 52.7 counts/0.1 s. The distribution of counting rates within 0.1 s observed has agreed with the Poisson distribution except tlis event. But this event was observed with just one station, so it is difficult to reject backgrounds, maybe electric noises perfectly. The one more station was constructed and the observation started with two stations on July 1, 1976. The coincidence rates of each station were measured with ratemeters and recorded by the pen recorder. The digital printer was not used. No large burst as described above has been observed. «Is- I MEAN-5029» \ MEAN > 46.996* COUNTS (111 \ COUNTS / ul» > STATION 1 Hf' / N(COUNTS;OLU> 20 30 «0 50 60 70 SO 90 10 20 30 40 50 60 70 80 90 Fig. ">. Differential nrnhahi1ity distribution vs observed counts, N within 0.1 s for station 1 and 2, where the broken lines are the Poisson distributions. 197 On March 9, 1977, the observation has started with the full system described in the section 2. \ч The differential probability \ distribution of the counting rate within 0.1 s of each station is measured by the PHA, \ where the probability is the ratio of frequency of the \ counting rate to the total number of sampling. And the result from March 9 to April 11 is shown in Fig.2, and compared with the Poisson distribution of average counting rate 50.3, 47.0 counts/0.1 s for station 1 EXPECTED FROM 4 and 2 respectively. THE POISSON DISTRIBUTION,** The integral probability distribution of both counting rates within 0.1 s being greater than or equal to N1,2 in 67 68 69 TO 71 11 .71 74 75 coincidence at the two stations Nt2( COUNTS #0.1») is obtained from the data obtained by digital printer from Fig. 3. Integral probability March 9 to April 11 and is shown distribution of counting in Fig.3. The distribution rates of both stations bs'ng observed is compared with that greater than or' equal to predicted by the differential Ni<2 in coincidence at two probability distribution of each stations. station (Fig. 2). This result shows a little increase of the probability compared with that expected, above ~73 counts/0.1 s. We will continue the observation, to reduce the statistical errors. References. Grindlay, J., and Fazio, G. 1974, Ap. J. (Letters), 187, L93. Klebesadel, R. W., Strdng, I. B., and Olson, R. A. 1973, Ap. J. (Letters), 182, L85. 19» GAMMA RAY PROOUCTION IN INTERSTELLAR SPACE G. D. Badhwar and S. A. Stephens» NASA Johnson Space Center Houston, Texas 77058 USA We have obtained a simple representation to the observed Invariant cross-section for the production of «* mesons 1n pp collisions. Making use of this representation, we have calculated the differential and Integral production spectra of gamma rays In the galactic space from the Interactions of cosHc ray nuclei with Interstellar gas. These spectra are compared with those from the existing calculations. 1. Introduction. The high energy gamma ray measurements made so far have provided valuable Information towards the understanding of the structure and contents of the galaxy. These observations were aimed at obtaining the bright ness distribution of the gamma ray sky without spectral Information. It- 1s generally believed that most of the observed gamma rays above 100 HeV are due to the decay of neutral mesons, created during the collisions of cosmic ray nuclei with the ambient Interstellar gas, and that the contribution from Inverse Compton and bremsstrahlung processes is small. Since the competition of these production processes differ from region to region in the galaxy, depending upon the physical conditions existing therein, it Is possible that the energy spectrum of gamma rays may not be the same everywhere. The COSB satellite results have already shown the capability of determining the , spectral shape of gamma rays 1n the energy region of 20 MeV to 2 GeV. It Is thus desirable at this stage to determine the expected gamma ray energy spectrum. Of the three modes of gamma ray production, the contribution from the bremsstrahlung and Inverse Compton processes can be calculated precisely for given physical conditions 1n the galaxy. On the other hand, the determination of the gamma ray spectrum through we decay requires a knowledge of high energy interaction characteristics also. Previous estimates of the ir° production 1n nuclear collisions were model dependent because of the lack of sufficient accelerator data at energies of interest. However, with the avail ability of extensive accelerator data at present on inclusive reactions, we have 1n this paper calculated the production spectrum of gamma rays due to the interaction of cosmic ray nuclei 1n Interstellar space, and compared it with other existing calculations. 2. Representation of n" Invariant Cross-section and the Gamma Ray Production Spectrum. The differential cross-section for the production of »* meson of energy Е by a proton of energy Ep In pp collision is given by do(E,Ep)/dE -Zt/P, (E^£) de (1) where (E d3 3 q (E dVd p) • UOiiriys (l-x) expC-BP^O.O+tapVs)] (2) | describes the invariant cross-section well. In this equation | *NASA/NRC Senior Research Associate now at Tata Institute of Fundamental Research '% 199 + + х * / The production spectrum of n0 mesons due to the interaction of cosmic ray proton with one hydrogen atom of interstellar space can be written as P. (E)dE = / j (ED)dEn {da (E, Е J/dE}dE_ (3) и £ V p p p where j(Ep) is the proton differential spectrum. In Figure 2, we show the proton spectrum observed in the neighborhood of the Earth during the period of minimum solar modulation. The data points are taken from Ryan et al. (1972), Smith et al. (1973) and from the compilation of Simpson (1971). The solid line is a general fit to the observations and can be approximated by three power law spectra in total energy and are given, as: 2.0 x 10ц En"2-" for 3 2 5 Ep > 70 GeV (Ryan et al.', 1972), 6.92 x 10 Ер" - for 3 GeV < Ep < 70 GeV and 3 1 68 2 2.81 x 10 En" - for Ep < 3 GeV (the flux being expressed In particles (m sr sec GeV)"1. The dotted line in Figure 2 is the demodulated proton spectrum using the demodulation parameters derived by Burger (1971), which explains satisfactorily the time variation and absolute modulation of electron and nucleonic components. Since the cross-section for the production of pions decreases rapidly below 2 GeV (Figure 1) the contribution of protons below this energy is less than 105S of the total gamma ray (*. 100 MeV) flux. Thus a factor of two uncertainty in the demodulated proton spectrum has a negligible effect on the total gamma ray flux. We have calculated the production spectrum of it" mesons by evaluating Equation (3) using Equations (1) and (2), and by making use of the proton spectrum as observed in the neighborhood of the earth (solid curve in Figure 2) and the proton spectrum in the interstellar space (dotted curve in Figure 2). In the case of interstellar space, we have assumed that the energy spectra of nuclei of charge >_ 2 are the same as that of protons. Therefore, we :оо have multiplied the »* spectrum derived . ""™ Ujr using the interstellar proton srfec- • < trum by a factor of 1.64 to take Into ••• ^-^^ ! account the cosmic ray composition and s the composition of Interstellar gas. ! (This factor Is consistent with that s< • derived by Stecker, 1971.) The gaunt / ray spectrum 1s then determined as Pj(E)dE 4 * * .».*Л ^-J-. / (2dE//(Ej„ mJ.))P¥.(E-)dE' (•> Е" * F1g. 1 - Total «• cross-section. The Here Е- « Е + m2/4E. solid line Is our calculation. 3. Results. The calculated 10* -I • i ii nini i i 1111MI—i i im differential spectra of gamma rays through ne decay are shown in Figure 3 1n units of photons/(sec sr GeV) per 2 unit hydrogen atom. cm" . Curve A In JO» this figure is the calculated gamma ray production spectrum for p-p OBSERVED collisions alone using the observed INTERSTELLAR proton spectrum near the Earth. This spectrum peaks at 67.5 HeV, which 1«и corresponds to half the mass of »° and the spectral shape is nearly > symmetric about this energy. At high energies, the spectrum gradually I .0< steepens and becomes above 20 GeV, a power law with the same spectral Index as that of protons. For comparison is shown in Figure 3 (Curve Al), the spectrum calculated by Stecker (1970). Й 10" Even though there is considerable agreement between the two spectral shapes, one notices that the spectrum of Stecker is not smooth and as a result, the spectral shape is different from our calculation (Curve A) in some • SMITH «tol (197$) ° RYAN etui (1972) energy regions. It can also be seen - •( COMPILATION OF that the absolute value of the flux SIMPSON (1171) differs by about a factor of two at iff* many energies even though the Input proton spectrum below about 70 GeV 1s exactly the same as that used by Stecker (1971). Moreover, at energies i i l inn i io4 ' Mil III £ 1 GeV, the spectrum calculated by 1 2 Stecker using an input proton spec 10 10 10 10' trum index of n,2.6 has a spectral ENERGY IN G«V index of about -3.0. This Is In Fig. 2 - The solid line with the data disagreement with our calculations, points 1s the observed proton enercjy which indicate that the qamma rav spectrum and the dotted line is the demodulate'* inpctrum. 201 spcctrun slowly steepens to an asymptotic power law spectrum with The calculated Interstellar production spectrum 1s shown by Curve В in Figure 3 and its spectrum above 20 GeV can be represented by a simple power law of Fig. 3 - The calculated production spec the type tra of gamma rays. Curve A are the differential spectra from observed cosmic P (£) - 2.52 x 10"27 E"2,75 ray protons with hydrogen, Curve В are Y those from cosmic ray nuclei with inter stellar gas and Curve C are the integral photon/ Burger, J. J.f Ар. J.. )66. 65) 0971). Cavello, G. «nd~GouTd, IT"J.. Nuovo Clmento. 132. 77 (1971). Carey, tt al.. Phys. Rev. Letts.. 33. 327 (19ДТ- D1dd#ns. A. N. and SchlupMnn, K., published In Landolt. Bornsteln (1971). Gangull. S. N. and Sreekantan, B. V., Proc. 14th Int. СомИс Ray Conf.. Munich, 2. 570 (1975). Ryan, M. J.,T)s*es, J. F. and Balasubrahmanyan, V. K., Phys. Rev. Letts.. 28. 28. 965 (1972). ~~ Simpson. J. A., Proc. 12th Int. Cowlc Ray Conf.. Hobart, Invited tnd Rapputeur ptptr, p. 324 (1971). Smith, L. N.. Buffington, A., Smoot. G. F. and Alvarez. L. M.. Ap. J.. ISO. 987 (1973). — sucker. F. W., As trophy s. and Space Set.. 6, 377 (1970). Whltmore. J., Phys. Reports (Section C of Phys. Letts.). 10. 273 (1974). 203 n-*v~— lawtlgatlo» at •paatta aat onapnaltlna of KLaaly патцаЛ tiff"*—>fa of to» «oaala raftlatloa alto • —#Mt1« mwUoMHct Г.Л^ЯпсдаоХру, Я|Л.ТиПг>ау» V*A«toama »«i«lattc ttyaloal-a.-ooaoUal Xmtlcuto» bwilngrart, UUUH. AaoaaOaal D UxgmrimmiMl [x] ПОЛО S3» reoulto of bnUoett fll#ito of a aagpatlo «peotrocoter •) oonalwtAng at а ДававМав аадааа» а apaak-atoabaa-oyatoa and GegmVtrm antiifttllariwn oasato* tobw>qp« ага рмит>< Zaa davloa ама аам asad to ааимиа taa aottpcoton - proton tin tati» la «fa» palawqr «oaate wUrttKm attain «to 2 - 5 Gar raaga» «to flax at dawtaaw «Ita «шщкзг 1 Qar/nael and ap—tea o* «a» aaonndaay todUtton alto aauasr abora 0*1 Oar. Яи aaawcaaasto мм austed «at at alddla latltoda* la 1902 - 1976. *) "A «aaraa for aat&piotoaa la toe prtaeiy ooaoio radiation wlto a aagoauio Aparcrowtoju, iroc* 12th ICCR, лзЬлг*, Zsaaaola» AavteaUa, Majw» 1971» 00-42» • (toijia mmi Miifirt if niwatij miju) Oooadlamtoat oa-i.8 м«1<14»д afldiaaat л«Г*1ава laTataUHEagastaal Xastatato 19«pi» laalasrad, mn» 3Q№C0Z59 TU= PROTON AND HFXIIN RIGIDITY SPECTRA FROM 10 TO SO C\ G. D. Badhwar, R. R. Daniel», Г. Cleghorn, R. L. Golden, J. L. Ucy, S. A. Stephens*, J. E. Zipse NASA Johnson Space Center Houston, Texas 770S8 USA A Magnet spectrometer flown fro» Palestine, Texas (September IS, 1976) has produced measurements on 1.5x10s protons and 2x10** heliun nuclei covering the rigidity range 10-50 GV. In this interval, the proton spectrum has been fitted to a power law with a spectral index of 2.51 ± 0.03. Similarly, the heliun nuclei rigidity spectrum can be repre sented as a power law with an index of 2.61 ± 0.03. 1. Introduction. Proton and heliun nuclei are the two dominant cosmic ray components'! It-is therefore of basic importance to study their spectral shape and relative abundance. Though numerous experiments have been performed in the past to study these components, there are, perhaps, only five1"5 in which direct measurements significantly greater than 10 GV have been made. In these experi ments, Anand et al.1 used the variation of the geomagnetic cut-off with zenith angle, Webber et al.2 employed a gas Cerenkov counter and the geomagnetic cut- - off, Verma et al.3 used a permanent magnet with an emulsion stack, Ryan et al.1* used an ionization calorimeter and Smith et al.5 used a superconducting magnet spectrometer. There is, however, considerable disagreement among the spectral indices determined by the various experiments. In the two experiments of Ryan et al.1* and Smith et al.5 with the highest statistical accuracy, the spectral indices of helium nuclei above V10 GV are respectively 2.77±0.05 and 2.47+0.03. The measurements of Verma et al.3 and Anand et al.' are in agree ment with a steeper index of 2.7 - 2.8. As for the proton component, it seems from the measurements so far made1* that in the rigidity range of 10-50 GV the spectrum is somewhat flatter than that of 2.75 now generally accepted for rigidities between 50 and 1000 GV1*. In the present paper, we describe an experiment with a „superconducting magnet spectrometer to determine the spectral shapes of protons and helium nuclei in the rigidity interval of 10-50 GV. 2. Experimental Details. The magnet spectrometer which is described in detail elsewhere" is shown in Figure 1. It consists of: (i) a gas Cerenkov counter, G, having a threshold, Yc40; (ii) scintillators SI and S2 each of 0.625 cm thick Pilot Y; (iii) a stack of five multiwire proportional chambers, MWPC, MI- MS (chamber M2 did not function throughout the flight) with a spatial resolution in the cathode coordinate of % 200pm; and (iv) scintillators P1-P7 each of 0.625 cm thick Pilot Y and each separated by %1.2 radiation lengths of lead forming a shallow shower counter of £7 radiation length. The signals from SI, S2, P1-P7 and G are all pulse height analyzed. The magnet was operated at a current of 120 amps producing a magnetic field of £ 40 KGauss at the center of the coil. Only events satisfying the trigger SI PI P7 were accepted.for analysis. This mode has a useful geometry factor of £ 275 cm2 sr. The instrument was flown from Palestine, Texas (cut-off 4.5 GV) on September 16, 1975, under 5.2 g cm"2 of residual atmosphere for a total exposure factor of 1166 m2 sec sr. «NASA-NRC Sr. Postdoctoral Resident Research Associate on leave from Tata Institute of Fundamental Research, Bombay, India. 205 Events were selected by requiring that (a) they represent single particle traversal through all the four MKPC, Cb) the particles have a downward a direction of notion as determined by the ti*e of flight between SI and PI, (c) the charge determined by SI, S2 and PI is consistent with Z»l or Z-2, and (d) the particles do not interact in the shallow shower counter. Using these selection criteria, we have a clean sample of 1.5x10s particles of Z-l and 2.Ш.01* helium nuclei. The overall efficiency for Z-l particles is 5St and for Z-2 is 23%. The №PC chamber alignment is made using a multiparameter minimizing routine to obtain best straight lines to tracks of a large number of sea- level muons which trigger the gas Cerenkov counter when the magnet is off. We define the magnetic deflec tion, D, as the inverse rigidity of the particle. Since D is proportional to the spatial deflection, its error distribution is related to the position measurement error distribution. (Con tributions to the error distribution due to mltiple Coulomb scattering are negligible at rigidities of concern here). The spatial error function of magnet-off sea-level nuons is shown in Figure 2. Ideally, one would like to obtain this function in flight with the magnet off. However, due to a premature balloon cut-off, this was not possible. In order to check that the ground level error function is applicable to the flight data, we show in Figure 2(b) the observed deflection of G-on helium nuclei, whose rigidities from their pulse height in the G-counter correspond to rigidities greater than 100 GV. Though the statistics of G-on helium nuclei is limited, one can, by super imposing the straight track distribu tion of 2(a), see that the agreement between the two distribution is very good. We have, therefore, assumed that the observed sea-level error function applies to the flight data. A Gaussian fit to the central section of the error distribution (D-± 0.05), -0.05 0 DEFLECTION, D, (OV1) :об gives an approximate maxuiun detectable momcntun of *70 GV. 3. Results 5.1 The "Proton Component. The deflection distribution for >1 events is shown in Figure 3. It is well known that the spectrun in interstellar space is significantly modified by the effects of solar modulation and geo magnetic effects only at rigidities £ 8 GV. Thus, if we restrict our interest to R K 8 GV we should be sampling the interstellar spectrum except for some very small atmospher ically produced meson background (S It). We have fitted the deflection distribution in the range 8.3 GV to 50 GV (deflection interval of 0.12 to «.u и m .02 GV'l) to a power law, J(R)«AR-Y, DVUCnON. O. To convert the observed deflection spectrum into an absolute flux, we 0.0 0.01 I .19 need to correct for all of our ИНКТЮИ, D. {OV-1) selection criteria and dead time. The overall efficiency is 0.35, with Fig. 4 - Observed deflection an uncertainty of ±154. The result distribution of Z=2 particles. The ing absolute differential rigidity solid line is power law in rigidity spectrum is J(R)=(S500±865)R-2-51io-03 R-2.6 1 convoluted with the error protons/m2sec sr GV at the balloon function of Figure 2(a). altitude. The absolute spectrum is 207 in good agreement with that of Stoith et •].' it R i 8 CV. 3.2 Helium Nuclei. The deflection distribution of Z*2 events is shown in Figure 4 ml as см be seen, it has « shape rather similar to that of 2-1 events. To avoid modulation effects, we have fitted in the interval 10 - SO CV a power law in rigidity in a manner similar to that for Z-l events. The result ing spectral index is 2.61 ± 0.03 with a confidence '.evel of 0.45. Table 1 gives a comparison of various experimental results. Table 1 Anand et al. Verma et al. Ryan et al. Smith et al. Present Y 2.74 ± 0.1S 2.8 ± 0.15 2.77 ± 0.05 2.47 t 0.03 2.61 ± 0.03 Rigidity Range in 12 -.40 22 - 150 20 - 800 8.3 - 100 10 - 50 GV The exposure factor for helium nuclei corrected for all selection critieria is 263 ± 40 m2sec sr. This gives an absolute rigidity spectrum J(R) « (1831±274) R-2.6i±0.03 above R > io GV at the balloon altitude. 5.3 The Ratio a/p. The difference in the proton and helium spectra can be more easily displayed by plotting the a/p ratio as a function of the deflection. This ratio uncorrected for different proton and alpha exposure factors is shown in Figure 5. We note that for R>10 GV, there is clear evidence for a rise in the ratio indicating that the proton spectrum is flatter than the helium spectrum. Because of different modulation effects at the same rigidity for helium and protons, ll the ratio below 10 GV cannot be used for this analysis. We note that displaying this effect in terms of a I" ratio has the virtue of removing all 6 4.0 instrumental biases which are common to both protons and helium measurements. 4. Conclusion. We have determined 1.4 - the rigidity spectra of protons and helium nuclei to have spectral indices U - of 2.51 ± 0.03 and 2.61 + 0.03 in the и' 1 1 1 J range 10-50 GV. Using cuts on the 0.0 м .10 .u a» / 1} x dl_ we hope to extend these BBUCIION, D, (OV)'1 measurements to 150 GV in the near Fig. 5 - The ratio of differential future. helium flux to the differential proton flux at the same deflection. References. T. K. C. Anand, R. R. Daniel, S. A. Stephens, B. Bhowmik, :о8 С. S. Krishana, P. К. Aditya and Г. K. Puri, Canadian .). of PhyMi:* -If-. S6S2, 196S. 2. T. T. Von Rosenvingc, N. R. Kcbber, and J. F. Oraes, Astrophyv anJ \-м. Sci. S, 342, 1969 and 3, 4, 11,69. 3. R. P.~Verma, T. N. Rengarajan, S. N. Tandon, S. V. ftimle ami Yash Pal. Nature 240. 13S, 1972. 4. M. J. Ryan, J. F. Omies and V. K. Bnlasubrahmanyan, Phys. Rev. Letts. J£. 985, 1972. 5. L. H. Swith, A. Buffington, G. F. Smoot, L. K. Alvares and W. A. KahliR. Astrophysical J. 180, 987, 1973. 6. R. 1. Golden, G. ТГГBadhwar, J. Lacy ajkl J. E. Z.ipse, Nucl. Inst, and Methods, 1977. \ / .l,.i, „ ^JJi THE COSMIC RAY ANTIPROTON FLUX: AN UPPER LIMIT AT THE PREDICTED LEVEL OF OBSERVATION G. D. Badhwar, R. R. Daniel*. T. Cleghorn, R. L. Golden. J. L. Lacy, S. A. Stephens*. J. E. Zip» NASA Johnson Space Center Houston, Texas 770S8 USA Data gathered fro* the September 26, 1976, balloon flight of the Johnson Space Center superconducting magnet spectro meter has been examined for the presence of cosmic ray antlprotons. The ratio of antlprotons to protons, p/p. 1n cosmic rays was found to be (0.0 ± 3.3) x Ю"* In the rigidity Interval 4.2 to 12.S GeV. The 67X confidence level upper limit for p/p Is thus 3.3 x 10"". This upper Hm1t Is In strong contradiction with the prediction of the closed galaxytnodel of Rasmussen and Peters, but is not Inconsistent with the prediction of the modified closed galaxy model of Peters and Kestergaard. It 1s nearly equal to the predictions of conventional propagation models. 1. Introduction. Antlprotons (p) are produced by collisions of cosmic ray nuclei with Interstellar gas. The production cross-sections are known up to 1500 GeV1»2. The energy of relatlvlstlc p does not change significantly during propagation. Thus, В measurements offer a very direct way of determining the amount of material traversed by cosmic rays 1n the galaxy. The determination of the p/p ratio has the potential of leading to knowledge of even greater significance. For example, If the p flux were far below the expected flux, support would be given to the speculation that the p has a finite lifetime3. Cosmic ray p have not yet been observed. The previous most stringent upper limit for p/p at relatlvlstlc rigidities 1s 6xl0-3 in the region 3-6 GV1». The experiment reported here is the first 1n a series of flights to study positron and p fluxes. This -flight was configured for optimum positron (e ) performance by virtue of the gas (air at one atmosphere) chosen for the gas Cerenkov detector. Because^ of tne lower refractive Index gas used for positron experi ments, the detector_was not able to completely remove the atmospheric meson background for the p experiment. Nevertheless, we have obtained an upper limit a factor of 20 below that of previous experiments and near the expected level for standard propagation models. 2. Method. A complete description of the apparatus 1s given by Golden et al.5. We present here only a brief discussion of the detector system and Us perform ance. The detector elements of the balloon-borne magnet spectrometer are shown in Figure 1. The elements are: (1) a multlwlre proportional counter stack (MI- MS) which-provides rigidity determination6; (2) a gas Cerenkov detector (G) using air as the Cerenkov medium; (3) scintillators (SI, S2, PI) to determine particle change; and (4) a seven radiation length lead scintillator shower counter (P1-P7).. Time of flight Is measured using SI and PI. The apparatus was flown on September 26, 1975, from Palestine, Texas (vertical geomagnetic cut-off of 4.5 GV) under 5.2 gm/cm2 of atmosphere for 12 hours. The geomagnetic factor for the S]Pi trigger is 320 cm2sr. •NASA-NRC Sr. Postdoctoral Resident Research Associate on leave from Tata Institute of Fundamental Research, Bombay, India. :ю Antiproton candidates were selected by requiring (1) a charge measurement \ i (from SI, S2, PI), consistent with proton charge, (2) time of flight corresponding to downward motion, (3) no cascade in the shower counter (This criteria is essential 1n eliminating electrons), (4) curvature in the proportional counter stack correspond ing to a negative charge, (5) the trajectory crosses sufficient magnetic field (>. 3.0 Kg-m) for the spectro meter to have a maximum detectable momentum7 of about ISO GV, and (6) no G pulse. Events satisfying these criteria are: (a) negative atmospheric secondaries (muons, pions, antiprotons, etc.) which did not trigger G, and (b) positive protons with very high rigidity which appear to have nega tive curvature due to the chamber measurement error or due to scattering (these are called spill-over protons), and (c) galactic antiprotons. The search for antiprotons was performed by computing the deflection Fig. 1 - Magnet Spectrometer distribution of the expected atmos pheric secondaries and spill-over protons. The computed distribution was then compared with the observed one. Any observed excess could then be attributed to galactic antiprotons. Since the number of p produced in the galactic medium and in the atmosphere are both proportional to the number of primary protons, we express the observation as the ratio of antiprotons to protons. This has the additional advantage of eliminating the uncertainties Introduced by separately measuring the absolute p and p fluxes. The geometry factor and detector efficiencies are the same for p and p. 3. Determination of the Expected Background. Determination of the expected background requires a detailed understanding of the production of background in the upper atmosphere and a thorough analysis of the experiment's perform ance characteristics. The critical parameters of the experiment performance are the efficiency of the G-counter as a function of rigidity, and the reso lution of the magnet-spectrometer system. In this section, We discuss briefly each of these factors, and the resultant background estimate. Due to the development of accurate analytic representations of the plon and kaon production cross-sections8, it 1s possible to calculate the atmospheric mesons fluxes with an accuracy of about 10X. Me have applied the method of Badhwar et al.8 to compute the negative meson fluxes at 6 gm/cm2. Figure 2B contains the computed fluxes normalized to the observed proton flux. Rejection of the meson background 1s achieved by the requirement of no G signal (selection criteria 6 above). The G threshold with air as the refractive medium 1s 4.24 GV for muons. The G-effic1ency varies from zero just below"geo magnetic cut-off to a maximum value at ^ 10 GV. The exact form of the 211 efficiency factor Is crucial to the estimation of the muon background. ThH efficiency function was determined for 60.000 ground-level muons by measuring the fraction of G triggers as a function of the measured muon magnetic deflec tion. The fraction of muons not triggering G Is shown as a function of magnetic deflection 1n Figure 2A. Tests of th: Instrument at the Stanford Linear Accelerator have shown the reso lution of the Instrument to be essentially Independent of rigidity above 4 GV. Thus it Is possible to construct the deflection distribution of "spill-over" protons by folding the resolution function with the appropriate proton deflec tion distribution. The resolution function was determined by operating the experiment at ground level with the magnet off. This simulates a source of muons of Infinite rigidity. The resolution function obtained with ground-level muons Is estimated to be 20* broader than the true resolution function due to the Coulomb scattering of the low energy muons. We have computed the spill-over curves using the observed sea-level muon resolution function, for all protons (see Figure 2B) and for G-off protons (see Figure 2C). In both cases, the spill-over is less than Ti of the muon flux 1n the 4-12.5 GV region. 4. Results. Figure 2B shows the deflection distribution of the antlproton candidates prior to the requirement that there be no G pulse. Note that the calculated atmospheric background agrees very well with the observed data for deflection where the spill-over is not significant (the reader is reminded that the normalization was determined only by the atmospheric depth and the number of positive protons in the data sample). Combining the spill-over curve with the computed background we are unable to observe any excess of observed particles above geomagnetic cut-off. If a sufficient number of ant1protons were present, we would expect an excess of observed events beginning at geomagnetic cut-off and extending toward zero deflection. Since the G-counter threshold for muons is 4.24 GeV/c, accepting only events which have no G-pulse should suppress the muon flux above 4.24 GeV/c and greatly enhance any excess due to p. The expected deflection distribution for "no-G" events is obtained by combining the G-eff1c1ency curve of Figure 2A with the mesons background curve of Figure 2B. The resultant curve, together with the observed data is shown in Figure 2C. Again, combining the spill-over curve with the meson background curve, we find good agreement with the observed flux. No excess Is observed 1n the region above geomagnetic cut-off. There are 66 events in the Interval .08 - -24 Gy-i (4.2 - 12.5 GV) on Figure 2C. The back ground estimate is 66.5 ± /(6.6)* + (8.1)* where 6.6 is the error associated with knowledge of the cross-sections and'8.1 1s the expected statistical error. There were 32,300 proton events In the data sample. Thus, the antlproton/proton ratio is: £= (0.0 ± 3.3) x Ю"* Regarding this as an upper limit, we find £ < 3.3 x Ю"* at 67* confidence and £ < 6.6 x 10-1* at 95* confidence. The production of antiprotons in the upper atmosphere is small relative to the galactic production. Two factors significantly limit the number of atmospheric p. First, the ratio of the production mean free path Tn air to that In hydrogen is 0.65. Secondly, antineutrons produced in the atmosphere do not have an 2i: opportunity to decay. Thus, tha atmospheric p flux In 6 gem"2 of atmosphere should be equivalent to that produced by -v 1.5 gm/cnt2 of Interstellar material. The atmospheric antlproton background was neglected 1n the above calculations. 5. Conclusions. Badhwar et al.1 have computed an expected value of | • «xlO"1* for both diffusive propagation models and leaky box models assum1ng:p (1) a mean matter traversal of 5 gm/cm2 at 5 GeV; (2) the density of target nuclei Is Independent of energy; (3) the artiproton 1s stable on a time scale of >20 x 106 years; (4) there are no primary cosmic ray antlprotons; (5) there is no significant amount of material in the accelerating regions. The present upper limit is therefore consistent with the assumptions made In these calcula tions though tantHirlngly close to it. We have computed the expected p/p ratio assuming the closed galaxy model of Rasmussen and Peters9 for cosmic ray propagation. With the assumptions (2) - IS) stated above, we find a predicted ratio p/p * (30 ± 10) x 10"*. Badhwar et al.10 have pointed out that the Rasmussen-Peters model is very difficult to reconcile with existing electron and positron observations. The upper limit presented here is even more inconsistent with the Rasmussen-Peters model. The Peters-Westergaard model11 is a closed-galaxy model modified by the presence of a nearby source. Steigman3 has calculated a p/p ratio of (6 ± 2) x lO"1* for this model. This value is about 2 standard deviations above the result of this model. A second flight of the magnet spectrometer with eight proportional counters was performed in June 1976. Analysis of this flight is in progress. A flight is planned for the near future with improved gas counter efficiency. The G- counter will be operated with sulfur hexafluouride, lowering the muon threshold to 2 GV (substantially below geomagnetic cut-off). The increased index of refraction greatly increases the light level and consequently gives a much higher G-efficiency. Recent testing with SFg in the G counter gave an efficiency of 99.4% in the 4.2 - 12.5 GV range. This change will reduce the muon background to a negligible level and should permit an acual measurement of the p flux. References. T. G. D. Badhwar, R. L. Golden, M. L. Brown and J. L. Lacy, Astrophys. and Spa. Sci., 37, 283 (1975). 2. T. K. Gaisser and E. H. Levy, Phys. Rev. D10, 1371 (1976). 3. G. Steigman, Phys. Rev. Letts., to be published (1977). 4. E. A. Bogomolov, N. D. Lubyanaya and V. A. Romanov, Proc. 12th International Conference on Cosmic Rays, 5_, 1730 (1971). 5. R. I. Golden, G. D. Badnwar, J. L. Lacy and J. E. Zipse, Nucl. Inst, and Methods, to be published (1977). 6. The superconducting magnet produces a magnetic field of 10-30K Gauss which is primarily horizontal 1n the chamber stack. 7. In this paper, we refer both to a particle's rigidity (momentum per unit charge measured in GV/c) and Its magnetic deflection (which we define at l/rig1d1ty). The definition Is useful because spatial deflection of a particle 1s inversely proportional to rigidity. 8. G. D. Badhwar, S. A. Stephens and R. L. Golden, Phys. Rev. D15, 820 (1977). 9. I. L. Rasmussen and B. Peters. Nature, 258, 412 (1975). 10. G. D. Badhwar and S. A. StephensT"Fnys.~Eev. 14., 356 (1976). 11. B. Peters and N. J. Westergaard, Astrophys. and Spa. Sci. (1977). 213 a 5 1.0 FIG. 2A о .8 FRACTION OF MUONS S NOT DETECTED BY G .6 IO N 5 D Z .4 3 .2 ON S u. 3 0 J ' • ' I L _L J I L S COMBINED OBSERVED BACKGROUND- FIG. 2B DATA ANTIPROTON 12 50 CANDIDATES (NO G TEST) £ 40 Ш u. 2300 i++i ATMOSPHERIC BACKGROUND- Zin | 10 SPILL-OVER 1 o I I I I I I -L GEOMAG CUTOFF 12.5 GV 1 50 FIG. 2C 2 40 ANTIPROTON ш CANDIDATES & 30 (GOFF ONLY) O к 20 ш GOFF ATMOS SPILL-OVER | 10 BACKGROUND— i 0 J 1 I I I I I—L -.44 -.36 -.28 -.20 -.12 -.04 -.40 -.32 -.24 -.16 -.08 DEFLECTION (GV/c)"1 214 Estimation of Primary Proton Spectrua bitimi 1012 and 1014«V T. Г. Caletar, lartol •.•••arch Foundation of th» rrenklla laatltute Svartfcaore, Poonaylvanla 10981 U.S.A. ' F. Slohan and C. 1. Yodh, Coda 661, KASA/Coddard Spaea Plight Canter, Creenaelt, Maryland 20771 U.S.A. Th. optical Q Bxpenfrantal (3 Both Q Abatract Th beh vlor of ,ллл * " proton-proton toul eroaa aactlon batman 2000 and 3000 GaV haa bean deduced uelng dlaperalon ralatlona and recently measured real part of tha forvard scattering amplitudes at CERH-ISR. Va derive, using Glauber theory, the proton-*lr In.la.tic eroia section and calculate fro. tha measured unaccompanied charged badron flux at «ounceIn altltudea tha primary proton apactrua above 2000 GaV. The resulting spectrum, la 1A good agreement 4lth an extrapolation of the spectra below 2000 GeV with no evidence of a band around 2000 GaV. Coordinates: OG 1.5 Nuclear Composition of Coamlc Raya Mailingaddren: Dr. G< B< Yodh Code 661 NASA/Goddard Space Flight Centex Greenbelt, Maryland 20771 U.S.A. :i5 Classification: Title: On the Proton Spectrin Measuraaenta Above 2000 C«V Author(s): C. B. Yodh. Ж. W. Ellaworth, A. Ito, J. MacFall, F. Slohan, R. E. Streltnatter, 5. C. Tonwar, V. X. Balaaubrahaanyan Sponaorlng Inatltutlon(a): NASA/Goddard Space Flight Canter and U. of Maryland Greenbalt, MD 20771 Postal Address for author underlined above: NASA/Goddard Space Flight Center Code 661 Greenbelt, MD 20771 Abstract The steepening of the proton spectrum beyond 1000 GeV and the the large rise In Inelastic cross sections between 20 and 600 GeV observed by rroton-1-2-3 satellite experiments are shown to be due to systematic effects of energy dependent albedo and energy sampling of the calorimeter. The proton spectrum corrected for these effects shows no steepening and can be represented by a power law spectrum with constant slope. :i6 THE IRON/PROTON PRD4ARY RATIO AHD RELATIVE СОвШС RAY COMPOSITION AT 10I,-1013 eV J. K. OrtndUy Cantor tor Aatrophyatea, Cambridge, Massachusetts IX Hartmaa and T. C. Week** Smithsonian Institution, lit Hopkins Observatory, Tucson, Arisona Theoretical • Experimental Q **» Щ We report the reeulU of simultaneous measurement» of the Cereokov pulae height apectnmi from the electron and muon componenta of the earn» extensive atr ahowera (EAS). Both the number of EAS O 20 000) and dynamic range of the pulae height record Ing represent a factor of >10 tnoreaae over a preliminary version of tela experiment reported by ua at the Denver conference. The total Cerenkov light pulae height recorded by two narrow-field reflector* at 188 m aeparatlon and pointing at the expected altitude of EAS maximum waa digitised. The Cereokov UV component (<4 000 A) of the aame EAS waa recorded by a third reflector near the baseline center at angles >1* from the mean EAS core direction. The ratio of the total Cerenkov radiation In these two components ts directly related to the EAS electron/muon ratio, respectively, and thus the primary cosmtp ray mass. The experimental ratio distri butions obtained are compared with the expected distributions resulting-from extensive Monte Carlo calculations of EAS and their Cerenkov components for l(r < E0 < 10 еЛ and several trial composition abundance ratios. Preliminary analysis of the data suggest aa increased iron abundance relwive to the low energy composition »s well as < 20% limits on tsotropy of the abundance ratios in a ~ 120* x 2* strip of sky. Final results will be presented. Coordinates: OG 1. S. Nuclear Composition of Cosmic Rays.(Medium-High Energy) OG 1.8.4. Sidereal Variations of Cosmic Rays Mailing address: Dr. Jonathan E. Grlndlav Center for Astrophysics 60 Garden Street Cambridge, Massachusetts 02138 :w ATMOSPHERIC COHHIXTIONS KUH СНЛКСГ M'l ( IU\ WDIM-M'I'I l( \ I .1 . OF HEAVY NUCLEI MEASURED AT HALHx>.\ Л1.1 П 1 i>l •> W. HE1NRICH UNIVERSITY ofSJECEN. DEPARTMENT of PHYSICS. GERMANY Theorelicil £] Ex|4iiiin-nl.il [~\ Both Q Fragmentation experiments with heavy ions -t Berkeley have shown a factorization of the partial fragmentation cross sections into a target and a projectile depending part. This factorization makes it possible to estimate the fragmentation cross sections needed for atmospheric corrections, based on the cross sections for collisions with protons. Using this factorization, energy spectra of different elements in deeper layers of the atmosphere were calculated considering the energy loss ana the fragmentation of the nuclei. From these spectra the atmospheric corrections for charge spectra of heavy nuclei can be derived for various thicknesses of residual atmosphere. Special emphasis is layed on the discussion of corrections for measurements of radioactive isotopes. ina es. o£, i ^ g Nuclear Composition of Cosmic Rays Mailing address: W. Heinrich • Gesamthochschule Siegen Fachbereich 7 - Physik 59 Siegen 21 Adolf-R eichwein - St r. FR Germany 218 MEASUREMENT! OF 1MB C/0 AMD 1/N ABUNDANCE RATIOS IN TM РЖ1НАКТ COSMIC RADIATION IN TNI EMEkCT IAMCI 0.5-2.0 GeV/NUCLBON HADE DURING SEPTEHkU 117» J.N. Derrlckson and T.A. Pintll Space Science* Laboratory, Marshall Space Flight Caster. muetevllle, AL JM12 USA J.C. Gragory Chemistry Dept., The University of Alabama u luntSTllle, ЪтсатШ*, AL 35*07 USA Theoretical • Experimental Q *otk Q K cosmic ray experiment consisting of two xenon-filled Ion chambers, a Teflon ind a Pilot *25'Cerenkov detector, one plastic scintillator and an 8-plane MVPC lodoscope was used to detect primary cosmlo raya t > 3 at balloon altitudes In September 1976. The Ce'renkov radiator* gave useful energy measurements in the range 0.5 - .2.0 GeV nucleon"1. Preliminary data from 15 hours of flight at 4-5 millibars are presented here. The geometric factor of the instrument was ipproxlmately 1400 cm? ateradlan. Coordinates: 0C 1.5 Composition and Spectra of Cosmic lays UsffingadoMK j.H# „.mckson ES 62 NASA* Marshall Space night Center BuntSTille. AL 3S812, USA :i9 MEASUREMENT OF THE Li/B ABUNDANCE RATIO ABOVE 650 MeV/nucleon B. Byrnolt, N. Lund, I. Lundgoord Ratmuuen, M. Rotenberg Danish Space Research Institute Lundtoftevej /, ГЖ-28О0 Lyngby, Denmark A balloon-born» coimic-ray telescope containing 4 Cerenkov counters and a neon flash tube hodo scope was flown from Sioux City, Iowa in September 1974. The data have been re-analyzed, using in formation of flashed tubes to reject events due to proton and He interactions. The new event selection has reduced charge overlap between He and Li to an insignificant level, so that abundance ratios in volving Li can now be measured with the instrument. Precise correction factors for events lost by ineffici ency of the hodoscope hove been derived from flight data; they lead to improved accuracy in measurements of light elements. Introduction Data obtained in a balloon flight from Sioux City, Iowa, in September 1974 from which some results were reported at the Munich Conference, hove been ге-onalyzed. The new event selection provides almost complete charge separation between He and Li, thereby permitting the Li abundance to be measured. An uncertainty was formerly introduced by the unknown efficiency of the neon flash tube hodoscope; the efficiency has now been measured as a function of charge so that an accurate correction can be applied to mea sured abundances of the light elements where the correction is most important. The Instrument Fig. I shows the configuration of 4 Cerenkov detectors separated by 3 flash tube trays. Each Cerenkov radiator was placed in a light diffusion box and viewed by twelve 5" photomultipliers. Further details have been given elsewhere'. Data Analysis Events for which a unique particle trajectory could be identified in the hodoscope were collected, and their detector signals corrected according to calibration maps obtained from flight data. The corrected signals in the two PVT counters were required to be con sistent. Background Rejection Li nuclei are 10-3 times less abundant than protons and He nuclei, some of which inter act in the instrument and produce a shower background from which Li events must be ::о extracted. Requirements of con listency between the PVT sig PVT PLAS4C n.lS9 nal» is not sufficient to reject this bock ground, sine* the to i > T !HAr lerance mutt bt wide enough No I to allow fo' the wide distribu tion of Landau fluctuation! AEROGEL n.106 characteristic of the light ele 1 ments; therefore the flash tube r :- ~ •"- - Я hodoscope was used for direct detection of particle showers. No ? For this purpose the top flash TEFLON n=l 32 tube troy was required to be free of extra flashes (i.e. j- FT-TRAY flashes not closely associated with the particle trajectory). No 3 The good charge separation ob tained in this way is evident PVT PLASTIC n:= 159 from Figs. 2 and 3; the bock- ground contamination that re mains in the Li region was cal CM culated to 37 events, and sub 0 5 10 tracted from a total of 995 ac F1G1 cepted events. By this selection strategy one will loose some events due to the fact that the hodoscope is sensitive to 6-electrons (although this sensitivity has been artificially reduced); the loss increases with charge, as can be seen from Table 1 (compare two first lines for charges above Be). Thus our use of the hodoscope for background rejection introduces a charge bias. The following procedure was used to correct for this bias: 1) the loss of Li and Be events was calculated from the known percentages of lost events for 5^Zi8 by a straightforward extrapolation, using a linear fit on a Z2 scale, and used to cor rect the number of accepted Li and Be events; 2) the number of accepted events for 5iZi8 was replaced by that obtained when extra flashes were allowed without restric tion. This procedure is based on the assumptions that the shower background due to pro ton and He interactions vanishes for charges above Be, and that all other interactions are rejected by the criterion demanding consistent signals in the PVT counters. The correction affects the measured Li/B ratio by ~ 2%. The resulting abundances were corrected for interactions in the instrument and 4gcm of overlying atmosphere^. Results are shown in Table I. The various correction factors are shown in Table 2. Sources of Experimental Uncertainty include: the approximation in the above linear extrapolation (*: 3% estimated); back ground subtraction (± 2%); statistics tt 2% originating from the measurement of hodo scope efficiency, ± 3% from other statistical limitations). Including a ±4% uncertain ty in corrections involving cross-sections, one arrives at a total uncertainty of ± 7% of the measured Li/B ratio. 221 Discussion rVevlouely publhhed abun- done»HWMWUBII typical ly correspond loJ-l/C ratios of 0.18^22 3^ог<|Ш/В ratios of 0.67-0.77 3"7, In agreement with our results of Ll/C - 0.21 10.02 and Ll/B = 0.70 ±0.05. A few measurements"'? or» In dic- ogreement with our results. It It Interesting to compere the maMtd Li/B ratio to theore tical predictions, line» the ratio depends on secondary collisions and thus contains information of the possible confinement of cosmic rays to galactic neighbourhood. 50 100 ISO A leaky-box calculation by Signal in lop PVT (Chann*l NumbtO Shapiro et ol.,0 yields Li/B ~ 0.71; on essentially simi lar calculation '' but inclu Fig. 2 ding energy dependent cross Plot showing corrected teflon vs. top PVT signal sections yields Li/B = 0.67. for accepted events, regions used for charge iden From a recent calculation tification (separated by solid lines), and expected assuming cosmic ray produc teflon signal for kinetic energy 650 MeV/nucleon tion in spiral arm» and con- ., (dashed curve). finement to the galactic halo one may find Li/B~ 0.74, to be taken with reservation since the model has not been adjusted to fit the l/M ratio accurately for energies below 1 GeV/nucleon'"'. These pre dictions are hardty significantly different, in view of uncertainties in the cross-sections; they all agree with the measurement. A high-energy measurement would be better suit ed to distinguish between the two confinement mechanisms, since the models differ ap preciably at high energies. On the other hand the present measurement supports the as sumption, underlying both models, that the entire L-component is produced in transit. Fig. 3 - Charge histogram for accepted events. Acknowledgements I he balloon tlight was part of the Soclay Soctay-Lyngby collaboration and sup ported by NASA through the HEAO- project. Sauaic Rov'. of vdocily-Corrtctcd PVT Sgnal 222 TABLE I Observed and Corrected Number of Event* above 650 MeV/nucleon LI Be В C O Li/6 Extra flashes not ol lowed 995 434 1256 3681 29X 0.79 Extra flcehes ol lowed 1318 561 1509 4577 3763 0.87 Corrected to top of instrument 1771 790 2328 7155 6069 0.76 Corrected to top of atmospheret 20.7 B.7 29.8 100. 88.1 0.70 (C= 100») TABLE 2 Coefficients used to Correct for Events lost in Instrument and Overlying Atmosphere Loss in instrument by 6-electrons Atmospheric correction LOSS in hodoscope by interactions Li 1.18 1.20 1.31 0.92 Be 1.20 1.13 1.35 0.87 В - 1.11 1.39 1.01 c - 1.П 1.42 1.11 N - 1.09 1.46 1.06 O - 1.08 1.49 1.15 References 1. Funch, O., Iversen, I.B., Lund, N., Rasmussen, I.L., and Rotenberg, M., 1973, 13th Int. Cosmic Ray Conf. (Denver), 4, 3023 2. Meyer, J.P., Caste, M., and.Westergaard, N., 1975 14th Int. Cosmic Ray Conf. (Munich), ]2, 4144 3. Balasubrahmanyon, V.K., and Ormes, J.F., 1973 AD. J. J86, 109 4. Brown, J.W., Stone, E.C., and Vogt, R.E., 1973 13th Int. Cosmic Ray Conf. (Denver), J, 556 5. Juliusson, E., 1974, Ap. J. 191, 331 6. Julliot, C. Koch, L., and Petrou, N., 14th Int. Cosmic Ray Conf. (Munich), ]2, 4118 7. Smith, L.H., Buffing ton, A., Smoot, G.F., and Alvarez, L.W., 1973 Ap. J. 180, 987 223 8. Arens, J.F., and Ormes, J.F., 1975 Phys. Rev. DU. 1920 9. Webber, W.R., Danle, S.V., and Kish, J., 1972 Ap. and So. Scl. ]5, 245 10. Shapiro, M.M., Slloerbere, R., and Tsao, C.H., 1975 Hth Int. Ccwmic Ray Conf. (Munich) 2, 532 11. W«steroaord, N.J., 1975 (private communication) 12. Peters, B., and Weitergoard, N.J., 1977 Ap. and Sp. Sci. (In print) 13. An improved version, including energy lot» by ionization, is reported in paper OG-110 of this conference 224 THE RELATIVE ABUNDANCES OF THE ELEMENTS SILICON THROUGH NICKEL IN THE LOU ENERGY GALACTIC COSMIC RATS M. Garcla-Nunoz, G. M.-.Mason «nd J. A..Simpson Enrico Fermi Instltuti, University of Chicago Chicago, Illinois 60637 (USA) We report measurements of tht relative abundances of the elements SI through N1 In the galactic cosalc rays In the energy Interval 72 to 450 MeV/nucleon using data collected by our cosmic ray telescope on the 1MP-8 satel lite. The Measured abundances are compared with propaga tion calculations using various distributions of path lengths. It 1s found that the Measurements favor an exponential distribution of path lengths truncated at short path lengths» The source abundances of Si, Ca, Fe and N1 derived by extrapolating the Measured abundances back to the source are shown to be comparable to the solar system abundances. The relevance of our measurements of Sc through Mn to the Nn54 radioactive decay Is examined. 1. Introduction. The elemental abundances of the nuclei silicon through nickel are of special Interest In cosmic ray physics since they provide unique Information on three questions bearing on the cosmic ray origin and on the nucleogenesis of the accelerated matter. First, the path length distribution In the Interstellar medium assumed for any theory of cosmic ray propagation must satisfy simultaneously a) the observed yield of light secondaries LI, Be, В which arrive through relatively long path lengths leading to a mean path length of 5.5 g/cm* (e.g. Barcla-Munoz «rt al. 197/a), and b) the observed heavy secondaries resulting from spallation oT primary Fe nuclei, I.e. Sc - Mn, whose yield Is sensitive to the assumptions concerning the relative abundance of shorter path lengths. Second, with a consistent model for cosmic ray propagation It Is then possible to extra polate observed relative abundances back to the abundances of the accel erated matter In the sources. These source abundances place meaningful constraints on permissible scenarios for nucleosynthesis of this matter. Third, the Isotopes of the heavy secondary components Include radioactive species such as 53Mn and 54мп which have been suggested as radioactive "clocks" (e.g. Reames 1970; Casse 1973), complementing the '°Be clock 1n the measurement of the age of the galactic cosmic rays. These factors have motivated our continuing program of i;ud1es for the elements Si - N1 beginning with our report of the first satellite measure ments with sufficient resolution to separate these elements (Garcla-Munoz and Simpson 1970; Cartwrlght et al. 1971, 1973). We now discuss briefly a new set of experiments on the satellite IMP-8 based on 19,167 hours of data between November 1973 and August 1976 taken at distances from the Earth of 24-46 earth radii. 2. The Experiment. Our IMP-B Instrument and our method of analysis are briefly described 1n the conference paper by Garcla-Munoz et al. (1977b). Figure 1 shows the ensemble" of histograms derived for the iTements S1 through Ni with the element separation measured In charge units. The number 225 of events correspond ing to each c) eatent 1s obtained by fitting an ow iimaa or in nun woo * o? maximum likelihood m тж идем* смие «m " curves to the corre J sponding histograms. I- »| These fits yield < I ,| charge resolution of 0.15 charge units for * SI to 0.30 charge units for Fe end N1. » The relative abundances ,? at 1 Ml have been - obtained by normal- I' 9| Izlng to а common «е* energy Interval based 1' on the Measured SI i: J and Fe differential energy spectra. No •a"i r* •••» « i м i и W*i. corrections for the _. . nuclear Interactions p19ure ' in the telescope were required because calculations show that a) nuclear Interactions In the material 1n front of the 04 detector would decrease the Sl/Fe ratio by only 0.8* and b) since, for many nuclear Interactions 1n the thick Csl crystal detector 04 the secondaries do not escape the crystal, the interactions at most decrease the measured Si/Fe ratio by 2.5X. The abundances resulting from the analysis are shown in the second column of Table 1, normalizing to Fe * 100. For comparison, we have included in the fourth and fifth column measurements at slightly higher energies by Webber et al. (1972) and Fisher et al. (1976) from balloon flights. In Figure 2 we compare these results wTtH-the solar system abundances as given by Cameron (1973). 3. The Path Length Distribution for Fe. Figure 2 shows the well-known fact that the measured abundances of the elements Sc through Mn are much higher than for the solar system abundances. If we assume that these elements are also very scarce In the sources of cosmic ray matter, their measured abundances may then be attributed to spallation during propagation of Fe between the time of acceleration and observation, yielding Information on the amount of Interstellar matter traversed by the cosmic rays. Similar informa tion Is obtained from the production of the secondary Li, Be and В mainly by the medium nuclei C, N and 0. However, because the total Inelastic cross- section 1s significantly smaller for Fe than for CNO, the Fe secondaries arrive at the Earth via shorter path lengths on the average than the CNO secondaries. Therefore, Fe interstellar propagation together with the CNO propagation provides a tool for the determination of the distribution of path lengths followed by the cosmic rays from the source to the Earth. It has been known for a long time that an exponential distribution of path lengths (derived for Instance from a homogeneous model of cosmic ray confinement and propagation) could explain simultaneously the propagation of CI» and Fe cosmic rays. Therefore, for the Fe propagation we have adopted an exponential distribution and calculated the'resulting abundances at Earth for different values of the mean path lengths, assuming a solar system composition 226 for the Interstellar TakU I matter. The calcu lations ware carried ммима т таш» ма анслми мшомса out using • propaga •IIM U F* • 40) tion program written by J. D. Anglln whose CtlcuUM general description 1s Тгмяс. l*>.. given 1n Garcla-Nunoz *MMrt4 4 • S tlaf. U. of [lant «. «f cut cm g • 0.1 «t/aH •м МаалМга et al. (1977b) except (72-4И H,«/») (ISO №»7« «t 1 Ml) (гю-|:о от/я) (lMwtOO №>/>) Inat for our purposes SI ITS • 7 II) »f7 1» «7 17» • 3 here we assume source t 4.1 • 0.* I.MII • • l.S 14 • 3 abundances from *He through 64NI . Some of s ».« • i.a 21.2 «4.1 M • J 4< 1 1 these abundances are CI S.I «0.» 4.0 • O.t 4.2 «1.2 • 1 ' zero and 1n particular Д w.i • i.o 12.3 • I.I 13 • 2.2 11*1 we have assumed that 4 K I.J • 1.0 S.i « 0.» « • I.I • i the elements Sc through Cl M.J »i.s 23.» « 34 27 1 1 Mn are absent from the »•! source. The cross- Sc 4.4 «0.C 4.2 «O.t S.7 • l.S St] sections are given by Tl l«.7 • 1.1 is.4 • 2.s If • 2.6 II • 1 an algorithm by » 1.10.4 a.c • ].з_ 7 il.S itl Silberberg and Tsao as is.» • i.a 20 ±2.7 If i 1 updated In SUberberg Cr 17.3 *l.t et al. (1976) with the . Ml ».f • 1.0 §.7 * I.S 12 « 2.2 is • i exception that the Fc loo • s 100 100 • i 100'13 values for Sc, V, Cr Co 0.4 • о.г 0.29 • 0.04 1.4 • 0.1 - and Mn have been adjust m 4.3 + 0.6 4.6 J 0.7 2.f * 1.1 ed to the experimental values given b-y Perron (1976) at 600 MeV/ . nucleon. Because the measurements are at : Спици iwi of tht Htov, Clcmrnt Abwdonca* u low energies, the effects of solar modula [~ fht Getocit; Соктмс Rain «nd in №• S*hy Sytltm F« tion on the abundances measured Inside the solar system have been taken into account (see Garcia-Muno2 et aU_ 1977b). The Inter stellar density has been assumed to be 0.2 at/cm3, as deduced by Garcia-Munoz et al. (1977b). ~ The results are shown 1n Figure 3 where the ratio between the calculated values for the different mean free paths and measured values have been plotted for the elements Sc through Fe. The shaded area encloses the range of uncertainties In the measured values. The measured ratios fall below the shaded area for 6 g/cmz and part of 8 g/cnr and only overlap the shaded area for 10 g/cm^. Since a 20% uncertainty must be ascribed to the calculations, it can be Figure 2 seen that this disagreement, though it 1s a trend, is not conclusively proven. However, taking this trend as a real effect, it is clear that the 5.5 g/cm2 mean path lengths, found by Garcia-Hunoz et al. (1977a) for the CNO propagation , underproduces the Fe secondaries. TnTs effect was recognized early by Garcia-Munoz and Simpson (1970) in the analysis of the Fe spallation 227 Figure 3 Figure 4 carried out with data returned by the University of Chicago IMP-4 telescope. It was noted at that time that the disagreement could be resolved by assuming that the fraction of short path lengths could be smaller than for a pure exponential distribution, and it was shown that this could be accomplished by assuming a composite of a 2 g/cm2 gaussian distribution of material in gaseous shells surrounding the sources followed by an exponential distribution for the interstellar propagation. Shapiro and Silberberg (1970) also recognized this discrepancy in the Fe and CNO propagation through a pure exponential distribution and invoked a distribution with a smaller proportion of path lengths (exponential distribu tion exp [-х/Ле] truncated by a linear rise from 0 to 1 g/cm'; Shapiro et al. 1973) to obtain a better fit to the propagation of both the Fe and CNO. We have performed propagation calculations using such a distribution with the results shown in Figure 4. Clearly a truncated exponential distribution with ^ = 5 g/cmz gives closer agreement with experimental values than the pure exponential distribution iif Figure 3. The abundances obtained using this truncated distribution are compared with measured abundances in the third column in Table 1. We conclude that although our measurements do not exclude the pure exponential distribution with 5-6 g/ш* mean path lengths, they fa"vor a distribution with a smaller proportion of short path lengths? A clear selection of path length distribution must await more accurate spallation cross-sections. These results may be interpreted physically either as requiring gaseous shells around the sources as suggested by Garda-Munoz and Simpson (1970) or by diffusion models in which the solar system is placed tangential or far away from the source volume (e.g. Owens 1976). 4; The Source Abundances of Si Through Ni. Using the above truncated exponential distribution with Л- = 5 g/cm2, we have calculated the source - abundances which after propagation would yield the closest agreement with the measured abundances. Table 2 includes these source abundances along with their corresponding solar system values, all normalized to Fe = 100. Specific values are found for Si, S, Ca, Fe and Ni with upper limits established for the source abundances of the remaining elements. We call attention to the abundance ratios (cosmic ray source)/(solar system) in Table i which point to a remarkable similarity of the accelerated matter to solar system abund ances of Si, Fe and Ni. These results suggest that the well-known enrichment factor between H, He and Fe found in cosmic rays relative to solar system abundances must be mainly occurring in the region He-Si and not Si-N1. Conversely,.the light elements appear depleted relative to the component Si-Ni. Furthermore, the presence of Ni at (5 + Z)% places a clear limitation 22$ on posit bit scenarios for labia I models of nucleosynthesis of соис ми шина «мемт or *um MUXI ('• • )M» the accelerated Mtttr. Mt found earlier (Garctt-ttunot CMMC **r S»l» (C«alc Rajr Sauna/ and Simpson 1970) end con torn» Salar InUal fined her» that Ca must be in 1.» In the source to account for » 1.» our uoservatlons. It Is clear that If ме assuae an uncer s to.t 1.) tainty of a factor of 2 In the CI o.r Ca solar system abundance, ме Д H.I «•.* cannot exclude a solar-like abundance for the cosmic-ray K o.s source Ca. Ca 1.7 I.« Sc O.0M We further note that the 0.1 cosmic-ray source abundances of sulfur and argon differ » O.M significantly fro* the solar Cr l.i system abundances. However, •h 1.1 the solar system abundances of Ft 100 100 1.0 sulfur and argon «ay be poorly known. They have been actual Co < 0.4 0.1 ly derived fro* nuclear Kt S* 2 s.a O.l statistical semi-equilibrium calculations to interpolate between Z8Si and 4Qca. The above tentative conclusion that the region S1-N1 In the source elemental abundances is like the solar system leads to Interest ing speculations on the required models for nucleosynthesis which could yield these abundances. Some investigations indicate how such models may be constructed (Woosley et al^ 1973). A further consequence of the results in Table 2 is that if preferential acceleration is occurring at the sources this mainly occurs in the region hydrogen through silicon, which was suggested as an alternative by Mogro-Carapero and Simpson (1972) from solar flare studies of preferential enhancements, now confirmed by Bertsch arid Reames (1976). 5. The Abundance of Hn and the Cosmic Ray Lifetime. We have explored the sensitivity of the abundance of the element Mn to tbe radioactive decay of the isotopes Mn53 (electron capture half-life, 3.7 x 106 years) and Mn5* (electron capture half-life 303 day; B-decay half-life, 2 x 1C« years; Reames 1970, Casse 1973). J. D. Anglln has incorporated In the propagation program used in our work the decay of electron capture isotopes with 0ppenhe1mer-Br1nkman- Kramer non-radiative capture cross-sections, Inverse photo effect radiative capture cross-sections and Bohr loss cross-sections (see Raistack 1974 for a discussion of these cross-sections). Assuming that the elements Sc through Mn are absent in the sources we have carried out propagation calculations through a ^ = 5 g/cmz truncated exponential distribution of path lengths for differ ent values/df the average Interstellar density. The values of the ratio R = Mn/(Sc + Ti + V + Cr) (which is essentially independent from the path length distribution) at 250 MeV/nucleon as a function of the interstellar density p are: p = 0.2 at/cm3; R = 0.21 + 0.036 p = 1.0 at/cm3; R = 0.23 + 0.039 p = 5.0 at c/3; R'= 0.24 + 0.041 where the uncertainty in the ratios arises from an assumed 15% uncertainty In «9 the calculated element abundance. The measured r«t1o tj O.?0 • 0.025. it» conclude that the calculation^ (and experimental) uncertainties «111 not penult the selection of a meaningful value of the average 6. Acknowledgements. The construction of the University of Chicago Ht*- B experiment and the processing of the data obtained with It Mas carried out by the staff of the Laboratory for Astrophysics and Space Research. Me thank J. D. Anglln for writing the elaborate cosmic ray propagation program and T. G. Guzlk and C. DeGrazIa for computational help. This work was supported in part by NASA contract MAS 5-П067 and grant NGL 14-001-006 and NSF grant ATM 75-20407. References Bertsch. D. L. and Reames, D. V. 1976, Goddard Space Flight Center, preprint X-662-76-293. Cameron, A. G. W. 1973, Space Sci. Rev., 15, 121. Cartwright, B. G., Garcia-Munoz, M. and Simpson, J. A. 1971, 12th Int'l. Cosmic Ray Conference, ]_, 215. 1973, 13th Int'l. Cosmic Ray Conference, 1, 232. Easse, M. 1973, Ap. CHARGE AND ENERGY SPECTRA OF HEAVY COSMIC RAYS AT INTERMEDIATE ENERGIES M. Garcia-Munoz, G. N. Kason, J. A. Simpson, end J. P. Meftl* Enrico Feral Institute, University of Chicago Chicago, Illinois 60637 (USA) The energy spectra and the charge composition of the primary elements C, 0, Ne, Mg and Si have been measured In both the low-energy and high-energy modes of the University cf Chicago telescope on-board the INP-8 spacecraft. Combining both modes of analysis yields differential energy spectra for each element from ъ 50 NeV/nucleon to ^ 1 GeV/nucleon. The charge ratios with respect to oxygen are found to be energy independent over this interval and are consistent with the results of cosmic-ray propagation and solar modulation calculations. The relative abundances obtained are In substantial agreement with previous Investigations in this energy regime. 1. Introduction. The energy spectra of cosmic ray nuclei has . acquired renewed significance in the past several years with the discovery that the ratio of secondary to primary elements and the ratio of lighter to heavier primary elements decrease with Increasing energy above several GeV/ nucleon (Juliusson et al. 1972; Smith et al. 1973). It has also been reported that the change with energy of the ratios of some primary elements continues down to ъ 0.5 GeV/nucleon (Lund et al. 1975; Maehl et al. 1976). These variations have been Interpreted as evidence for, at least, two distinct cosmic-ray sources with different characteristic energy spectra (Lund 1975). The crucial experimental Issue involved in this interpretation 1s the detail ed differential energy spectrum for each primary element and the derived elemental ratios. It 1s necessary to determine the precise change with energy for a selected ratio and the region of energy Independence, If any, In order to separate possible propagation effects from source variations. This paper reports differential energy spectra obtained from the University of Chicago telescope on-board the INP-8 spacecraft. Combining data from the "high-energy" channel, reported for the first time In this paper, with the results from the "low-energy" mode (c.f. Garcla-Hunoz et al. 1977a) provides differential energy spectra covering the range from several tens of MeV/nucleon to % 1 GeV/nucleon for the primary elements C, 0, Ne, Mg and SI. These measurements continue to be the only satellite results so far reported for the study of the energy dependence of the cosmic-ray composition. 2. Experimental Details. A cross-section of the IMP-8 telescope 1s sfiown on Figure 1. Elements Dl - D3 are Hth1 urn-drifted silicon solid-state detectors, D4 is a thick (11.5 g/cm*) Csl crystal scintillator, and D5 Is a sapphire Cerenkov counter (index of refraction - 1.80). The entire telescope is shielded by a cylindrical scintillator D6 which, together with the curved solid-state detectors, restricts the acceptance cone to an angle of 60° with a geometrical factor of ^2 cm2-sterad1an. IMP-8 was launched In October 1973 and has remained highly stable over the last four years. The data presented here cover the time period 1974-1976. Three detectors are pulse- height analyzed for each event; Dl (D5), D2 and 04 where analysis for the :л Cerenkov detector 05 rtplicn 01 'or "high-energy" events. Coot-rate accumulators record the total number of ШР-rVB particles penetrating to various depth* 1n the Instrument, thereby permitting the number of analyzed events to be normalized to the total number of incident cosmic rays. It should be emphasized that the two analysis regions (high/low energy) are part of the same telescope; therefore normal ization, calibrations and systematic effects remain relatively constant 05 over the entire energy Interval studied. ,. DS* The analysis of events in the Г "3 ^ D6 low-energy mode has been described - " elsewhere (Garcia-Munoz et al. 1973, 1975a) and the high-energy analysis \ \ follows the procedure described for the Tube IMP-5 instrument (Mason 1972) with some gem minor changes. The sapphire Cerenkov ScoW- counter has a strong scintillation component which permits analysis of events which exit D4 with energy below Figure 1 the Cerenkov threshold. Thus, the "high-energy" data spans the region between the stopping (in 04) particle mode and the saturation of the Cerenkov signal at ^ 1 GeV/nuclepn. Both forward and backward-moving (entering from the bottom of Figure 1) particles are observed in the high-energy matrix, and the regions of overlap between forward and backward-moving events have been excluded in this analysis. The events for each charge were selected in the 04 versus OS matrix, and background was removed by requiring consistency in the dE/dx detectors D2 and D4. At high energies, multiple charged particles produce high-energy knock-on electrons which can trigger the shield scintil lator 06. These events were recovered since this instrument records and analyzes, with a unique identification code, events triggering D6. The energy assignment was provided by the solid-state detector D2 which displays a linear response to particle energy-loss over the region of interest. The energy scale was calibrated (to * 3%) using the polnt-of- penetration for each detector, fit to standard range-energy relations. The Cerenkov signal was not used for energy determination because the response of the sapphire Is a complicated function of energy. Two points about .the energy calibration should be emphasized. At the hlghtst energies, the distribution of signals from 02 becomes asymmetrical due to Landau-Vavllov fluctuations, and the energy assignment Is unreliable. This region of the spectrum was combined Into a single Integral flux point having an energy limit below the Vavllov region. Secondly, the energy Intervals adjacent to the Integral point are quite sensitive to uncertainties 1n the energy cali bration, and this leads to a large asymmetrical error In the differential flux points. For lower energies, the calibration uncertainty becomes progressively smaller. 3. Results and Discussion. Figure 2 shows the differential energy spectra for the elements 0, t, Ha, Hg and S1 Note that the peak of the :з: ю" ^—^гт-м-тгр •• The Uravemty 01 Crucogo ^ IMP 6 ч OMerentol Entrxji Spectra Oxygen t-i- 1974-1976 1 • -r- r-1 i-r I1йгЕ—•— t мин] 1—ч-1 I мн|—л 5 Е 10' |.o3i 4 КЗ 10 ЮО 1000 100 1000 Kinetic Energy (MeV/nucleon) Kinetic Energy (MeV/nucleon) ти Figure 2 distribution is Included in the region of the measurements. The division between the low and high energy analysis modes varies from % 200 MeV/nucleon for carbon to * 300 HeV/nucleon for silicon. The Integral points have been converted to a differential flux and plotted at 1100 HeV/nucleon for an assumed power law spectrum 1n total energy per nucleon with у * -2.65. The uncertainty limits are dominated by systematic uncertainties 1n the energy calibration and normalization corrections, and the gaps in the spectra represent the deleted Intervals due to backward-moving events. The energy dependence of the charge ratios 1s shown on Figure 3. The oxygen spectrum was selected for normalization since this element showed the cleanest data and was less subject to normalization uncertainties than carbon. The uncertainty In Interpolating the oxygen flux values contributed a negligible amount to the error in the ratios. The striking result shown on Figure 3 Is the energy independence of the charge ratios. With the exception of the lowest energy point, the data are consistent with a constant ratio as a function of energy up to 1 GeV/nucleon. The points at 20-50 MeV/ nucleon for C, Ne and Mg are low because the oxygen flux 1n this region contains an additional component ~ the "anomalous" oxygen which has been reported by McDonald etal. (1974) and Hovestadt et al. (1973). Since this component decreases rapidly with Increasing energy, it represents'only a small fraction of the oxygen flux at the energy of the first silicon point, for which no depression Is observed. The solid curves superimposed on Figure 3 show the results of a cosmic- ray propagation calculation assuming identical source spectra for all Isotopes through nickel and propagation with an energy Independent exponen- 233 -I Г—ГТГТП——T 1 I | 1 f 1 ttil path-length distribution го characterized by • емп of 6 g/ar. Tht results have u> been Modulated at tht level -N=tt*+T-Hf^: appropriate to this Им -+ ptrlod (stt Garcla-Nunoi as ' ' ' ' 'l tt al. 1977b and references 1 1 1П ITT — '' —T - T !—Г T If I I IKereTn), and the calculated a» oxygen spectrum 1s In excel lent agreement with the data 0.15 shown on Figure 2. The solid ^FFttjHH curves on Figure 3 tend to rise slightly with Increasing Q07 l ii i il P energy as a result of the o 0.3 ~1 1 1 |"T)'I1| I I I I HI difference in energy-loss a: rates for the two species 2 4 during propagation. With 0J5 qr+ "+ll Iff)' the exception of Ne/O, all of the Measured ratios are ooe in excellent agreement with -J 1 • I • • ' • I I L__J ' • • • the calculations. The in > ' "4 . Si ferred Ne/O curve is 10-Ш аг above the measured data which Indicates that the 0.1 Н~-Н^+Н+ЧНН-; source ratio of Ne/O 0X37 (-0.135) must be decreased -I Mill -I 1 I—i I I I ll Ю Ю0 Ю00 by a like amount, and this Kinetic Energy (MeV/nocteon) тя would bring the curve into agreement with the data. Figure 3 Lund et al. (197S) have presented data from a Cerenkov counter telescope which shows a change of ~ 30% in the C/0 ratio between •>» 600 and 1000 MeV/nucleon, but Maehl et al. (1976) observe a constant ratio over a similar energy interval. The present results indicate no variation (to within "x. 20%) in the C/0 ratio below 1 SeV/ nucleon and are consistent with the results of Maehl ejt al. and w'th the compiled data presented by Garcla-Munoz et al. (1973). Variations 1n the C/0 ratio have been reported at energies > 2iT"GeV/nucleon (Juliusson 1974), but these changes would not be observable in the integral point measured here. Thus, it appears that any energy dependence in the C/0 ratio must be confined to the very high energy region. Maehl et al. (1976) report variations in the Mg/0 and S1/0 ratios (also observed atiSign-energies by Juliusson 1974) but not in the Ne/O ratio. Again, the present experiment does not observe such variations. However, the S1/C ratio of Maehl et alj, departs significantly from energy Independence only for energies greater than ^ 1.5 GeV/nucleon. Our integral point for Si/0 shows no indication of an increase in the ratio at higher energies, but within the experimental uncertainties, such a variation cannot be ruled out. A similar conclusion applies to the Mg/0 ratio. It should be noted here that the absolute value of the silicon flux shown in Figure 2 is 1n agreement with earlier work (Garcla-Munoz et^al. 1975b) and is approximately 30% higher than the value reported by Maehl et al. at 1 GeV/nucleon. The curves in Figure 3 are the results of modulating the local inter stellar spectra to the levels observed in 1973. The levels deduced for 1974' :м 1976 art similar. For no solar modulation the Si/0 r*tto, which thcnn xr* largest Modulation effect, decreases frow-v- 1.45 at 1 GeV/nuclton to a «alue of 1.20 at 70 NeV/nucleon. This unmodulated curve Is consistent Kith the data within the uncertainties, but the nodulated curve gives a better fit. Due to the «diabetic deceleration suffered by cosailc rays tn penetrating the solar cavity, modulation has the effect of reducing any energy variation present in a specific ratio 1n local Interstellar space. The relatively constant ratios observed 1n this experiment Indicate that there can be no sizeable differences 1n the spectral form for these elements outside the hellosphere. The relative abundances at Earth obtained from these data for the two energy Intervals ^ 45-300 MeV/nudeon and > 450 HeV/nucleon are given In Table 1 normalized to carbon. For comparison, the table gives the recent results of Lund et aJL_ (1975) and" Fisher et aK, (1976) and the "best" values compiled by Shapiro and Silberberg (1974)Trom data available through 1973, generally at significantly higher energies. Our results are 1n substantial agreement with other Investigators. The maximum departure is about one stan dard deviation for the Mg and SI abundances > 450 MeV/nucleon. The Ne abun dance observed in both energy Intervals is slightly lower than the compiled value and the measurements of Fisher et al., as is the result of Lund et al. A glance at Figure 3 shows that, of aTT tbT ratios, Ne/0 is the only one which might be interpreted in terms of an energy variation, and this, is due, in part, to the smaller relative abundance in the 45-300 HeV/nucIeon) interval compared with higher energies. A variation several times that shown by the curve in Figure 3 is permitted within the uncertainty limits. Additional work on this question is clearly required, but the simplest interpretation of the Ne/0 ratio favors an energy Independent composition up to 1 GeV/nucleon. Table 1 Relative Abundances at Earth £ P. Ne 1Й. 1L This experiment 45-300 HeV/n 100 94 ± 2 13 ± 0.5 17 t 0.7 12.5 ± 0.5 This experiment > 450 MeV/n 100+ 89 ± 6 14+2 16 ± 2 12 ± 1 Lund et al. 1975* 0.6-I70~5eV/n 100+ 90 ± 1.3 14 ± 0.7 18 t 0.6 13.5 i 0.5 Fisher et al. 1976* 350-600"76V7n 100+ 94 ± 1 16 ± 0.5 20 ± 0.8 14 ±0.2 Shapiro and Silberberg (1974) R > 4.GV 100+ 91 ± 2 16 i 2 19 ± 1 14 ±2 '''Normalization *Bal1oon observations 4. Conclusions. The differential energy spectra reported between sev- eral tens and 1 GeV/nucleon for the primary elements C, 0, Ne, Mg and SI are In agreement with the results of energy Independent propagation calculations combined with solar modulation. The charge ratios show no variation with energy up to •v 1 GeV/nucleon which limits any substantlaTenergy dependence In these ratios to higher energies than those studied here. Thus, these 2JS results show no_ evidence for two or more different types of cosm1c-r«y sources. 5. Acknowledgements. The University of Chicago IMP-8 experiment was constructed and the data processed by the staff of the Laboratory for Astro physics and Space Research. He particularly thank E. Murphy, M. Nixon and T. Cheung for computational assistance, and are grateful to J. 0. Angltn for developing the cosmic-ray propagation code. This work Mas supported in part by NASA contract HAS 5-11067 and grant NGL 14-001-006 and NSF grant ATM 75-ГО407. J. P. Hefel acknowledges partial support from the NcCormlck Fellowsh.j and Compton funds at The University of Chicago. 6; References Fisher, A. J., Hagen, F. A., Haehl, R. C, Ormes, J. F. and Arens, J. F. 1976, Ap. J.. 205, 938. Garcia-Munoz, N.. Mason, G. M. and Simpson, J. A. 1973, Ap. J.. 184. 967. 1975a, Ap. J. (Lett.), 201. L141. 1977a, this conference, paper OG-61. 1977b, this conference, paper 0G-84. Garcia-Munoz, M., Juliusson, E., Mason, G. M., Meyer, P. and Simpson, J. A. 1975b, Ap. J., 197, 489. Hovestadt, D., Vollmer, 0., Gloeckler, G. and Fan, C. Y. 1973, Phys. Rev. Let., 31, 650. Juliusson, ГГ, Meyer P; and Mu'ller, D. 1972, Phys. Rev. Let., 29, 445. Juliusson, E. 1974, Ap. J., 191, 331. Lund, N., Rasmussen, I. L., Peters, B. and Westergaard, N. J. 1975, Proc. 14th Int'l. Cosmic Ray Conf., Munich, 1_, 257. Lund, N. 1975, Proc. 14th Int'l. Cosmic Ray Conf., Munich, 11, 1746. Maehl, R. C, Ormes, J. F., Fisher, A. J. and Hagen, F. A. T576, NASA Goddard Space Flight Center pre-print X-661-76-132. Mason, G. M. 1972, Ap. J., 171, 139. McDonald, F. B., Teegarden,~BT J., Tralnor, J. H. and Webber, W. R. 1974, Ap. J. (Lett.), 187, L105. SNpiro, M. M. and SlTbTrberg, R. 1974, Phil. Trans. R. Soc. Lond. A., 277, 319. Smith, L. H., Buffington, A., Smoot, G. F., Alvarez, L. W. and Wahllg, M. A. 1973, Ap. J., 180, 987. 236 пиеш* сомк* m« or тж COSMIC ми кгтх о. s «ив 10 ОЖГ/NUCUOM В. Byruak, M. band, I. bmdmaard М—им» and N. Kotaaber*. Danish Space Рмамгса Institute Uedtoftevej 7, OK 2800 I/ngby, Duaark Theoretical Q ExiMriMMtal Q] Bo(h Q A balloon borne counter teleacope waa flown for a total expot'ir* timm of 60 hours from Sioux City, low» in the fall of 1974. An analysi» of 160000 erenti baa been performed. The significant .variations within a limited energy interval of the relative abundance of.B,C,N,0 which were reported at the Minich conference have, been reinvestigated with aore data. The new results will be presented and their astrophysical implications discussed. Coordinates: OG 1.5 (Nuclear Composition of Cosmic Rays) Mailing address: r0 Lundgaard Rasmussen Danish Space Research Institute Uutdtoftevej 7, 2800 Iyngby Denmark 237 ТИ DH.ICATIOW ТОК GALACTIC FKOrACATICM FROM A *ASUWD*JfT OF ПОЕ ОШа (1-5-26) AMD KMIBCT (0.5-5 BtV/nuc) SPKCTKA OF COSMIC UTS J. A. Lecalak, W. ft. Webber, J. C. KUh «id C. A. Simpson Fhyaica Department University of On Hampahlre Durham, M. H. 0382* 0. S. A. The energy spectra of cosalc raya of charga 5-26 have baan measured ov«r a kinetic-energy Interval of >0.5 to "50 BeV/nuc using solid and gaa Cerenkov datectora flown aboard hlgh-altltude balloooa In the summer and fall of 1974. Tha data la examined and Interpreted In tarsal of an equilibria "leaky box" galactic propagation aodel which haa an exponen tial dlatrlbutlon of path lengths P(x)-e X« where the es cape mean' free path X, la a decreaalng function of energy. 1. Introduction. Tha energy spectra of cosmic raya of charge Z-5-26 ware measured отег a kinetic energy Interval extending from "0.5 to 'SO BaV/nuc utilizing the signals from a aolld Cerenkov detector (C) and a gaa Cerenkov detector (G). These detectors are identified In Figure 1 where the experi ment used to make the measurements Is shown In cross section. This experi ment was flown in the summer and fall of 1974 at Ft. Churchill, Canada and Sioux Falls, South Dakota. The accumulated geometric factor for the meas urements Is 7600 m2-eterad-sec. A more detailed presentation of the data, the charge rstlos etc., will be made In a future publication after the 1974 data has been combined with ad ditional measurements obtained from a balloon flight in 1976. Here a more limited data presentation will be made with the main purpose of the present paper being an Interpretation of the groas properties of the 1974 data in terms of n awwccsMciw galactic propagation. fl w"!£wS 2. Data Analysis. Briefly, consistency cri teria are applied between the two dE/dx scintillation detectors S^ and S2 (see Figure 1) to remove background and enhance charga Sj+S» resolution on cross plots of ^ versus C. Once charge resolution has been optimised, pulse height, distributions are then obtained for each of the Cerenkov detectors C and G for each charge. These pulse height distri butions are then deconvolved (Lezaiak, 1975) to extract the underlying energy spectra. Finally these spectra are corrected for the removal of events due to the Imposition of consistency criteria and for Interactions within the instrument as well as Interactions ЯС45"NEK» and energy loss In the overlying atmosphere. A more complete discussion of the data anal ysis will be presented In a future publics- F-tguAe 7. A anoii action oj tion. оил 1Ч14 coimic-лау zxpeMment. lit 3. Date гт—antatlom and Inter pretation. A charge histogram for kimattc' ««rtlM >13.1 lev/ nue obtained from the |м Oeren- kov 4*tector data Is praaaotad In Figure 2. This histogram summarise* the charge compos1- tloo of cosmic ray* at high en ergy and can beat be lntarpret- ad through a comparison with similar data obtained at a much lower energy. Tigwu t. ChaAQt киюдлля {o* iote&c Such a comparison la tntAgit* >13.1 StV/nuc obtaintd with ahovn In Figure 3. Bere the оил дал Ctxuikov dtttcXoA. ratio of the Integral Intensity above a kinetic energy of 13.1 BeV/nuc (from the gaa Cerentcov detector) to the Integral In tensity above a kinetic energy of 0.45 BeV/nuc (froa the solid Gerenkov detector) la presented as a function of the cosmic-ray charge. The data for the cos mic-ray primaries (even charges) are shown as open circles, and the data for the cosaic-ray secondaries (odd charges) are shown as solid circles. Above the charge 2-8 various charge groups are sunned as Indicated to increase the' limited statis Ptgu/te 3. InttgAat intimity Aatioi tics available at high energy. veMui cna/tge. This data shows some remarkable trends discussed by:other authors (e.g., Balaaubrahmanyan and Ormes, 1973; Jullusson it at., 1972; Smith tt at., 1973; Webber it at., 1973) but never presented In this format before. Firstly, the Integral Intensity ratios demonstrate a trend to systematically Increase with charge for both the secondaries (dashed line) and the primaries (aolid line). And secondly, the Integral Intensity ratio for the cosmic-ray secondaries (dashed line) Is systematically lower than the similar ratio for the cosmic- ray primaries (solid line). Thus at higher energies, secondary cosmic rays are less prevalent relative to the primaries than they are at lower energies, and the higher? charges are more prevalent than they are at lower energies. All this qualitatively agrees with the equilibrium "leaky box" model of cosmic-ray propagation first put forward by Cowslk (1967). If the escape mean free path is allowed to be a decreasing function of energy, then cosmic rays will escape the galaxy easier at higher energies and thus their secondary pro duction will be diminished relative to that at lower energies. Also as 'charge Increases the interaction mean free path decreases, thus making fragmentation more competitive with galactic escape and thus making the observed spectra progressively flatter than the source spectra. A differential total-energy spectral Index scale corresponding to the Integral intensity ratios is indi cated at the right In Figure 3. и» In order to explore the nature of Cbe escape aean free path «ore quanti IIM| ГТТТТТу Т Г1 I MT' tatively, we'll cooatdar seconder? to primary nuclei ratio* obtained fro» the differential intensity spectra of nu clei from boron to oxygen. ТЪаае are •hows la Figure 4. Tha data for boron to oxygen contain! moat of tha actual events aeee by tha instrument at kinet ic energlea »10.5 leV/nuc (the gaa Car- епкот detector threahold energy) and s thua contalna aoet of tha Information ' 1 which we have available at high energy. ^ For tha Incite Cerenkov detector only i ЛС«0) data at kinetic energies £1 BaV/nuc are conaldered in Figure 4 in order to ig • Р(ж)-Л*« к«-к,<Т|**//мс»>' nore the effects of galactic lonisation energy loaa and aolar modulation ef 1 '"I * ' fect* in the aubaequent analyaia. Ю oo Kknttc Ewty T.BeV/iwe The M/0 and B/(OH>) ratio* *hown in Figure 4 are fitted with the expres- flguAt. 4. VlHuuiutiat лм&ииОц alons indicated. These expressions, Katioi vvam елелду. were obtained froa an equilibrium "leaky box" propagation aodel where the effecta of diffusion out of the galaxy can be deacribed in ten» of an exponential path length distribution P(x)- e e . The "Kt" and "Ay" parameter» are known total and partial fragmenta tion aean free paths. The escape mean free path \^ is taken to have a power- 2 law dependence on kinetic energy Xe-X0 (T(BeV/nuc))V gm/cm where the param eters X0 and у are determined by obtaining the beat nonlinear weighted leasts- squares fit to the data. Bast fits are obtained separately for the N/0 and B/(0f0) data and are Indicated by the solid curved lines. Note that a aource abundance of nitrogen equal to 10X of oxygen was assumed independent of ener- gy. T T" The best-fit values of the parameters XQ and у so obtained are shown in Figure 5 In terns of tiff elliptical regions on a X0 vs. у plot. The two aeparate esti mates of the parameters are cloae ГIQ2(T(BtV/noclf * to one another on this plot given aaV. ^irumi the approximate model employed, C and we take a weighted average value of.tha parameters to give VMT< •*">«"» the expression 10.2 (T(BeV/ --365 4 nuc)) gm/or as the best ~f ( KttfftOQOTW I ea£ie»t«_i>! the escape mean free path from this"data. , I. • , i It is of interest to note that the value of the exponent ao Пдиле. 5. Слон plat o& the. poAcutotoiA deduced lies In the -1/3 range, \0 and у uizd to рслая&Хл1сдйц dz&cMbe. which is what one would-expect if the. ucapt mean put path. 240 the lataretallar maaMtlc-flald power •peetnaa caa ha described la tana of a Koloaogorov amactnm. The Kolomo- пт т шая» у» mat и»ч ТУ— gorov apectnaa la a aatural caacada •Ut%MMUl a»l*>*a*tl tJMtaa> <-»ф •Ms) )tl> * • * apactnsa (Oaodraaekhar, 1*49) which «1 la> 4» tot/M vary llhely axlata In tnterateller • 1 apaoa. » 1 i.« •a.»*»"* -i »w u -* «XI •* c I. Ml. «Hit'* -глнм Tabla 1 aoamarlaaa tha meaaura- I ».t 1.} -«-* -j.m.n manta diacuaead .hare In tana of flta 0 2.Ml.Ma»"* -a.Mt.U -*.>*>.«! PNMA1 l.l>a.M«l*>-* -X.Ua.M of tha data to total-energy apactra. •ИЩИ1 l.Mm-.M-.t-' •J.Wl.ll -* *•«.« •>.*•*.u Tha flrat eoluan Indicate* the charge. •Mt.ll.W -i in.» Tha next two coliama give tha param- aW21*t) -2..П.4» etera Е and Y. obtained by fitting Pa, -1.4*1 U tha differential Intensity data from "1 to "SO BeV/nuc to tha axpreaaloa ](Е)-ГВУ nuclei/(cm2-eterad-eec) Table 1. Total tMAgy dititAtntuU «hare Е la tha total energy par nu •itittnitdf ilti to tht data. cleoli axpreaaed In unite of BeV/nuc. The leat column" Indicate* the value of the spectral index parameter determined In a different aanner--aaaely from the integral flux ratloa presented In Fig- ure 3. The apactral exponenta derived by the two different approaches agree wall with each other when both values exist. The spectral exponenta Indicate that, in general, the odd or secondary nuclei have steeper spectra then the even or primary nuclei. Also the spec tral exponents for both the even and odd charges become flatter as charge In creases. This la consistent with the galactic propagation model discussed earlier. It la also interesting to point out that ivon haa a spectral exponent which la only about .18 of a power in total energy flatter than carbon or oxy gen nuclei. This can be explained In terns of the same model used to analyze the B/(C+0) and N/0 differential Intensity ratios without having to Invoke different source spectra for these primary nuclei. 4. Summary and Conclyaiona. We have interpreted the charge (Z-5-26) and en- ergy («0.5 to =50 BeV/nuc) data obtained from our cosmic-ray instrument flown several times during 1974 in terms of an equilibrium "leaky box" model of galactic cosmic-ray propagation. Katios of the integral Intensity data obtained at high and low energies show (1) the effect of a galactic escape mean free path which decreases with Increasing energy, allowing the cosmic rays to more easily escape at high en ergy and thus result In lower secondary production, and (2) the effect of a fragmentation mean free path which decreases with Increasing charge and thus makes fragmentation more competitive with escape at the higher charges, re sulting In apectra which become progressively flatter than the source spectra as charge increases. ' Secondary to primary ratios of differential spectra were analyzed to obtain a parametric description of the escape mean free path. The best-fit 365 2 expression was found to be Ae-10.2T~" gm/cm wheie T is the kinetic energy expressed In units of BeV/nuc. 241 The data from "1 to «SO BeV/uuc waa alao fit to total-energy dlfferem- tlal epectre. Ska data again llluatrataa the fact that the eecoodary nuclei have ataapar energy epectra than the primary nuclei, and that the energy apactra become flatter aa charge Increase*. Lastly we wish to note that, although the preaent Interpretation has been in tezmo of a single "leaky box" nodal for galactic propagation, a multiple "leaky box" model (Covalk, 197S) nay be necessary to explain the low anlaotropy of coealc rays at high energies. Measurements of secondary and primary cosmic rays st even higher energlee are neceaaary to eetabllah whether such more complicated models need be Invoked. teferences Chandrasekhar, S., 1949, Aa trophy*. J., 110, 329. Cowalk, R., Pal, T., Tandon, S. N.. and Vena, R. P., 1967, Fhya. Rev., 15S, 1238. Covali, R., and Wilson, L., 1975, Proc. 14th Intl. Cosmic Ray Conf., 2, 659. Balasubrahmanyan, V. K. and Ones, J. F., 1973, Ap. J., 186, 109. Jullusson, E., Meyer, P., and Mailer, D., 1972, Phys. Rev. Ltrs., 29, 445. Lezniak, J. A., 1975, Hucl. Inst. Heth., 126. 129. Smith, L. M., Bufflngton, A., Snoot, G. F., Alvarez, L. U., and Wahllg, M. A., 1973, Ap. J., 180, 987. Webber, W. R., Lezniafc, J. A., Klsh, J. C., and Damle, S. V., 1973, Nature, 2U, 96. •ч4 Cosmic IUy Composition Batvaan 5 and J00 C»V/Kucl»cn is1, V. K. liluubrhuunykn1, J. F. On»1, V. K. H. Schmidt'-' I. Simon3, F. Slohan2, H. Splaeelhaoer }. C. I. Yodh1'* 1. NASA/Coddard Space Flight Canter 2. NAS/MtC ««search Associate at NASA7GSFC. Craenbclt. KD 3. Max-Planck-Institute filr Extreterrcatrlache Physics. D-8046 Carctilng. U«at Germany 4. Phyalcs Department, University of Maryland, College ParK, KD Theoretical • Experimental [5Г] Both Q Cosmic Ray energy apectra from the analysis of a large area thin calorimeter will be presented. The data available at the time this abstract Is written indicate that individual changes are well resolved and the energy distribution of all the abundant elenents frcn C f> *> can be determined. Final results are expected to be ready for the conference. Ue anticipate that this experiment will yield detailed energy spectra of individual elements up to 300 GeV/nuc for Fe (and higher for more abundant nuclei) for the first time. The instrument will be described in detail to this conference in a separate paper (J. F. Arens e_t a^.). The inclusion of a gas Cerenkov counter with well known response into the instrument enables us to determine the energy, distribution up to 60 GeV/nucleon independently of the calorimeter. This at the same time is used to calibrate the calorimeter so that its energy response function is now known much better than previously. The results and their implications to the source mechanisms and propa gation models of cosmic rays will be discussed. Coordinate»: 0G 1.5 (Nuclear Composition of Cosmic Rays) Mailing addien:. Wolfgang K, H, Schmidt Code 661 NASA/Goddard Space Flight Center Greenbelt, MD 20771 U.S.A. Астсяес CHARGE COMPOSITION AND ENERGY SPECTRA Of COSMIC-RAT NUCLEI AT ENERGIES ABOVE 5 GEV PER NUCUON John H. Caldwell and Peter Meyer Enrico Feral Institute, University of Chicago Chicago. Illinois 60637 USA A scintlllatlon-Cerenkov counter telescope, with three gas Cerenkov counters for energy determination between 5 and 90 GeV/n, has been exposed for a net total of 4.S w? sr hr In two balloon flights In 1974. The measurement yields the chemical composition and energy spectra of cosmic-ray nuclei In the charge range 5 * Z s 28. The differential spectral Indices of oxygen and the Iron group nuclei are measured to be 2.67 ± 0.04 and 2.SO t 0.08, respectively. The results are Interpreted 1n the context of the "leaky-box" model of cosmic-ray confinement and propagation. The energy dependence of the cosmic-ray leakage path length, A, 1s 59 -09 found to be X(E) « Etot~°' * ° . After Incorporating this energy-dependent propagation Into the model, It Is found that the measurements are consistent with an energy- independent source composition for the cosmic rays. 1. Introduction. A powerful method in the study of the origin and propagation of the cosmic radiation is the measurement of the charge composition and energy spectra of cosmic-ray nuclei. The nuclei that are detected at earth are composed not only of primary particles, which originated in the cosmic ray sources, but also of secondary nuclei which were produced 1n nuclear interactions of source nuclei with the interstellar gas .during their propagation. From the relative abundances of the primary particles, one can infer the nuclear composition at the source. On the basis of spectroscopic measurements and nucleosynthesis calculations, 1t 1s believed that some nuclear species which appear as secondaries are completely absent at the cosmic-ray sources. A measurement of the abundances of such species can yield a value for the amount of matter traversed by the progenitors of the secondary nuclei. Measurements reported In the last few years have conclusively shown that the relative abundances of cosmic-ray nuclei are different above about S GeV/n than they are below a few GeV/n (Juliusson, Meyer, and Muller 1972; Smith ejt aJL 1973; Balasubrahmanyan and Ormes 1973; Jullusson 1974). The most pronounced change observed Is the decrease at high energies of the relative abundance of secondary nuclei when compared to primary particles. This has been interpreted as indicating that high-energy particles traverse less matter than those at low energy, and thus produce fewer secondary nuclei. An unambiguous explanation for the origin of this energy-dependent path length, however, has yet to be advanced. 2. Instrument and Balloon Flights. The detector used in this experiment is a scintlllatlon-Cerenkov counter telescope, which 1s shown schematically in Figure 1. The two scintillation counters, CI and C4, are the charge determining elements. Both, scintillators are made from Pilot Y and are 11 mm thick. They are situated in light diffusion chambers to maximize uniformity of response over the detector surface. Another technique used to Improve the charge resolution of the instrument was to shape CI and C4 to reduce path г** length variations In these counters. The high pressure 9» Cerenkov counter CJ Is the major мм component of the tnstnawnt. It Is a light diffusion chamber that 1s filled with ethylene gas at • pressure of IS.6 atmospheres it 20* C. The thres hold kinetic energy of this detector 1s 5 GeV/n. The mean path length through the counter Is 24 cm, along which a singly-charged highly relatlvlstlc particle produces a signal corresponding to 12 photoelectrons. The other two gas Cerenkov counters C2A and C2B operate at a pressure of one atmosphere. C2A Is a light diffusion chamber, filled with Isobutane gas, which has a Cerenkov threshold of 17 GeV/n. C2B, on the other hand, Is a focussing gas Cerenkov counter which contained CO; during the first flight and Freon-T2 for the second. The thresholds of these two gases are 30 and 19 GeV/n, respectively. The instrument also contained an annular guard scintillator. This counter 1s Fig. 1 Schematic cross-section of •;red to define the coincidence the Instrument. requirement, CI • ТГ • C4, which has a geometric factor of 910 cm? sr. The guard counter is also used to help in the rejection of non-nuclear background. It, and all pther counters In this instrument are pulse height analyzed. The instrument was flown twice from Palestine, Texas 1n the fall of 1974. The first flight remained at a ceiling altitude of 4 g cm-2 of residual atmosphere for 12 hours. The second flight lasted over 42 hours, although during the second night the balloon descended from 4 to ab?''t 11 g cn*z. The total net exposure factor for the flights amounted to 4.5 m* sr hr. 3. Data Analysis. In order for an event to be accepted for further analysis, it must fulfill two main selection criteria. These criteria are designed to select nuclei that pass through the Instrument without interacting and without accompanying background particle's. The first criterion rejects events that produce a large signal in the guard counter. The guard counter .cannot be in strict anti-coincidence, however, since delta-rays produced by a nucleus can cause a signal in the guard, although the primary particle passed through the hole in the guard counter. The selection, criterion used is that the signal In the guard cpunter be less than 82 of that signal which would have occurred, had the primary nucleus passed through the guard scintillator. The second selection criterion is consistency fn the pulse heights from the scintillators CI and C4. In order to select non-interacting nuclei, it 1s required that the ratio of tne"~5lgna1s produced In the two identical scintillators CI and C4 vary from unity by less than a factor that is-equal to 202 at beryllium and decreases linearly as a function of atomic number to 8X at Iron. A charge- dependent criterion 1s applied In order to Insure that Interacting particles are effectively eliminated. The charge of a nucleus 1s determined from the average of the square roots of the signals in the two scintillators CI and C4. The scintillator signals were corrected for saturation, drift, and finite channel-width effects ana small :*5 non-linearities In the pulse height analyiers. A Msto- , дгм of nuclei with energy „ greater than 5.4 GeV/n that ,. satisfy tht guard counter and *, scintillator agreement sal- f, ectlon criteria Is shown In ° Flgure-2. The standard dev- * latlon of the charge measure- t 4 V ment at carbon Is 0.18 charge ~ units. J The energy Measurements made « > ' ' 5 ' « « 5 3 5 т by this Instrument are accom- »ншм гм* и .« ri^w^t.^r?*^ г?д «9- 2 Charge histogram for nuclei «1th Cerenkov counters C3 and CZA. ' „„.J,,, ».»I.. »>,».. с л r«u<« For the purposes of this ener^ ^n r*ater tnan 5'* GeV/n' paper, we divide the energy region measured by the Instrument Into three parts. Low energy nuclei will correspond to particles above the geomagnetic cutoff, yet with Insufficient energy to trigger the high pressure gas Cerenkov counter. This region thus contains energies from 1.6 to 5.4 GeV/n (median energy Is 2.5 GeV/n), although the lower limit is approximate because of the drift of the instrument in geomagnetic latitude during a balloon flight. Medium energy data will refer to those events that do produce Cerenkov light in the high pressure gas counter, and are below the Cerenkov threshold in the C2A gas counter. Thus this energy region corresponds to 5.4 to 17 GeV/n. Nuclei with energy greater than the C2A gas Cerenkov counter threshold, 17 GeV/n, are denoted as high energy. In order to extract an energy spectrum from a Cerenkov pulse height distribution, one must 1) deter mine the location of the Cerenkov threshold in the presence of residual scintillation, 2) determine the В =1 point of the distribution, and 3) cal culate the deconvolution corrections that must be applied to remove the effect of finite detector resolution. The 8 = 1 point is the signal that a particle traveling with the speed of light would produce 1f the detector had infinite pulse height resolution. The deconvolution correction factors are calculated by the statistical deconvolution technique described by Juliusson (1974). 4. Results. Events that satisfy the two main selection criteria are divided into charge and energy bins. A background correction was made for boron, amounting to the removal of 6% of the events In that charge interval for the low and medium energy data. We also require that a high energy event possess a large signal in the C3 Cerenkov counter. The number of nuclei with 5 lAtn oo*>fir^urir>«f«e«nm rei-OO^OfO>-p-r- o o o r- o o o o o o 0.0 Or-Or-OOOOOOOO <\||Ле\|.»Ог— Or—Or—r—r— oooooooooooo doododddddo'd dddddddddddd rj Sioj < Or-Or-Or-r-r- oa-S6f ONOr-or- ood^oooodooo OOOr-OOOOOOOO «con_ n _1ЛСр00Г^г^*ЛГ'г cvCVJLf>{\JuT)©rj - — Or—Or-r-r— o-oooooooooo o or— d.— oooooooo oc> dddddddddd +1 +1 +1 +1 +1 +1 +1 +1 •*•! +1 +1 +1 oooooooooooo Г0<Л№О№01О1Лг*00г-Гч ^niAQCVIinnotnSi-Ol nmnQi- momf*«NCO NONOOi- Or— Or— f— O OOr- OOr- ONOr-r- r- Or—Or—oooooooo O — O—OOOOOOOO CM U) CNJr*,— <\jcnr—rtf-OJ^CfJfr» olsslslsisss S — Or-OOOOOOOO oooooooooop o doo odddododo +1 +1 +1 -H +1 +1 +1 +1 +1 +1 +t +1 +1 +1 +1 +1 +1 +Г +t +Г +t +1 +1 «s 1Лг— cOQr^mi- fj»rtf-r— r— m С0Г-ФОГМ91СЧО1Г000ФО r-r—r—OOr—Or—OCVJr—t— Or—Or- OOOOOOOO evjr—cvjoor—Of—Or—Or— dw-^dr^-dddddddd СМЛСМШОг- Or- Of— r— r— OOOOOOOOOOOO dddddddddddd ooooooooooo +1 +1 +1 +1 +1 +1 +1 +1 +1 +1 +1 +1 dddddddddddd «Jf^OOOlCONOr- СГ1 •* r— +1 +1 +1 +1 -H +1 +1 -H +1 +1 +1 +1 r%f—r^OC\JLf)Cv»*OC4JCMCOO 8 r^r— CO O r— ONCnr>teON CMOCMOOr—Or— Or— Or— fr>r>OOr— Crt r— OOr- —OCMO Or—I Or— r— r— di-^dr^dddddddd dddr^dddddddd c ,зот1 oiSr- a O fo -r- ХГг— Ш O fflr-Tl C DUZOU.ZUlI"tW>M ; '. '. , '. '. '. '. '. '.«*cd • •, <\J CM II O f—C\J ГО *d" 1Г> 1П U>tONC0O\r- r— r— r— r- r— M :•»- The differential energy spectrum of oxygen TMlE J is determined from the observed number of SPECTRAL INDICES FOR C0SKIC RAYS oygen nuclei, the net exposure factor, and ABOVE S.< 6EV PER HUCIEOW* the absolute value of the correction factors. A power law fit tn total energy г Element 'Spectra l Index yields the spectral index of 2.67 : 0.04 for oxygen over the energy range 5.4 to 90 5 Boron 3.3S i 0.10 C-eV/n. To calculate the spectral indices 6 Carbon 2.76 s 0.06 of the other elements, a best power law 7 Nitrogen 3.07 t 0.10 fit was made to the abundances relative to 8 Oxygen 2.67 t 0.04 oxygen for the combined medium and high 9 Fluorine 3.01 i 0.21 energy data. The spectral Index differ 10 Neon 2.88 i 0.08 ences were then applied to the measured 11 Sodium 3.14 s 0.25 spectral index of oxygen. The results of 12 Magnesium 2.60 i 0.06 this calculation are shown in Table 3. 13 Aluminum 2.52 i 0.10 14 Silicon 2.63 ± 0.U 5. Discussion and Conclusions. The 15-24 VH Group 2.64 i 0.07 energy dependence o the cosmic-ray • 25-28 Iron Group. 2.50 i 0.08 leakage path length can be determined by *Spectra have been fitted to a relating the observed energy dependence power law in total energy. Only of B/0 to the predictions of a propa- statistical errors are shown, gation model for a variety of path lengths. Since boron is assumed to be completely secondary, the amount of matter traversal that the model requires to reproduce the measured B/0 at a given energy, yields the mean value of the amount of interstellar material through which the cosmic rays have propagated at that energy. A least-squares fit to our derived leakage path length yields X(E) = Etot ± • This is consistent with the result of Juliusson et al. (1975), who arrived at the value of -0.49 ± 0.05 for the spectral index of the leakage path length. This energy-dependent path length is incorporated into our propagation model. The assumption of energy-independent source composition is retained to test its validity. The propagation model is used to calculate the energy dependence of the elemental abundances relative to oxygen, and these predic tions are compared with our measurements. We find no significant disparities that would force us to reject the notion of energy-independent source composition for any of the relative abundance measurements reported in Tables 1 and 2. In Figures 3 and 4, we present our data and our propagation model predictions for C/0, C/S1, N/0, and 0/1ron group, as well as a compilation of recent measurements. To obviate the possibility of introducing systematic error by the use of differing atmospheric corrections, we have applied our-atmospheric corrections also to the data of other authors that are shown in these two figures. Acknowledgements. The authors would like to thank Dr. E1nar Juliusson for his work on the original version of the instrument, and for many valuable suggestions concerning this project. They are also greatly Indebted to the many members ot the staff of the Laboratory for Astrophysics and Space Research who contributed to this effort. The balloon flights were conducted by the National Scientific Balloon Facility of the National Center for Atmospheric Research. This work was supported in part by NASA under Grant NGL 14-001-005. References. Balasubrahmanyan, V.K., and Ormes, J.F. 1973, Ap. J., 186, 109. 24* Caldwttl. J.H. 1977. tubaltted for publication. Jullmton, E. 1974, Ap. J.. J9i. 331. Jultwsson. E.. and Ncytr, P. 1975, 14th ICRC (ИщИсМ. 1, 256. and private COMMltCttlOfl. JuHusson, E., Ntyer. P.. and Nulltr, D. 1972. Hurt. Rev. LttUrt. 28, 926. Jullusson, E.. Ctsarsky, C. NtntouKl. N.. and Cass*. N. 1975. 14thlCRC (Munich). 2. 653. Leznlak. J.A.. and Hebbtr. H.R. 1975. 14th IOC (Nunlch). 12., 4107. Lund. N.. Lundgaard Ra$«isstn, I., Peters, B., and Mestergaard, N.J. 1975, 14th ICRC (Nunlch). 1, 257. Ornes. J.F., Fisher, A.. Hagen, F., Nathl. R., and Arens. J. 1975. 14th ICRC (Nunlch). 1, 245. Orth, CO., Bufflngton, A., and Snoot, G.F. 1975, 14th ICRC (Nunlch), 1, 286. Sarith. L., Bufflngton, A.. Smoot. G., Alvarez, L., and Uahllg, N. 1973, Ap. J., 180. 987. т 1—i i i mil 1—n i mil TTT| 1 1 1 I f I I l| I I I" I I I I I o) * f 01 i I - i i" •e 5 го ПЙСЯГ J [•0t s -1 I I | I llll 1 I Ю ICO Kinetic Energy (GtWhucteonl MOM Fig. 3 Fig. 4 Relative "abundance measurements for C/0,. C/S1, N/0, and 0/1ron group. The solid lines are the propagation model predictions with energy-dependent leakage path length and energy-Independent source composition. :дч CHtNlCAL COMPOSITION OF COSMIC-RAY IRON CROUP NUCLEI AT ABOUT I GeV/Wuc Peter Meyer «rid Gary N1n«9«w* Enrico Fermi Institute, University of Chicago Chicago. Illinois 60637 Data from • high-resolution measurement of the abundance distribution of the cosm1c-ray Iron group nuclei (245 Z 5 28) меге gathered 1n two high-altitude balloon flights 1n the fall of 1975. The total exposure factor of the experiment exceeded 20 mz sr hr. The detector system consisted of two plastic scintillators, two Cerenkov counters {n • 1.27 and n * 1.49), and a six-plane multi- wire proportional counter hodoscope. In the energy Interval from 1300 to 2100 MeV/nucleon we find Cr/Mn/Fe/ N1 to be 12/11/100/4.5. 1. Introduction. - Knowledge of the elemental composition of the nu- clear components Is essential to the understanding of the origin and pro pagation of the cosmic radiation. Over the years many studies have covered most or all of the range of nuclei from Hydrogen to Nickel, but separation of individual elements in the iron group has not been achieved well since it is technically difficult. We have made two balloon flights with an instrument specifically designed to overcome this difficulty and to pro vide sufficiently high resolution to yield charge separation for the elements fromCarbonto Nickel, and accurate energy measurement in the range from 400 to 2500 MeV/nucleon. In this paper we present results on the relative com position of the iron group nuclei around 1 GeV/nucleon. 2. Instrumentation and Flights. The instrument used in this experi ment is a combination of two scintillation counters, two Cerenkov counters, and a multiwire proportional counter hodoscope similar to that described by Dwyer and Meyer (1975). Figure 1 shows a cross section of the detector system. The counters labeled Tl and T3 are identical 1 cm. thick Pilot Y CERENKOV COUNTER MWPC A scintillators, which are used to determine the charge of penetrating nuclei. TO is a plastic Cerenkov SCINTILLATION COUNTER counter with wave shifter (Pilot 425) having an index of refraction of 1.49. T2 is a 2 cm. thick liq LIQUID CERENKOV COUNTER T-2 uid Cerenkov counter with Freon MWHC В of index of refraction, n = 1.27. It, too, contains a waveshifter. SCINTILLATION COUNTER The two Cerenkov counters combine MWPC C to provide a contlnous and over lapping velocity measurement for particles from 400 to about 2500 MeV/nucleon. Each of the counters GUARD is enclosed in a white, light In tegration box to gain uniformity 2" P.M. TUBES 0 25 CM. of signals over the entire counter Figure 1. area. An annular scintillator, G, Instrument cross section. :5d surrounds the bottom of the detector stacK and serves to reject star *>••.«.-<-. The nulttwtre proportion*! counter hodoscope consists of three ch*r£cr-.. design»ted NWPC A, B, and C in Figure 1, each aith wire grids In the > anj • directions. The hodoscope determines the trajectory of the particle and atsu define* the geometrical factor of the instrument at 0.?5 r>? sr. Jt serves to correct the pulse heights from the counters for angle of Incidence and for small, residual nonuniform!ties. The data were gathered on two high-altitude balloon flights launched from Muskogee, Oklahoma (vertical rigidity cutoff of 3.4 GV) 1n September anj October, 1975. The first flight floated at 4 om cm"? and dipped to almost A gm cm"' on the last night; It lasted 60 hours. The second flight stayed at •! gm cm"' for all of Its 40 hours. The total exposure factor of this experi ment is greater than 20 m2 sr hr. 3. Data Analysis. An event, to be accepted, must satisfy several selection criteria. FTrst, the trajectory of the particle, as measured by the proportional counter hodoscope, must appear as one and only one straight line. Also, the event must have a low guard counter signal. Finally, the pulse heights from the scintillators, Tl and T3, must agree within narrow limits, which indicate that the particle has not interacted In the detector stack. The requirement that an event caused by a penetrating nucleus have a unique straight line not only makes pathlength corrections pos sible, but essentially removes all background. The percentage of particles removed by this criteri- oin on is independent of charge over the whole range of elements cover ed by this experiment. The limits for agreement of the scin tillator pulse heights are applied in a charge dependent manner. The window for agreement narrows Ft- If on Group quickly from Carbon to Silicon and remains constant from there. For 0.1SC2/C2m„«S0.8 higher charges, allowances must be made for particles which have sufficiently low energy to slow £ 120H down appreciably in the instru ment. For these events, charac terized by low Cerenkov signals, the window is adjusted to permit the acceptance of particles with slightly higher signals in the lower scintillator. To display the resolving power of the instrument Figure 2 shows a histogram of the charge function, i.e., the square root J of the sum of the signals from I. A.1 the scintillators, for the charge 25 26 interval from Titanium to Nickel, Charge and taken from a portion of the first flight. It is assembled Figure 2. under the criteria described a- Charge histogram of data from the bove, with agreement of the first flight with energy between 700 and 2000 MeV/nucleon at the Instrument. :5i sc'ntlllators set at i S.3S. These d«tt are selected fro* between 10 and 801 of the beta • 1 point (C2max) of the liquid-filled Cerenkov counter, corre sponding to an energy range fro* 700 to 2000 NtV/nucleon at the detector. The standard deviation of the charge function, a Measure of the resolution of the Instrument, Is 0.25 charge units, even for this broad energy range. Han- ganese, for the first tine at these energies. Is unambiguously resolved. After selection, the data are corrected for abundance overlap between neighboring charges, for Interactions within the instrument, and for propa gation through the atmosphere. Except for the elements next to Iron, namely Manganese and Cobalt, the overlap corrections are less than a few percent. 4. ___R*su1*si Table 1 gives the abundance of each element In the Iron group (24 i Z . 28) relative to Iron at the Instrument and at the top of the atmosphere. It also has the correction factor for Interactions in the Instru ment and the correction to the top of the atmosphere (per gm cm"' per 1000 Fe) to be added to the ratios. These data are selected from the second and part of the first flight In the energy range from 1300 to 2100 MeV/nucleon. Only statistical errors are given. These results are preliminary since no detailed evaluation of the systematic errors, due to the selection criteria and the corrections, has been made, and since only a portion of the available data has been used. Our results agree well with the work of Benegas et al. (1975) which was obtained at a slightly lower energy (870 - 1400 HeV/nuc.) and with the low energy results of Garda-Hunoz et a!. (1977) between 72 and 450 MeV/nuclecn. In particular, our measurement of the ratio Mn/Fe = 0.11 is consistent with those of Benegas et al., Garcia-Munoz et al., Lund et al. (1975), and Webber et al. (1972), but not with that of F1snir~et al. (T576T who find Mn/Fe = 0ТТУ+0.01. TABLE 1 IRON GROUP ABUNDANCES, 1300 < T < 2100 MeV/NUCLEON Atmospheric Instrument Correction Interaction (Additive) Top of the Observed Correction Factor per gm cm"z Atmosphere z Abundance (Mu1t1pl1c1t1ve) per 1000 Fe Abundance 24 0.189 + 0.013 0.981 -8.9 0.123 + 0.008 25 0.147 + 0.011 0.991 -5.3 0.109 + 0.008 26 1.0 +0.028 .1.0 0.0 1.0 + 0.028 27 0.011 + 0.003 1.014 -0.2 < 0.010 28 0.046 + 0.006 1.013 +0.2 0.047 + 0.006 252 5. Acknowledgements. Ue thank Dr. R. Deyer and Dr. J. Caldwell for major contributions to this project. Me are grateful to the »uff of the Laboratory for Astrophysics and Space Research, «specially W. Johnson, G. Kelderhouse. and M. Hollls, for the design, construction, and preparation for flights of the instrument. The balloon flight program MS ably handled by the staff of the National Scientific Balloon Facility. This work Is supported In part by the National Aeronautics and Space Administration under grant NGL 14-001-005. References Benegas, J. C., Israel, M. H., Klarmann, J., and Haehl, R. C., Conf. Papers, Uth International Cosmic Ray. Conf., 1. 251 (1975). Dmyer, R. and Meyer, P., Phys. Rev. Lett.. 35. 601 (1975). Fisher, A. J., Hagen, F. A., Haehl, R. C., Ormes, J. F., and Arens, J. F., Ap. J.. 205. 938 (1976). Lund, N.. Rasmussen, I. L., and Peters, B., Conf. Papers, 14th International Cosmtc Ray Conf., 1, 263 (1975). Webber, W. R., Damle. S.~V., and Klsh, J., Astrophys. and Space Sd., 15, 24S (1972). ~~ COfJCIC RAT ЛЯТГОАМСЯ ГНПН 'TTHDCTt: T) "Г-С 'JsTtfC A cftWLoet WITRAK PLASTIC D-:TX;TOR O.S»JCalntt>,T.3.Bhetla «d suaan Paruthl. Physics Departaent. Pan Jab University, Ch*ndlr«rh.i4 ( India.) Theoretical П lip»rUirt«® Bote Q *n atteapt baa been aada to study the chare* coaposltlon of low energy ooaate ray nuclei fro» ftltronen to zinc, A a tack of aaelear eaolsloas and Datcall Cellulose nitrate sheets azpoaad over Fart Churchill (Canada) la July 19«B has b*an need. After ohaaleal etching, about 6000 conical pita belonclng to about Ю00 stopping oo«ale ray tracks пате bean analysed. Iaproved procedures of aeasureaant and analysis yield a charge reaolotloa of 0*2 charge onlts for nuclei oato sulphur end 0*3 to 0*4 miti of charge for heavier elements. The present oxperlaent, therefore, shows that Cellulose Nitrate detectors can tie aged aueoessfufly for the study of cosale ray nuclei with eh areas ae high as 33* Ooaposltloa of nuclei with z»7,30 baaed on our •easuraaeotst «Ш be presented in the paper* Coordinates» OQ 1,5 (Nuclear Coaposltlon of Cosmic Rays) Hailing address* Dr.r.3.Baatla. Physics Departaeot, Pan jab Uafreraity. Chandlgaih-ld00l4 WDIA. :я VH COSMIC RAY MEASUREMENTS WITH A 6.6 n:sr ELECTRONIC DETECTOR J. Tueller, P. Love, J. W. Eptseln. M. H. Israel and J. Магдагт McDonnell Center for the Spice Sciences. Department of Physics. Washington University, St. Louis, Missouri 63130. U.S.A. Results on the relative abundances of Iron group elements 1n the Cosmic Radiation have been obtained from 28 hours of a balloon flight of 6.6 m sr electronic detector. Excellent resolution combined with a significant statistics yields the ratio Nl/Fe » 0.0491 0.003, Independent of energy 1n the Interval 600s Е s 1100 MeV/amu. We also place an upper limit on the abundance of Co as ' Co/Fe< 0.008, which Indicates a delay time of >1 year between the synthesis and the acceleration of Co. 1. Introduction. This detector was designed and built at Washington ilnlversity to measure charge abundances for cosmic rays with Z > 12 with higher precision than hitherto possible. The improvements over our previous detector system (Benegas et al. 1975) involve a) an increase in geometry factor from 0.9 m*sr to 6.6 m2sr, b) an Improvement in hodoscope resolution from 4.8 cm to 2.5 cm, c) an increase in the depth of the 1on chambers from 30 cm to 40 cm. We report here results on the composition of the Fe group in the energy range 600 2.3 cm -II- "(IONIZATION HODOSCOPE UPPER ALUMINIZED - [IONIZATION CHAMBER #1 MYLAR <= ELECTRODES IONIZATION CHAMBER «2 t=J25cm WHITE OORVON CHERENKOV COUNTER <гш @ @ © @ f=^95cm PILOT 425 RCA 4525 ' ^ [IONIZATION CHAMBER *3 3 [IONIZATION CHAMBER*4 [[IONIZATION HODOSCOPE LOWER - 165cm- -I Figure 1. Schematic of Detector Module. 255 Corrections to tht ran pulst heights for time variation (due to temper- «ture changes) and for are* non-uniformities «re generated by selecting appropriate sets of Iron data. In noreal analysis ме require that no Ion cheater pulse height differ from the mean by no re than 1 charge unit. This criterion eliminates -HOY of events In agreement with the expected number of Interactions and statistical fluctuations. For highest resolution (as for the determination of the Co abundance) we "«quire agreement to within 0.5 charge units. This stricter requirement rejects about half of the particles but results In slightly improved charge resolution. The Nl/Fe ratio 1s the same for both criteria but the relative abundance of particles with charge assignments near Z»27 Is significantly reduced. 3. Results. Figure 2 Is a charge histogram with the 0.5 charge unit criterion of the Fe region for events with energy 400< Е < 1000 MeV/amu at the detector. Charge assignment 1s based on the empirical Ionization vs Cherenkov curve obtained for Fe in flight. The curves drawn for Z» 24,25,28 are scaled from the Fe distribution to fit at the peak channels. We obtain a ratio Ni/Fe « 0.049 ± 0.003. Our large statistics enable us to divide the data Into energy bins and obtain the ratio Nl/Fe as a function of energy as shown in Figure 3. There is no evidence for a difference in the energy spectra as has been suggested (Lund 1975). 400 600 800 ЮОО ENERGY (MtV/omu) Figure 2. Charge Histogram at Fe Figure 3. N1/Fe Ratio vs Energy for 400 Figure 4 is an expanded view of the Co region in Figure 2. There clearly is no peak visible at Z = 27 and many of the particles assigned Zs27 are due to Fe nuclei. As a result we cannot derive a flux for Co. Instead, we derive a conservative upper limit to the Co abundance by fitting a scaled Fe peal: to the observed abundance at Z = 27, thus neglecting the Fe contribu tion entirely. This procedure results in a ratio Co/Fe « 0.0074 ± 0.0014. Taking 0.009 as the upper limit and extrapolating we obtain an upper limit for Co/Fe of 0.008 at the top of the atmosphere. 256 4. Conclusion. Flour* S shows curves taken fro» Soutoul et »1. (Ш5). the variation of the N1 and the Co abundance are shown at • function of deity tie* between synthesis and acceleration. The reason for variation with time Is that both N1-56 and Co-57 decay by K capture only. Thus they must survive once the nuclei have lost their orbital electrons during acceleration. The solid curve for Co assumes Co-59 at the source, the dashed line assumes no Co-59 at the source. ] i Figure 5. N1, Fe, and Co abundances vs Delay Time Between 27 26 Synthesis and Acceleration. CHARGE Figure 4. Expansion of Figure 2. Our Co data are only marginally consistent with the presence of Co-59 at the source, and require a delay 2l year between synthesis and acceleration. The N1 abundance alone places a lower limit of 100 days on this time. Our N1 abundance,0.049±0.003,1s slightly smaller than the Cameron (1973) solar system abundance relative to Fe (0.058). While the difference Is small, It Is substantially larger than the uncertainty In our measurement. Acknowledgements. We thank RaVen Industries for their successful balloon flight operations. This work was supported in part by NASA Grant NGR 26- 008-001. Re.erences Benegas, J. C, H. H. Israel, J. Klarmann, R. C. Naehl, 1975; 14th ICRC, Munich I, 251. Caeeron, A. 6. U.. 1973; Space Scl. Rev. 15, 121. iS7 Lovt, P. L. J. Twllir, J. N. Epstein. N. H. UrM). J. tliTMM. H77i nm u$. 5*». Lund, N.. I. L. RtMMttm, I. Httrt, N. totMbtrf. N. J. tftsUrttN, 1)75; 14th IOC. Munich 1. 2Ю. Soutixil. A.. N. Castt. E. Jullusson, 1975; Uth ICRC. Nuntch I. 455. IS* ULTRA HEAVY COSMIC МГ NEASUROCNT3 WITH A t.6 e'jr tUCTftCMK DCUCTC* • . Love, J. Tueller, J. W. Epstein, N. H. Israel end J, lUnw» McDonnell Center for tht Sp«ct Sciences. OeptrtiMpt et 4mytlci Mashlngton University. St. Loub. Missouri «150, U.S.A. A 6.6 m %r electronic detector was employed on • 28 hours balloon flight to obtain abundances of nuclei of Zi 30 In the cotalc rays. Me obtain • Zn abundance of Zn/Fe • (7 i 2) x 10"\ The abundances relative to Iron of nuclei with 30 *Z $40 are lower than In the solar system abundances of Cameron (1973), but Individual elements do seta to follow the trend of the solar system. 1. Introduction. Measurements of abundances of nuclei of 2г30 is hampered by extremely low fluxes and hence require large exposure factors. Light weight passive detectors have been extensively used, but do not exhibit single charge resolution (Shirk et al. 1973). Electronic detectors afford much better resolution but have so far yielded only limited exposures (Jullusson et al. 1975, B1nns et al. 1973). He have constructed a large area electronic detector system (see paper OG-71, this conference) which has been flown successfully for 60 hours at the time of this writing. Me report here preliminary results obtained on the first 28 hour exposure at 5.5 g/cm1 float altitude. 2. Charge Resolution. Figure T shows a charge histogram for Z<30 for events of 400 i Е si 000 MeV/amu at the detector. This is a small sample of our data to demonstrate the charge resolution at low energies. Charge has been assigned by normalizing to the Fe peak and assuming a Z* dependence. The SI peak is observed to fall at Z • 14.00 ± 0.01, showing no deviation from linearity in the region 14sZs26. The sigma at the Si peak 1s 0.23 charge units, similar to ths sigma observed at Fe (0.25 charge units). Clearly resolved peaks at every unit charge demonstrate our single charge resolution. i П " 1 1 20 10 14. 16 20 22 24 26 2 Z3 30 CHARGE CHARGE Figure 1. Charge Histogram for 400< Е F1 2. Charge Histogram near < 1000 MeV/amu at the Detector. Z-30 for 400<Е < 800 MeV/amu. :s« Flgur* 2 shows th* region around Zn for 400 « £ « 800 NrV/*au *t tfte detector. Th* line corresponds to • scaling of th* Fe peak yielding Zn/Fe • (7*{)x10~V Th* resolution based on th* F* peak IMM to fit th* date «ill «ltd th* peak 1s clearly resolved. Flgur* 3. • theoretical plot of Ionization vs Ch*r*nkov shows that lh* charge becomes a Multivalued function of C and I In tht high energy region, и* hav* assigned charges on th* assumption of no r*lat1vlstie rise and plotted the resulting charge histogram for Е >1 G*V/aau In Flgur* 4. It ts seen that Ignoring the relativists rise results 1n a tall towards highw charges, such that N1 can no longer be resolved fro* Ft but Nn can be re- solved. Since aaong th* high energy data we do not observe any events of 2«33 we are confident that this "relatlvlstlc tall* does not contaminate the data of Z>32. Furthermore H Is apparent fro* Figure 4 that our charge assignment ignoring relatlvlstlc rise will be correct for the vast Majority of all events. We therefore Included data for energy Е >1000 HeV/amu In our results for Z>32. 5000 3000 2000 1000 0.4 -0.6 08 10 \Z 26 • 27 SOUARE ROOT CHERENKOV CHARGE Figure 3. Theoretical /Tvs /TPlot. Figure 4. Charge Histogram for Е>1000 rteV/amu at the Detector. Our statistics do not allow us to make Individual charge peaks and thus prove our resolution. However we have plotted (Figure 5) the fractional charge for all our data of Z*32. This distribution is consistent with the observed Iron resolution (o » 0.25 charge units) within our meager statistics. 3. Results. Figure 6 shows the abundances of all charges greater than Fe normalized to Fe * 10*. For Z>32 the high energy data are included. The background histogram is the Cameron (1973) solar system abundances. The lowest experimental data points correspond to one observed event, elements with no data points were not found. Our results are also presented In Table 1 In comparison to the solar system abundances and to the 1973 summary of ultra heavy nuclei (Fowler 1973), (abundnaces normalized to Fe - 1.00). 260 -i—i—i—r—|—i i i—r—r * ' i~~i r~~\ •I • I • I • * * * * • •.i I.I i.« t.t •.« so м 40 «a 4o ss to rMCTMML I NUCLEAR CKAMC Figure 5. Histogram of Fractional Figure 6. Abundances of Elewnts Charge Values for Z > 32. * with Z>26. Fe « 10*. (The histogram Is the Cameron (1973) Solar System Abundances.) TABLE 1 Charge This Experlaent Caatron (1973) Fwlcr (1973) 27 <8 x 10"3 2.7 x 10"3 28 (4.9 ±0.3)x10*2 5.8 x 10"2 4 29 <8 x 10"* 6.5 x 10" 30 (7 ± 2) x 10"4 15.0 x 10"4 32 (1.0 ± 0.5) xlO"4 . 1.4 x 10"4 34-40 (1.0 ± 0.3) XlO"4 2.3 x 10"4 0.65 x 10"4 (35-39) 41-49 <3 x 10~5 1.3 x 10"8 4 X 10"S (40-49) 50-54 (3.5 * 1.7)x10"S 2.0 x 10"5 0.8 x 10"5 > 55 <3 x 10"5 2.1 x 10"5 2.3 x 10"5 The most striking, and statistically significant result Is the under-: abundance of Zn In the cosmic rays which Is only one half of the solar systbm abundance. In the whole region 30 tnowlidgemnts. «to thank Raven Industries for their successful balloon Йlen t operations. This uort wet supported In part by NASA Grant HSR M-OM-OOI. References Huns. V. R.. J. I. Fernandez. M. H. Israel, J. Klamann. R. C. Naenl. R. A. Newaldt. 1973; 13th ICRC. Denver. 1_. 260. Caaeron. A. 6. M.. 1973; Space Set. Rev. 15., 121. Fowltr. P. H.. 1973; 13th ICRC. Denver, 5. 3627. Jullutson. E.. P. Meyer. 1975; Ap. J. 201. 76. Shirk. E. K., P. B. Prlct, E. J. Kobttlch, M. Z. Osborne, L. S. Plnstl. R. 0. Eandl. R. B.pushing. 1973; Phvs. Rev. D 7. 3220. гь: AVEKAGE AN'RUANCES СГ CALAvTK CUIC НАУ5 '4 I '.".•: .' > KROK STUDIKS ок иъткок:::: L.'L:YIK.*.S V.l'.Terel Vgin» S.u.Sti'U.i-nk^, J. L!..-ifv.4 • ,.r<-1:, v . •• '. • . < :. .'eutit Institute for N'uc-lu.'ii- fini'.'i"."., 'J.b;. , - -*.i r.Pollas, LaboruU'i у oi" Mi norali r v , !.,::•, >:.•..• B.Jukupi, University of Priatiim, Pr.r. • . •. . , Y .r- ::l.ivi . Abstract By applying the ;H'luc'„. v'c i.tct.iii(.' teol.:...... ;.. from Marjalahti and Eagle Station moteoiutet; iv. :• • •• > • :-,к-k_ longer than 1^0 um have been found to be due to (.viliiciio cr:£ic ray nuclei with Z>50. The data on the track distribution are compared with the known values of the relative abundances of the 50 ^Z 6-92 elements in the solar system. Satisfactory agreement is obtained assuming the group of the longest, 720-900 um tracks to correspond to the Th-U nuclei. I.INTRODUCTION Studies of the atomic number distribution and energy spectra of the WH component of galactic cosmic rays are extremely comp- licated due to the low intensity of this component, about 4x10 of the intensity of the Fe group nuclei. Systematic studies of WH nuclei were initiated in 1967 by 1} Fowler et al •, who carried out experiments usine nuclear emul sion stacks.exposed in balloons in the upper layers of atmosphe re. Since 1969 similar studies have been performed by the groups of Fowler, Price, Walker and Fleischer with nuclear emulsions, polymer foils and Cherenkov detectors '. However the total exposure time for the 10 years of these intensive studies does not exceed 3-Ц- m year sr. To increase considerably the sensitivity of such investigations one should expose in space for a long time tens and hundreds of square me-r ters of nuclear track detectors. Another method of investigating the WH component of cosmic rays makes use of the capability of extraterrestrial silicate minerals to detect and record for several tens and hundreds of million years the tracks of galactic nuclei with Z>20 JJ, гь.» Simple estimations she» th.->t one cuMc '<• ,'ite crystals from the furfnee ltsvera iV the c.< in exposure n^e of 10'-10 .vciir:; cont.-il-i:, !u jf tracks due to the thorium-uranium nuelii •) previous studies of hyperstenes frcn Johnstown c ц К;пьг pigeonite only three tracks of :',>, 0 r.i '.IV. ui.-ntif ied. However the tot;il number of the vi'..:i ги -isured in refs. Ъ' b)' is several tim»s cniiilei- t( of the tracks revealed in ixperiments Ьн:н-Ц on tlie i.:-<-^- :•- ding of galactic cosmic ray nuclei, 2 I.STUDIES OF THE TRACKS OF WH NUCLEI IN OLIVINES Here we present the results of our studies of the tracks of WH nuclei in silicate crystals from meteorites, obtained 4) since 1974 » The aim of these experiments has been to increase significantly the experimental sensitivity to the WH component of galactic cosmic rays. The studies were carried out using transparent homogeneous crystals of olivines from pallasites. Olivine makes ?0% of the pallasite volume, the-crystal size usually amounting to several millimeters. At the first stage of these studies a total of over twenty pallasites have been examined and'a few samples located very closely to the preatmospheric surface of these meteorites have 7-101 been found In these la'tter papers it has been shown that Lipovsky, Marjalahti, Pavlodar and Easle Station meteorites have ft o locations with the density of Fe group tracks of 10 -1-5x10' per 2 cm of olivine crystals. The tetudies of the tracks due to WH nuclei required a sig nificant improvement 10) of the etching technique proposed by ,11) Krishnaswami et al , the unambiguous identification of these tracks against the background of dislocations and capillar • inc- 12") lusions , the calibration of meteoritic olivines with accele- 10 П) rated ions ranging from Tl to Кг and Xe ' » and the detailed studies of the thermal fading of old and new tracks in these 8 4) crystals »7 . Our previous studies of the WH tracks detected in 62 mnr of olivines from the surface layers of Marjalahti meteorite» 2М carried out la 1976» permitted the doterolnot,Ion оГ the rrlativ- abundances of the Z >50 nuclei, averaged ovor the erp-osure afe 101 of this meteorite equal to 175 million y?nra '. The f. resent investigations of olivines from Earjalahti and Eagle Station meteorites are an extension of these studies. We have selected, mounted in epoxy and polished over 600 crystals & 2 na in oirt of olivine from Marjelahti meteorite. The crystals had the den sity of the VH nuclei tracks of ('*-6)x10 per cm2. We have blue examined eeveral tens of >1 mm crystai3 of Eagle Station mete- orite with е density of VH nuclei tracks of (5-15)x10 per cm1". The etching of these crystals was carried out in hermeti cally sealed ampules at 100-110°C for a period of several days (refs. ' J). In the crystal volume tracks of WH nuclei were revealed in natural cracks, capillar inclusions and dislocations 10") inaide the crystals . Olivines without such structural de fects were irradiated perpendicularly to their surface with 40 4-0 7.3 MeV/nucl xenon ions, or with Ar and Ca ions with a flu- 14 -2 ence of 10 cm end energy of 300 MeV through several slits with a 15-20um width and 150-200 urn spacine. Such a procedure per mitted a rather efficient etching of WH tracks inside the crys tal volume. The microscopic examination of the crystals and measurement of the tracks were done at a magnification factor of 500-1500 X. All the tracks longer than 40 Jim. were measured. To revea"1 the total etchable length of the inner tracks, the repeated etching 10) of the crystals was used . III.RESULTS AND DISCUSSION FiK. 1 shows the length distribution of the 1905 tracks measured in 100 пшг of olivines from the surface parts of Mar.ia- lahti meteorite. The distribution of the lengths of the 370 tracks revealed in a the 16шпг volume of olivine from Basle Station is presented in fig. 2. It should be noted that in the region of L > 300 jum about 50% of the tracks went beyond the crystal edges. As follows from figs. 1 and 2, the track length distributions of cosmic ray nuclei in olivines from Marjalahti and Eagle Station meteorites (having exposure ages of 175 and 50x10 yr, respectively) have similar shapes. In order to US increase, the statistics for the heaviest component of the YVH nuclei »e hsve examined en additional amount оГ 180 ma of Мл: . lahti olivine? *ith • track density of *-5x10 per cm". In tiau case onlj the trecks longer than 260 |im were meuaurod. *4g. 3 shows the result* of this examination together with all dam obtained previously. As follows from fig. 3, the total llnrjulah- ti olivine crystal volume examined amounts to 260 mm , and the number of the measured tracks of L>260^m is 1^ьч. M MARJALAHTl OLIVINES 120 INI» VaOtW? Г eo 40 /. -~ . 1 . ^*U iyi««a 200 400 600 eoo 1000 Fig. 1. т 1 1 1 r Eagle Station' !N-370tr V-16mm3 400 eoo L Aim Fie. 2. :бб 1*0 MAfUALAHTI OLIVINC4 «о •0 40 0 2CO The main difficulty in decipherine the measured spectra of the track lengths is the establishment of an Unambiguous cor respondence of the etchable track lengths to the atomic numbers 101 Z of WH nuclei. In ref. ' two versions of deciphering track length spectra have been proposed; first, by a semiempirical ex trapolation of the data on the etchable length of the tracks due to the accelerated nuclei ranging from Ti to Кг into the Z region of 92 *»°'. In this approximation (taking into account the partial fading. of the tracks of WH nuclei) the track lengths of L> 560 yum correspond to the region of Z^86. The total number of such tracks is 275. If °ne takes into account the processes of slowing down and fragmentation in meteoritic matter, the abun dance of the Z>86 nuclei relative to the Fe nuclei is 9.7x10"', 1 -2) which is in agreement with the data obtained by Fowler et al ' According to another "assumption ^that -only L?720iim tracks are due to the U group nuclei one can compare the relative abun- dances of different groups of WH nuclei and elements of the so lar system "f'). The results of this analysis shown in table I indicate that the abundances of WH nuclei and the elements of the solar system agree within a factor of 0.36-2^5. In terms of this extrapolation only one track longer than 1110 jam (fiK., 3) J67 ia due to a 2 ^97 r.ucleua (the low-energy рогЧол of the track •sent beyond the crystal boundaries). The absence of trnckn loit er than 1.* mm in the spectra shown in figs. 1-3 nllo»a one to set the upper licit on the abunlance of hypothetical superheavy elements (Z>110) in galactic cosmic rays ct 2-.Sx10~" of that of 10) the VH nuclei. Second, in agreement with ref. ' it is possible to set the upper limit on the flux of magnetic monopoles of G = n(137 e/2 ) in the region of n£U at ЗхЮ"!9 CD"2S sr. TABLE I THE ABUNDANCE OF COSMIC RAY NUCLEI OF Z >50. Charge Track lengths, Number Abundance of Abundance in^-j groups urn tracks cosmic ray nucl.solar system ' NZi/20 "* z *3° Nzi/2° * z *30 50 References A. P.H.Fowler, R.A.Adams, V.G.Cowen, T.M.Kidd. Froc.Roy.Soc. A301,39(1967). 2. R.L.Fleischer, P.B.Price, R.NL..Walker. Nuclear Tracks in Solids. University of California Press, Berkeley, 1975. 3. M.Maurette, P.Pellas, R.M.Walker. Nature, 204,821(1964). 4. O.Otgonsuren, V.P.Perelygin. At.En. 37,164(1974-). 5. M.Maurette, P.Ihro, R.M.Walker, R.Webbik. Meteorite Res. 12, 286(1968). 6. P.B.Price, R.Rajan, E.K.Shirk, Proc. the II Lunar Sci.Conf. 2,V 3(2621) MIT Press (1971). 7.3f.P.Perelygin, H.B.Wiik, O.Otgonsuren. JINR Preprint P13-8359, Bubna, 1974. 8. G.N.Flerov, O.Otgonsuren, V.P.Perelygin. Izv.AN USSR,ser.fiz. 39, Я2, 388(1975). • 9. G.N.Flerov, T.P.Zholud, O.Otgonsuren, V.P.Perelygin, H.B.Wiik. Geoch.Cosmochim.Acta, 40, 305(1976). 10. O.Otgonsuren, V.P.Perelygin, 'P.Pellas, S.G.Stetsenko, N.N.Gav- rilova, C.Fieni.Astrophys. J. 210, 253(1976). 11. S.Krishnaswami, D.Lai, W.Prabhu, A.S.Tamhane. Science, 174, 287(1971). 12.,G.I. olivo-Dobrovolskaya, V.D.Kolomensky, N.N.Gavrilova, V.P.Perelygin, S.G.Stetsenko. Geochiaiya №10, 1476(1976). 13. P.B.Price, D.Lai, A.S.Tamhane, V.P.Perelygin. Earth and Planetary Sci. Lett. 19, 377(1973). 14. A.G.W.Cameron. Space Sci.Rev. 15, 121(1973). 26* UN Cosmic toy»: Possible Origin In Massive SUrs John P. Hofol* and David N. Schramm Enrico Fermi Institute. University of Chicago Chicago, Illinois 60(37 luSA) and J. Bernard llikt The Aerospace Corporation Los Angeles. California S0009 (USA) The origin of the Z > 28, ultra-heavy, cosalc rays In supernova explosions of Massive stars, N i lOp, IS considered. For Z > 70, the UH data Is dominated by an r-process source distribution, but for the elements just beyond Iron, 29 <. Z < 36, the data can not be explained by any single process of nucleosynthesis. This problea Is solved naturally In a massive star eedel by secondary neutron-capture reactions occurlng during core helluai burning (a United s-proeess) and during explosive carbon burning. Inter stellar propagation calculations have been performed with these episodes of synthesis as source distributions, and the results offer an explanation for the current UH cosmic-ray data. Further, the heavy element synthesis during explosive carbon burning Is re examined using more realistic Initial conditions given by the post-helium-burnlng configuration of the star. These results are compared to earlier work and to the UH cosmic-ray data. Some effects of preferential acceleration, based upon Ionization poten tial, are considered, and experimental tests for this model are discussed. 1. Introduction. UH cosmic rays (Z > 28) can provide unique tests for many ideas aoout heavy element nucleosynthesis and cosmic-ray acceleration and propagation. Extensive data has been obtained from plastic/emulsion detectors on long duration balloon flights (Fowler et al. 1976a; Fowler 1973 and refer ences therein; and on-board the Skylab space station (Price and Shirk 197S), and from electronic counter-telescopes (Binns et al. 1973; Jullusson and Meyer 1975). The measured charge spectrum for Z > 70 Is inconsistent with solar sys tem (Cameron 1973) relative abundances but, rather, shows a distribution char acteristic of the.rapid (r) neutron capture process of nucleosynthesis (Schrama 1972; Blake et aK 1977). However, for the lighter UH nuclei, Z < 60, the In terpretation is more complex. Figure 1 shows a compilation of UH measurements - plotted over the solar system (SS) abundances and normalized to iron. The data . points represent abundances sunned over the Indicated charge intervals (• from l Jullusson and Meyer 1975; t from Binns et al. 1973; a from Blanford et al. 1973a. b; v from Shirk at al. 1973; • fromTowler et al. 1976b; a compiled re- suits given by Fowler 197377 Although there are inconsistencies between varl- i sus experiments, the data generally follows the SS distribution. Below Z * 40, | the SS .elements are believed to be produced predominantly by a slow (s) neutrons apture process with only a small r-process component. Thus, the general slml-Щ larlty between the -osmlc ray and the SS measurements argues for a similar щ origin. m However, the astrophyslcal site for the formation of the SS elements just jf beyond iron,.29 <. Z < 36, is Itself unknown. Current models for the synthesis W •Robert R. McCormick Fellow 1 169 of tht bulk of the s-process fit ments through lead—helium shell UH COSMIC RAYS flashes In red giant stars with and M < 10 Mg--underproduce the ele SOUR SYSTEM ABUNDANCES ments just beyond Iron (Ulrlch 1973). Furthermore, the r-process forms few nuclei below mass 78 In current calculations and cannot explain the neutron-rich nuclei In this Interval. For Z > 36. the current models for s- and r- tt process nucleosynthesis repro ftfei" duce the SS abundances nicely> Щ but for the elements just beyond m Tt iron additional sites/processes rf- are needed. Perhaps cosmic-ray M\ Ln data can provide information on these sites. ~ 40 « 50 60 NUCLEAR CHARGE In this paper we consider a possible origin for the UH cosmic Figure 1 rays and the SS elements just be yond iron in supernova explosions of stars with M > 10 Kg. These massive stars are a plausible source since they are believed to be (1) the source of the ele ments carbon through iron via explosive nucleosynthesis 1na supernova, (2) the site of r-process synthesis, and (3) the birthplace of the cosmic rays. Two episodes of secondary neutron-capture nucleosynthesis peculiar to massive stars have been investigated. The helium-burning s-process (HBS-process) forms heavy elements through slow neutron captures on seed nuclei during hydrostatic core helium burning. The supply of neutrons for the HBS-process is limited, and, consequently, only isotopes below A \ 90 are formed (Couch et al. 1974). The second type of heavy element synthesis (the ECB-explosive carbon burning pro cess) occurs during the supernova explosion. A shell composed principally of carbon and oxygen—the residues of helium burning—is shock-heated to tempera tures above 2 x 109 °K and undergoes explosive processing. Large quantities of neutrons are liberated, some of which are captured by 'seed'nuclei 1n the shell forming neutron-rich Isotopes up to A-76 (Howard et al. 1972). \EC8-««iv. 4, Figure 2 shows the relative abundances of the isotopes, plot ted against the mass number, for the three processes of nucleosyn thesis considered. The abundances are normalized to the solar sys *щ tem value for the nuclide Indi it—*- cated 1n the lower left. The r- rf MASS WMKR MISS MJMKR process represented the SS r-pro V i • i i 1 ' cess component obtained by de R-wcass "иИ * • composing the Cameron (1973) abundance table for A i 78, and 0 the HBS data was obtained from - У ^у. Lamb et al. (1976). Cosm1c-ray I propagation calculations have i, ' i • A • if • k • k • Jo been performed for each of these distributions and the results Figure 2 are discussed below. 270 2; Propagation Calculations. The source distributions shown on Figure 2, assuring no Modification or selection effects In the acceleration process, were used as Input to a fully developed propagation calculation which followed all of the stable or long-lived isotopes between 5"Fe and 1><2Ce. Fragmenta tion cross sections were obtained from the semi-empirical equations of Sllber- berg and Tsao (1973). A power law 1n total energy per nucleon was assumed as the source spectral form, and the cosmic rays were propagated through an ex ponential path-length distribution with a mean of 6 g/cnr in an Interstellar medium of density • 0.7S atoms/cc composed of H, He, and heavier elements. Only energies greater than 500-600 MeV/nucleon at Earth we<*e considered. Define the cosmic-ray source distribution to be the superposition, n Source - [ Fe"]+«R[R] + «S[HBS] + «C[ECB] where R, HBS, and ECB refer to the relative abundances shown on figure 2, HFe" Is the solar system abundance of the Iron-peak elements and «R, «5 and «c are the scaling or normalization factors for the three processes. The parameters 6{ give, effectively, the amount of material, from each nucleosynthesis epi sode, which is mixed and accelerated with the Iron peak elements. The value of 6R was determined by the measured charge spectrum for Z > 45 since, in this model, the r-process 1s the sole source of these cosmic rays, 'pper limits for 6S and SQ were obtained from the cosmic ray source ratios C+O/Fe and Mg/Fe respectively (Shapiro and Silberberg 1974), by assuming that all of the cosmic ray C and 0 derives from helium-burning regions and the Hg is formed during explosive carbon burning. These limits are 65 £ 1-U and SC *£ 1>0 for the normalization employed on Figure 2. T I I I I I 1 ) I I ГТ"| I 1 I I I I I I 1 [ 1 I I I I F T I 1 I Figure 3 compares the calcu 10 ^R+HBS*ECB lations to the data, plotted as the o ratio of the data to the calcula tions versus nuclear charge. A 5 t value of <5R * 0.6 ± 0.2 brings the 5 iL тим ^ITTT _d>_i .• j calculations into accord with the 3 'ГТЧ^Г •H^*FП^Г^ - 1 data for 45 <. Z £ 55 and implies that the r-process peak at A * 130 is under-represented in the cosmic V0J6 5,=U7 V&5 rays relative to SS material. For OJ I 1 1 1 1 I 1 1 1 1 I 1 1 1 1 I 1 1 1 1 I 1 1 1 1 I 1 1 1 1 I the HBS-process the maximum nor 30 35 40 45 50 55 60 malization value is needed to pro NUCLEAR CHARGE vide sufficient abundance to ex plain the measurements for 36 i Z .Figure 3 < 40. A more stringent test would be provided by the interval 32 <. Z < 36, but here the data is contradictory. Finally, a value of Sr = 0.5 gives agreement with the Zn measurements of Juliusson and Meyer (1975) and Shirk et al. (1973), but the calculations are lower than the Zn measurements of Blnns et al. (1973) and Fowler et al.(1976b) and all of the Cu results. Clearly, Improved data for Cu-Se would permit a better determination of 6$ and &Q. There Is general agreement between the cosmic-ray data and the calcula tions on Figure-3-except for the charge intervals 40-44 and 55-59 which are not reproduced by the present model. Focusing on the data summarized by Fowler (1973), the calculated values are about 2 and 5 standard deviations, respectively, below the observations. A fission origin, from elements formed in the r-process at A ^ 280, has been suggested for these discrepant inter- \ Is (Blake and Schramm 1974; Wefel et al. 1977), but other possible explana tions are discussed in subsequent sections of this paper. Additional 271 Information on the nucleosynthesis end cosmic-ray propagation and a detailed Interpretation of the results from this model for the UH cosalc rays art given by Wefel et al. (1977) and references therein. 3. Explosive Carbon Burning. The relative abundances on Figure 2 for the ECB-process were obtained with a SS 'seed' distribution, because these calculations were performed before the importance of the HBS-process was real ized. The initial configuration of the region undergoing ECB-rtactlons fs determined by the previous helium burning period, and the HBS-process modifies these conditions in the following manner: (a) the neutron supply is reduced because neutrons are removed by the s-process, and (b) the 'seed' distribution is radically different with considerable abundance located above the Iron peak, as illustrated on Figure 4. Noti the sharp Iron peak region In the pre-H8S distribution which is reduced 1n size and widened follow ing HBS processing. The post-HBS curve also shows en hancements by * 100 In the mass range 80-90. The heavy element yield 1n the ECB-process has been re-calculated by using the neutron flux obtained from carbon burning calculations and the post-HBS 'seed' dis tribution as input to a generalized n(neutron)-process calculation (Blake and Schramm 1976). The time evolu tion of the 'seed' isotopes was followed—specifically the competition between neutron capture reactions and beta decay. The preliminary results for a peak tempera ture of 2.1 x 109 °K are shown on Figure 5, compared to the earlier calculations at temperatures of 2.05 and2.IS billion degrees. The present results show a lower pro duction in the mass range 65-67, due to the reduced num ber of iron peak 'seeds', but for A K 68 the results fall above the previous calculations. The odd mass Isotopes are underproduced on Figure 5 because (p,n) reactions Fiqure 4 have not yet been 1nclude<* in tne calculations. The importance of these results for the UH cosmic rays lies in the fact that the distribution on Figure 5 extends to A * 100. With these nuclei In the cosmic-ray source, the value of 6$ can be reduced, but more significantly an additional source of nuclei in the charge interval 40-44 is obtained. Using the 6n- factors on Figure 3, the calculated abundance for Z * 40-44 now falls within one standard devia tion of the experimental datal Thus, the ECB-process,, using a realistic initial configuration, plays potentially a major role In the nucleosynthesis of the UH cosmic rays.. 4; Acceleration Effects. The previous discussion assumed that there is no modification of к>'г the relative abundances between т9 T9-2.05' nucleosynthesis and acceleration T» > 2.10 -Thit Work to cosmic ray energies. This as ub U_i 70 75 80 85 sumption has been questioned Mass Number (Casse et al. 1975 and refer ences therein) since injection Figure 5 272 of particles Into the acceleration process may depend upon the atomic proper ties of the elements. Casse et al. present data which can be Interpreted 1n terms of a reduction In the source abundance (relative to SS material) with Increasing values of the first Ionization potential. Such a preferential ac celeration yields several noticeable effects for the UH cosmic rays. In par ticular, the r-process peak at A * 130 Is composed principally of the elements Te and Xe which have large Ionization potentials {•*• 9 and 12 eV respectively). Fitting an exponential to the data given by Casse et al. yields a rela tion between the 'enhancement' 1n the source abundances and the Ionization potential. Relative to Iron, the enhancement for the elements Te +'Xe 1s » 0.5. This 1s comparable to the normalization factor «R • 0.6 t 0.2 found In section 2 for the r-process. Thus, preferential acceleration appears to explain the apparent underabundance of the A * 130 r-process peak- 1n the cosmic rays relative to SS material. For the heavier elements Ba and La, an enhancement of * 2 relative to Iron 1s predicted, and this enhancement can reduce the disagreement between the calculated and measured abundances shown on Figure 3 for the Interval Z » SS-S9. There Is essentially no enhancement predicted for elements in the range Z - 40-44, but effects can be anticipated for the elements Se, Kr, and Sr. The question of preferential acceleration Is an Intriguing one, since it produces modifications 1n the UH source abundances and clouds any direct in terpretation of the data in terms of nucleosynthesis. It Is interesting to note that preferential acceleration tends to Improve the agreement in the Z > SO region which Is due solely to the r-process. Clearly, additional work is needed to resolve these questions. 5^ Conclusions. The episodes of nucleosynthesis occurring 1n massive stars prior to or during a supernova explosion provide a natural explanation for the UH cosmic rays. The combination of the r-process, the HBS-process, and ECB-synthes1s can explain the current data. The elements just beyond iron, Cu-Zr, are especially Important because here the various processes over lap. The ECB-process appears to be a major component of heavy element syn thesis and detailed nucleosynthesis calculations for this episode are needed. Preferential acceleration may play a role in the UH charge spectrum, but ad ditional work Is required to establish Its Importance. However, the best Indicator of the validity of this model will be Improved cosmic ray measure ments which should be forthcoming 1n the next several years. 6_j Acknowledgments. Thanks are due to J. D. Anglin for developing the propagation code ana to J. N.. Truran for providing some of his calculations prior to publication. This work was supported in part at the University of Chicago by NSF Grant AST 76-21707 and NASA Grant NSG 7212 and at the Aerospace Corporation by the Company Financed Research Program. J. P.° Uefel acknow ledges partial support from the HcCorrolck Fellowship and Compton funds at the University of Chicago. References Blnns, N. R., Fernandez, J. I., Israel, M. H., Klarmann.J., Maehl, R. C, and Newaldt, R, A. 1973, Proc. 13th Int'l. Cosmic Bay Conf.. Denver, 1_, 260. Blake, J. B.t and Schramn, D. N. 1974, Astro. Sp. Scl., 30, 275. 1976, Лр.-J., 209. 846. Blake, J. B., Halnebach, К.ТГ, Scjiranni, D. N.. and Anglin, J. D. 1977, preprint. 273 Blandford, Jr., G. E., Friedlander, И. U., Klarmann. J., Pomeroy, S. S., Walker, R. M., Uefel, J. P., Fowler, P. H.. Kldd, J. M., Kobetlch, E. J.. Moses, R. T., and Thome, R. T. 1973a, Phys. Rev., 06, 1707. Blandford, Jr., G. E., Friedlander, N. W.. Klamann, J., Balker, R. M., and Hefel, J. P. 1973b, Phys. Rev.. 08, 1722. Cameron, A. G. W., 1973, Space Scl. Rev.. 15, 121. Casse\ M., Goret, P., and Cesarsky, C. J. "T975, Proc. 14th lnt'1. Cosmic Ray Conf.. Munich, 2, 646. Couch, R. G., SchmleHekamp, A. B., and Arnett, W. 0. 1974, Ap. J., 190, 95. Fowler, P. H. 1973, Proc. 13th Int'l. Cosmic Ray Conf., Denver, 5_, 3527. Fowler, P. H., Alexander, C, Clapham, V. M., Henshaw, D. C, O'Ceallalgh, C, O'Sulllvan, D., and Thompson, A. 1976a, Proc. 9th Int'l. Conf. on Solid State Nuclear Track Detectors, Munich. Fowler, P. H., Henshaw, D. L., O'Ceallalgh, C, O'Sulllvan, D., and The** an, A. 1976b, Proc. 9th Int'l. Conf. on Solid State Nuclear Track Detectors, Munich. Howard, W. M., Arnett, W. D., Clayton, D. D., and Woosley, S. E. 1972, Ap. J., 175, 201. Juliusson, E., and Meyer, P. 1975,'Ap. J., 201, 76. Lamb, S. A., Howard, W. M., Truran, JT~W., anJTlben, I. 1976, Bull. Am. Astron. Soc, 8, 558. Price, P. B., and Sliirk, Е. К. 1975, Proc. 14th Int'l. Cosmic Ray Conf.. Munich, 1, 268. Schramm, D. N. 1972, Ap. J., 177, 325. Shapiro, M. M., and Silberberg, R. 1974, Phil. Trans. R. Soc. Lond. A.. 277, 319. Shirk, E. K., Price, P. B., Kobetlch, E. J., Osborne, W. 2., Pinsky, L. S., Eandy, R. D., and Rushing, R. B. 1973, Phys. Rev., D7. 3220. Silberberg, R., and Tsao, C. H. 1973, Ap. J. Suppl.. 25, 335. Ulrich, R. K. 1973, in "Explosive Nucleosynthesis", e'ds. D. N. Schramm and M. D. Arnett (Austin: University of Texas Press), p. 139. Wefel, J. P., Schramm, D. N.. and Blake, J. B. 1977, Astro, and Sp, Sci., in press. 274 ON THE DEPENDENCE OP r-PROCESS YIELDS UPON EXPERIMENTALLY UNKNOWN NUCLEAR PARAMETERS; ULTRA-HEAVY COSMIC RAY ABUNDANCES J. B. Blake Spade Sciences Laboratories The Aerospace Corporation D. N. Schramm Enrico Fermi Institute The University of Chicago Calculations of the relative abundances of heavy nuclei (A > 180) in the r-process are presented for three different mass laws and two /3-rate formalisms. The implications of these results for the study of ultra-heavy cosmic rays are discussed briefly. INTRODUCTION Calculated relative abundances of the heavy elements produced in the r-process are important parameters for the interpretation of experimental ultra-heavy cosmic-ray data. The nuclei along'the path of the r-process flow to higher mass are far to the neutron-rich side of ^-stability.' At the present time the nuclear parameters required for an r-process calculation have not been experimentally determined but must be estimated by theoretical analysis and extrapolation from known nuclei. The two major nuclear physics inputs in a r-process calculation are a nuclear mass law and a /3-rate formalism which describe the very neutron-rich nuclei i.iat take part in the r-process. Results are presented in this paper of r-process calculations in which different combinations of recent mass law and )S-rate formalisms are used. The calculations show the nature of the uncertainties in the theoretical r-process predictions due to uncertainties in the nuclear physics inputs. The numerical results are in a form suitable for use in a cosmic-ray propagation calculation and for predicting elemental and charge- group ratios. In turn, high resolution ultra-heavy cosmic-ray measurements (UK-6 and HEAO-C) will provide guidance in selecting mass-law and 0-rate formalisms which | adequately describe Nature, assuming propagation and possible preferential acceleration § effects can be sorted out (Blake and Schramm, 1974; Blake et al., 1977). | 275 CALCULATIONS Two modern dMcrlptiona of the r-process 0-rates are currently available. The Fermi description hai been prepered for convenient astrophysicel application by Senbetu (1173). 1be Gross theory hat been similarly described by Kodama and Takahaski (1975). Varioua там law*, have been wed in r-procen calculation*. The moat commonly used ones have been the liquid-drop model of Myers and Swiatecki (1966) and the more recent droplet model (Myers, 1976). in these models the next magic proton number after Z * II is Z * 114, and the next magic neutron number after N = 126 is N = 184. The mass- law of Soeger and Howard (1975) differs in a very important way from the droplet model in that the next magic neutron number after N = 126 is N = 164 and there is a magic ridge running to N * 184. The droplet model (Myers, 1976) uses the shell correction formalism of the liquid-drop model (Myers and Swiatecki, 1965). A new shell correction description for the droplet model has been described by v. Groote et al. (1976). The calculations were carried out using the three nass laws and two /8-rates in the six possible combinations. As in previous, similar calculations (Seeger etal., 1965; Schramm and Fowler, 197Г; Blake and Schramm, 1974; Mathews and Viola, 1976), the r- procen temperature and density were selected to yield the abundance peak due to N * 126 at A = 195 as is observed in solar-system material. Figures 1, 2 and 3 show the resulting production vs. mass curves for the three mass laws with the curves for the same mass law but different /J-rates plotted in the same figure. The points are normalized to Pt = 0.473 (Cameron, 1973). The high-mass cutoff selected was based upon the fissjon properties of nuclei along the valley of /{-stability (Schramm and Fowler, 1971). T.i-; points not connected by lines at the high-mass end of the curves are in the mass regio.i in which spontaneous fission occurs. For example, 25*Cf decays almost entirely by spontaneous fission and thus does not feed lower masses. There is no yield above A = 263. The low mass limit was chosen at A = 180, which is below the platinum peak. A major difficulty in calculating the production abundances in the rare-earth region is due to the very .uncertain contribution of fission. If the r-process terminates before reaching the next magic number after N = 126, then the fission contribution will be small. If the next magic number is reached, then the fission contribution can be large, cf. discussion in Blake and Schramm (1974). Figure 1 Figure 2 Figure 3 276 The calculated abundances plotted in the three figures are tabulated in the appendix. DISCUSSION The production curves shown in Figure 1 for the droplet model exhibit a "classic" shape; away from the magic-number peak the production is quite uniform. Furthermore, delayed neutron emission during the decay back to the valley of ^-stability from the t- process path during freezeout will cause an appreciable additional smoothing (Blake and Schramm, 1974). It is obvious that use of the Gross theory for the /J-rates substantially enhances the interpeak production compared with the Fermi theory results. As a result, the ratio of the sum of the actinides to the platinum peak The authors would like to thank Lynn Priesen for carrying out many r-proeess computer calculations, Keen Hainebaoh for many useful discussions, M. Yamada for his Gross theory /J-deeay code and H. v. Groot for his shell correction computer code. This work was supported by The Aerospace Corporation Company-Financed Research Program.' REFERENCES Blake, J. B., and Schramm, D. N., Astrophys. Space Sci., 30, L7S, 1974. Blake, J. B., Hainebach, K. L., Schramm, D. N., and Anfin, J. D., Ap. J., in press, 1977. Cameron, A. G. W., Space Sci. Rev. IS, 121,1973. Kodama, T., and Takahashi, K., Nucl. Phys., A239, 489, 1975. Mathews, G. J. and Viola, V. E., Nature, 261, 382,1976. Myers, W. D., and Swiatecki, W. J., Nucl. Phys., A81, 1, 1966. Myers, W. D., At. Nucl. Delta Tables, 17, 411,1976. Schramm, D. N., and Fowler, W. A., Nature, 231, 103,1971. Seeger, P. A., Fowler, W. A., and Clayton, D. D., Ap. J. Suppl., 87, 11, 121. Seeger, P. A. and Howard, W. M., Nucl. Phys., A238, A91, 1975. Senbetu, L., Phys. Rev., C7, 1254, 1973. v. Groote, H., Hilf, E. R. and Takahaski, K., At. Nucl. Data Tables, 1£, 418,1976. 278 APPENDIX v. Groote et al. Seeger-Howard Myers A Gross Fermi Gross Fermi Gross Fermi .022 .008 .028 .008 .001 .001 180 .017 .010 .028 .010 .004 .004 181 .019 .006 .017 .008 .001 .002 182 .023 .013 .027 .009 .012 .009 183 .016 .007 .036 .016 .004 .005 184 .024 .013 .022 .013 .022 .012 185 .023 .010 .037 .015 .018 .010 186 .028 .024 .031 .021 .033 .022 187 .041 .017 .053 .023 .058 .021 188 .072 .063 .075 .059 .082 .064 189 .100 .051 .147 .075 .129 .060 190 .159 .174 .162 .178 .168 .163 191 .206 .110 .291 .151 .266 .131 192 .319 .332 .319 .343 .330 .313 193 .377 .212 .508 .268 .486 .275 194 .473 .473 .473 .473 .473 .473 195 .406 .271 .572 .328 .477 .402 196 .214 .220 .246 .226 .198 .289 197 .160 .090 .238 .113 .184 .134 198 .141 .105 .190 .123 .118 .118 199 . .085 .038 .112 .045 .079 .051 200 ! .089 .049 .075 .039 .071 .042 201 .083 .031 .083 .028 .059 .033 202 ' .067 .032 .022 .010 .037 .020 203 .068 .022 .028 .009 .032 .018 204 .075 .030 .014 .005 .021 .011 205 .052 .018 .030 .008 .015 .011 206 .072 .023 .036 .010 .013 .006 207 .056 .020 .089 .022 .016 .008 208 .054 .019 .061 .013 .018 .005 •- 209 .059 .017 .048 .014 .032 .007 210 .050 .020 .133 .028 .034 .008 211 .046 .013 .110 .036 .039 .010 212 .053 .019 .066 .017 .047 .009 213 .036 .013 .040 .012 .049 .014 214 .042 .014 .083 .020 .043 .010 215 .036 .012 .045 .013 .052 .014 216 .034 .012 .038 .00» .049 .011 217 .029 .009 .030 .008 .046 .016 218 .035 .012 .033 .008 .050 .010 219 .030 .010 .034 .009 .050 .017 220 279 APPENDIX (Cont'd) v. Groote •taL Seeger-Howerd Myers A Otoe» Fermi Gross Fermi Gross Fermi .036 .011 .049 .009 \ .044 .011 221 .036 .010 .054 .015 .052 .015 222 .042 .013 .058 .0i2 .049 .011 223 .036 .011 .050 .013 .046 .016 224 .045 .013 .071 .014 .051 .010 225 .041 .014 .060 .017 .051 .017 226 .044 .015 .074 .014 .045 .011 227 .040 .015 .068 .018 .053 .015 228 .048 .020 .065 .025 .048 .012 229 .038 .018 .052 .014 .048 .016 230 .042 .026 .068 .016 .050 .014 231 .041 .019 .043 .012 .053 .017 232 .034 .033 .058 .018 .045 .018 233 .042 .018 .052 .015 .056 .016 234 .035 .032 .046 .020 .048 .019 235 .049 .019 .071 .017 .048 .017 236 .039 .023 .019 .014 .049 .015 237 .065 .020 .017 .006 .044 .017 238 .044 .015 .018 .013 .044 .012 239 .072 .019 .014 .008 .042 .016 240 .050 .012 .031 .028 .041 .011 241 .068 .017 .034 .017 .052 .018 242 .055 .012 .047 .032 .035 .011 243 .054 .013 .064 .026 .053 .016 244 .070 .013 .076 .031 .034 .012 245 .048 .013 .151 .043 .037 .012 246 .063 .011 .200 .055 .037 .010 247 .052 .013 .307 .067 .030 .011 248 .055 .011 .374 .103 .025 .008 249 .006 .001. .052 .012 .003 .001 250 .060 .012 .413 .126 .026 .008 251 .046 .013 .283 .082 .025 .006 252 .057 Ml .186 .057 .028 .006 253 0. 0. 0. 0. 0. 0. 254 .056 .012 .100 .032 .028 .006 255 .003 .00 i .008 .002 .002 .001 256 .055 .012 .111 .031 .027 .008 257 0. 0. 0. 0. 0. 0. 258 .038 .010 .062 .016 .023 .007 259 0. 0. 0. 0. 0. 0. 260 .030 .008 .030 .008 .019 .006 261 0. 0. 0. 0. 0. 0. 262 .007 .002 .006 .001 .006 .002 263 280 The r-PROCESS PATH IN THE SUPERHEAVY REGION K. L. Hainebach J. B. Blake The Aerospace Corporation Ivan A. Getting Laboratories P. O. Box 92957 Los Angeles, California 90009 D. N. Schramm Enrico Fermi Institute University of Chicago Theoretical studies indicate that neutron-induced fission may halt an r-process which makes the bulk of solar system material before the superheavy region is reached. However it appears that an r-process proceeding along the neutron drip line could reach the predicted superheavy element region. Such an r- process could not have synthesized the bulk of solar system r-process material. INTRODUCTION Blake et al. (1977) have used simple arguments to show that there may be flux of superheavy elements in the galactic cosmic rays as high as - 0.1 the flux of Th and U, or as little as ~ 0.0005 of that flux. Their argument is based upon the work of Anders et al. (1975) who find evidence of superheavy fission xenon in carbonaceous chondrites, in particular the Allende meteorite. Since there is substantial theoretical speculation (although little hard experimental evidence) that superheavy elements cannot be made in the r-process that makes the bulk of the solar system neutron-rich nuclei, it is of substantial interest to consider other scenarios. In this paper, another possibility is examined. CALCULATIONS Fission barrier heights and neutron separation energies for nuclides in the superheavy region, Z < 130, N < 250 have been calculated, and classical (static) calculations showing the r-process path in this region have been carried out. Two different mass laws have been employed: the liquid-drop model of Myers and Swiatecki (1966), and also their improved droplet model in its latest version by Myers 281 (1976). The semi-empirical shell effects described in th« liquid drop model have been used throughout. The last knawn proton shell closure (magic number) occurs nZ< 12. There u uncertainty among theoretical predictions as to whether the next proton mafic number is 114 or 126; therefore calculations have been done with both. The (unknown) neutron magic numbers were taken to be 184, 224 and 256. The results of our calculations are presented in the Figures 1, 2, 3 and 4 which show contours of fission barriers, labeled by the height in MeV. In general, positive barriers indicate nuclei stable against fission, either spontaneous or neutron-induced. The shaded portion represents the solar-system r-process path through the region, and the dashed line is the neutron drip line, i.e. the zero neutron-separation energy contour. DISCUSSION The two mass laws show the same qualitative effects since, in the superheavy region, fission barriers are primarily determined by shell effects, and these are the same for the two mass laws used. The r-process. path passes through the roughly circular valley between neutron magic numbers 184 and 224. If 114 is a magic proton number, then this valley still has positive fission barrier heights. If not, the r-process path crosses the zero fission barrier contour and fission terminates it. For the droplet model, the path nicks the zero fission barrier contour below N = 184 and probably terminates there. However, as Lattimer et al. (1977) point out, it is seen that if the r- process proceeds along the neutron drip line, fission can be avoided in every case in the region contained in the figures. Thus, even if neutron-induced fission terminates the ordinary r-process path below the superheavy region, as may be the case for the droplet model mass law, or in either of the two mass laws if Z = 114 is not a magic number, superheavy elements might still be synthesized in a low temperature decompression of neutron star matter (Lattimer etal., 1977). ACKNOWLEDGEMENTS The authors are indebted to Lynn Friesen for carrying out many numerical computations. This work was supported by The Aerospace Corporation Company- Financed Research Program. т Figures 1-4 show fission barrier contours, labeled by helfht in MeV (solid lines), the r- process path (shaded area), and the neutron drip line (dashed line), in the superheavy region: 80 < Z < 130, 170 < N < 250. Each figure is labeled by the mass lew and proton magic numbers used. Fig. 1 LIQUID DROP MODEL, Z-114,125 MAGIC Fig. 2 LIQUID DROP MODEL, Z=12B MAGIC 283 190 ZOO 210 220 230 240 250 Fig. 3 DROPLET MODEL, Z = 114.126 MAGIC 170 180 190 200 210 220 230 Fig. 4 DROPLET MODEL,Z = 126 MAGIC 284 REFERENCES Anders, Е., Higuehi, H., Gros, J., Takahaski, H., and Morgan, J. W., 1975, Science, 190, 1262. Blake, J. B., Hainebach, K. L., Schramm, D. N., and Anglin, J. D., Ap. J., in press, 1977. Letti.ner, J. M., Mackie, F., Ravenhall, D. G., and Schramm, D. N.. Ap. J., 213, 225, 1977. Myers, W. D., and Swiateckj, W. J., Nuclear Physics, ol, 1, 1966. Myers, W. P., Atomic Data and Nuclear Data Tables, 17, 411,1976. 285 ТНВ0ЖЖТ1СА1. CALCULATIONS 0» THE 55»e ABD 57Co AB01TDAHCES IN GALACTIC COSMIC RAYS £.I.£uzn«teoY«tA.Y.Pi8«nko and A.Z.Lavrukhina V.I.Vernadsky Institute of Geochemistry and Analytical Chemistry.Academy of Sciences of the USSR,Moscow,USSR According to data on abundances of the K-capture radioisotopes in cosmic raya qu estion about the duration of the process of cosmic ray acceleration may be solved. We have made theoretical estimations of the *Fe and 5'Co abundances in galactic cosmic rays at the expence of their formation by nuclear reactions in shock waves being ge nerated by the flash of Supernova and by the fragmentation reactions of the coumic ray heavy nuclei during their passage thro ugh the interstellar matter and at explosi- 57 ve silicon burning.lt is shown that ^'Co in galactic cosmic rays (~-1jt of the 5 Fe con tent) has been mainly formed at the explo sive silicon burning.Therefore it may be uee for determination of time acceleration of cosmic rays.An experiment for ''Co detec tion galactic cosmic rays including the ir radiations of the Al-or Si-targets in cos- •57 mic space and measurement of tha •*'Co-radio activity has been suggested. Determination of the time having passed between the mo ment of the cosmic rays nucleosynthesis and their accele ration till relativistic energies (at) is importance for the solution of the problem of the origin of galactic cos mic rays and the mechanism of their acceleration.In paper fl] the observations of-the gamma-and x-ray radiation du- 286 ring the flash of Supernova at the moment of tha Increa sing and the maximum of ita luminicity haa described.The duration of tha intensive X-ray radiation (-7 daya) coin cides with tha time of cosmic ray acceleration (—11 days) following from the theory of shook ware spreading in a shell of Supernova [2 J .JL same duration (—10 day a) require nuclear spallation raaotiona by tha high-energy protons proceeding in tha shall of Supernova and having lead to the formation of passing isotopes [3j. For the laet years the methods of the д t value deter mination appeared which are based on experimental measure ment е of the Ye,Hi and Co contents in galactic cosmic raya [4f7J and it was shown that the values of acceleration time are equal to : 14 daye-сд t -< 1 year [4] iAt«1 year [5] ; A t» 6 days [7] . More reliable data on the A t value may be obtained when direct measurements of contents of the I-capture nuc lides in galactic cosmic rays will be carried out.from all K-capture nuclides which may be present in the galactic cosmic rays most suitable for this purpose are thi V, •^Mn, Fe and "co radioisotopes (according to the values of half-life,See Table)./or the purpose of determining the possibility of carrying out such an experimentfwa have ma de theoretical estimations of contents of these radioiso- topes and the J Fe content in galactic cosmic rays in the near-Earth space.The isotopic composition of galactic cos mic rays in the source was taken out as analogous to the composition of matter being formed in Supernova by equi librium process of the explosive silicon burning [ej.fe have examined changes of the cosmic ray isotopic compositi on which were caused by spallation reactions proceeding da ring the flash of Supernova,as will as during the passage of galactic cosmic rays through the interstellai matter (4 g/сш of hydrogen) and at the expense of the d.ecay of radioisotopes.Calculations were carried out for the K-cap ture radioisotopes with energy ^ 400 MeV/nucleon,therefore the probability of decay of these nuclei by electron: cap- 287 ture may be disregarded [9,10j .The experimental valuoe of the nuclear reaction сговв-eoctiona in the region of the iron group elements were used [l1,12l .The result of the calculations are given in Table. Table.Change of the 49V,54Mn,55Fe,57Co and 56Fe contents in galactic cosmic rays at the expense of spallation reactions of nuclei which were formed at the explosive silicon burning. 49T 54lln 54Mn * 55,e "co 5**e T1/2 331d 313d 31?d 2.7y 271d stable 25 27 26 28 H 1.2 *10 - - 2flw10 5jO*10 4.5-Ю 1 24 25 24 25 23 23 H 5.8*10 2 4.5*10 9.0-10 6.9*10 3.1-Ю 8.8*10 24 25 24 26 25 27 Ж 2.8*W 1.0*10 1.6-10 3.7»10 8.9*10 ap*io 3 4.6*1026 3.8»1025 4.3-Ю 2б 1.1*1027 6.5*1024 1.8*1025 Я4 H 4.6*1026 3.8-1025 4.3-Ю26 1.5*1027 9.6*1025 8.0*1027 5.8x10"2 4.8*10"5 5.4*10*2 1.9*10~1 1.2И0"2 1.0 J Pa where M. - is the number of nuclei being formed at the explosive silicon burning,in 1cm3 [8ji Bo - is the number of nuclei being formed in spal lation reactions at the Supernova flash, J (ES25 MeV) - 5«1021 proton/cm2; H, - is the number of 'nuclei having been remained after the passage of galactic cosmic rays o through the interstellar matter (4g/cm of hydrogen); H. - is the number of nuclei which had been formed 4 in spallation reactions during the passage of cosmic rays through the interstellar matter; N - is the isotope content being expected in ga lactic cosmic rays in the near-Earth space; 288 •/•** - 1» the isotope content relative to the ' Pe * CCJt«lt{ - under tae condition that the 56Ii and 56Co isotope* had been decayed to the noaent of the acceleration of galactic ooanic rays. It follows from these data that for the determination of the Д t value la suitable only the ''Co radioisotope. The nuaber of its nuclei remained in galactic cosaic rays from nucleosynthesis (under the condition thai Д tj&lyear) is approzlaately 10 times greater thanvwhich was foraed by spallation reactions during the passage of cosaic rays through the interstellar aatter.The other isotopes are in this respect of little interest as aain part of their con tents in oosaio rays are caused by nuclear spallation reactions in the.interstellar space.In this connection da ta on the content in galactic cosmic rays will have a greater interest.This isotope is not formed at the pro- oess of explosive silicon burning.Its presence in galactic oosaio rays is caused by spallation reactions in the inter stellar space and its content ohanges by one order of va lue in dependence on that whether these reactions have proceed on the ' !i nuclei or on the ' Bi and ^ Co nuclei. Therefore the -ratio may be used for the deter mination of the lower limit of the д t value.The **V and 5Sn isotopes nay be used also for the determination of the natter thickness through which galactic cosmic ray pass froa the site of their formation till the Earth, for the detection and measurement of ''Co content in galactic cosmio raya we suggest to carry out an irradi ation of the 11-or Si-targeto in the cosmic space and then,after the radiochemical procedures of cobalt separa tion from the target matter,to measure the "Co radioac tivity by means of low-level counter.For determination of the tia» which is necessary for irradiation of the target and of the target else we have used data on the' denuity of tracks froa VH - nuclei (f y%) of galactic cosmic rays 289 in meteorites.Vhese tracks are mainly caused by nuclei of the iron group elements Г13] .On the Я ун *nd *S7 ^*S6 ~ ratio data we calculated that after a °° f * half-year irradiation of the Al-or Si-spheric target of the 50cm diameter and the 5-10cm thickness the ''Co radio activity will be equal to «-* 50 decay/day .Such a level of the 5^7 Co radioactivity may be measured on a low-level counter. References 1. Sprott G.P.,H.V.Bradt,G.W.Clark,W.H.G.bewin,H.W.Schnop- per .b.Pigatto.b.BosinotAstrophys.J. ,19,1 .H3.part 1 ,739, 1974. 2. Gurevich I,£, ,A.A.Rumyansev t Astnphys.and Space Sei., .22,79,1973. 3. bavrulchina A.K. ,l.I.Kuznetsova > Izvestia Acad.Bauk USS2,ser.fie.,XXXVJJ.,H6 ,1159,1973. 4. Soutoul A.,H.Casae,E.Jul insson j Froo.of 14th Int. Cosmic Bay Conf.,£,455,1975. 5. Bartholama K.P. ,G .Siegznon.W.Enge t Froc.of 14th Int. Cosmic Bay Conf. ,J.,384,1975. 6. Benegae J.C.,H.H.Israel,J.KLarmann,B.C.Machl « Proc. of 14th Int.Cosmic Hay Conf, ,.1,251,1975. 7. Beeves H. t "Origin Cosm.Bays,Proc.Hato,Adv.Study Inst, Durham 1974",Dordrecht-Boston,135,1975. 8. Truran I.W.,W.D.Arnett,A.G.W.Cameron t Canad.J.Phys., 4^,2315,1967. 9. Baisbeck G.M.,G.Comstock,C.Perron,?.Yion t Proc.of 14th Int.Cosmic Bay Conf.,2,560,1975. lO.Tamahane A.S.,V.3.Yenkatavaradan i Proo.of 14th Int, Cosmic Bay Conf. ,2,,564,1975. 11.Lavrukhina A.S.,b.D.Bevina,V.7.Halishev,L.H.Satarova, Su £hunchu,I.S.Kalicheva,L.D.lirsova i Hadiokhimia,5_, 6,1963. 12.LaVrukhina A.I. t "Huclear reactions in cosmic bodies", "Kauka".Moscow,1972. 13.Fleischer Е.1..,P.B.Price,H.M.Walker .M.Maurette t J.Geop^a.Bee.,22,N1 ,331,1967. 290 Monopole, Antinucleus, or Superheavy Nucleus? '«orne. L-S. P.nskv. Physics Oept.. Univ. of Houston. Houstc 7700J. F.B. "rico. E.K. Shirk. Physics Hept.. Univ. of Calif.. Berkeley. CA 91.720 .4. Hlg!trr-, L^rence Berkeley Laboratory. Univ. of Calif.. Berkeley. CA 91,720 Thcoreticui Q Experimental [7] Bo,h Q Calibrations of the Lexan detectors showed that the particle interpreted in 1975 as a nonopole had [Z/el-l 1 ^ through -J.J, g/cmJ -jf natter. Measurements both visually and with an automated microscopic inage dissector showed that in three independent nuclear emulsions the particle produced a far lo-er density of high-energy knock-on electrons than did normal nude! with similar l/B and е г 0.6. These new measurements, together with the most conservative non-controversial interpretation of the Lexan data and Monte Carlo calcula tions of electron energy deposition and fluctuations in the emulsion, estab lish to an extremely high confidence level that the particle could not have been a known nucleus, even if it fragmented several times in the detector V»/2iS^i" fits.of *he data t0 three c'asses of hypothetical particles with |Z/6|=1 J»: a slow (0.35 5 3 < 0.5) particle of very high mass (>103 amu) • a„ „ Д?;7 S 6 S °-9) a"tinucleus with JZ|>80; and an ultrarelativistic IB г. 0.99) superheavy nucleus with Z «•111». Images in the Cerenkov film detectors, now being analyzed, may enable us to decide between the fast and slow hypothetical particles. Coordinates: 0G 1.5 Nuclear Composition of Cosmic Rays Mailing address: Professor W.Z. Osborne Department of physics University of Houston Houston, Texas 77001! 291 THE HADROH EHERGY SPECTRUM AT A 10 0/om2 DEPTH IB THE STRATOSPHERE V.V. Abulora, S.B. Ignatyer, M.A. Ivanova, K.V. Mandritskaya, I.V. Rakobolskayav O.P. Sashina, H.V. Sokolskaya, M.I. Tulinora, A.Ya. Varkovitskaya, B.A. Zamchalova, G.T. Zatsepin, V.Z. Zataaplo Institute of Kuolaar Phyaice, Moacow State University; Moscow 117234, USSR. Chaoretieall Experimental X Both The angular and energy distributions of the electron-photon cascades detected with many-layer X-ray emulsion chambers at a 10 G/cm depth are presented. The total exposure time is 100 m .hour. ГЬе results obtained and the possibilities of studying the chemical composition of primary cosmic rays by the method of X-ray emulsion chambers are discussed. Coordinates: OG 1.5 (ЛКсс-ве<М, ОзмЬо^/У^ <р/ Cos»?'C &#*,) Hailing address: Dr. I.V. Rakobolskaya , Institute of Nuclear Physics, Moscow State University; Moscow 117234, USSR. 292 MEASUREMENT OF HELIUM ISOTOPIC COMPOSITION Ш PRIMARY COSMIC RAYS. M.G.Jodko, V.K.Karakadko, V.A.Romanov. A.F.Ioffe Physical-Technical Institute, 19402"!, Leningrad, USSR. He-VHe^r He* ratio in the energy range Е £.1.13 Gev/nucleon was measured by means of scintillation counters telescope, directed or ganic glass Cerenkov counter and two gas Ceren- kov counters. The value of Не3/Не3+ He4= 0.15 *• 0.04 was obtained. This investigation was car ried out at middle latitudes using balloons. Introduction: A lot of measurements of the He3 countent in primary cosmic rays refer to the region of comparatively low energies from 100 to 400 llev/nucleon [1,2^ . The spectrum of particles in this energy region is subjected to the strong mo dulation in the solar sistea. We tried to measure the HevHe •+ He4 ratio in the region of higher energies (~1 Gev/nucleon), where the modulation is relatively weak. The experiment is based on He* and He^ sepa ration in the magnetic field of the Earth. The geomagnetic thresholds for He and He^ at the point of measurement are 0*7 Gev/nucleon and 1 Gev/nucleon respectively. The intensity of ot-parti.-.les in different energy intervals was measured by the scintillation and Cerenkov counter telescop. Having measu red the intensity of ot-particles in the interval 0.7-1.0 Gev/ nucleon (intensity of He ) and in the region of higher energi es (>1.5 Gev/nucleon, where the sum isotopes He4 and He3 is recorded} and knowing the spectrum of particles [3,4] , it is easy to calculate the ratio of the isotopes in primary cosmic rays. The separation of the isotopes by the geomagnetic field takes place before the penetration of the particles into the 293 itmofiphcre. Tncrefore corrections for the change of ieotopic -„mioGition during the passage of the particles in the residu al ntmosphere above the device should be email in this Method. l'reliminnry results of the 1972 flights пате been publi shed earlier [5] The ratio He3/He3t- He*« 0.2010.06 for Ь';. Л ,ЪЪ Gev/nucleon was obtained with the assumption that spectrum of He does not differ con iderably from the spec trum of all oC-particles. In the data procession we intro duced the correction for under - threshold dL-partiolss,recor ded by gas Cerencov counter due to 8*-electrons formed while еж-particles pass through the gas of the counter. Pressure used in the counters being high this correction is large and its wrong account can leard to enormous mistakes. In [53 the calculated values of the correction factors were used. In ex periments we tried to minimize this correction. 2. Experiment. Measurement of o^-P&rticle flux were car- ried out by the telescope with geometrical factor Г=13«1 cm ster. The telescop was lifted on the balloon up to 8-10 g/cm of the residual atmospere. The device consisted of two scintillation counters deter mining solid ancle of telescop, the directed Cerencov counter with the plexiglass radiator and two gas Cerencov counters. Gas pressure in the first gas Cerencov counter was kept con stant and of such a value as to provide the measurement of c?L-particles in the enorgy range 0.7 -1.0 Gev/nucleon during all the flight. Pressure of the second counter was changed in the course of the flight according to the programme. The ran ge of energy thresholds of the second counter was from 1 to 5 Gev/nucleon. Inpulses from the upper scintillation counter and from directed Cerencov counter avere analyzed by a* multi channel anali"i»r. The additional thresholds were established in electronic recording schemes of gas counters and marks for recordind particles resulting in pulse amplitude below and above those thresholds. It was necessary to estimate the influence of S - electrons on the registration of oC-particles in the gas counters. All this increased to some extent uncertainty in the ?С4 energy thresholds, >uit allowed ua j eliminate the counts due to 5"-electrons. Supplementary mi ltlchh.mel analizer was estab- lished in order to estimate whether tne thresholds were Steele. This apparatus analyzed pulses from the gas counters and the Becond scintillation counter in turn. 3. Results. The data of daily flight which was carried out in September 1976 were used. Gas pressure in the gas coun ters kept constant. Measurements were carried out practically at constant latitude. Diagrams of o£-particles recorded during the flight at different thresholds of electronic schemes of v -J gas Cerencov counters were drawing. In order to calculate He^ we used intensities obtained from the thresholds at which the main part o-electron events were not recorded. As before [5] we used the calculated value of penetretion of penumbra £=0.49. Energy thresholds of the gas counters were estimated from the data of the gas pressure and temperature and the ratio of o^-particles intensity at different thresholds. While estima ting the thresholds we used also calculated values of light output in the gas Gerencov counter and calibration curves ob tained with the muons during tuning procedure. The data of 9 hours of the flight 1976 were examined. For Б^.1.86 Gev/nucl. the ratio He3/He3+ He4= 0.14± 0.05, for the 1.13-1.86 Gev/nucl. range He3/He3+ He4= 0.14- 0.09 were obtai ned. With the account of the earlier obtained results we have the following: for B^. 1.13 Gev/nucl. He3/He3+ He4= 0.15±0.04. The measured flow of oL-particles with Е ^1.86 Gev/nucieon coincides within the limits of mistakes with the flow obtained in work [ 3 ~\ . References. 1. Famadurai S. and S.Biswas, Astroph.Sp.Sci., 30, 187, 1974. 2. W.R.Webber and N.J.Schefield, Proc.Hth Intern.Conf.Cosm. Ray, Munchen, ±, 312, 1975. 3. R.S.Preier, G.J.Waddington, J.Geophye.Res., 73,4261,1968. 4. W.R.Webber and J.A.Lezniak, Astroph.Sp.Sci.,30, 361, 1974. 5 М.Г, Иодко, В.А.Романов, Изв AH CCCP.cep физ,38, 1813, 1974. 395 ГКЕ ISOTROPIC CUNTOMTiON OK Ll-lle «ud H F. B. McDonald and J. H. Tratnor NASA/Coddard Space Flight Center, Greenbelt, Maryland J077: r.s.A. and W. R. Webber University of New Hampshire, Durham, New Hampshire 0382-, U.S.A. Theoretical Q Experiments! jTJ Both Q jPreviously it has been shown that the Coddard-University of Mew Hampshire ' •cosmic ray experiment is capable of resolving individual isotopes for elements ! extending from Hydrogen through Nitrogen. This superior resolution is due in part to the use of three parameter analysis (2 measurements of <*F.'dx i. well as the residual kinetic energy of particles stopping in the telescope) ; and the small opening angle of the telescope. A typical energy interval for three parameter analysis is 30-56 MeV/nuc for Helium nuclei and 55-75 MeV/nuc for Carbon. Complete separation is achieved for the isotopes of Li, Be and B.^However, becaus-e of the very snail solid angle of the telescope (2.1 x 10 m2-ster) it has been necessary to accumulate data ever long periods of time. Data has now been accumulated from Pioneer 10 since March 1972 out to 12 A.U. and from Pioneer 11 since April 1973 out to 4 A.U. Data on the low energy portion of Li, Be and В spectra will be presented along with the isotopic composition. Coordinates: Primary cosmic ray isotopes 0G(1.6) Mailing address: Dr. Frank В. McDonald Code 660 NASA/Goddard Space Flight Center Greenbelt, Maryland 20771 U.S.A. :96 OKMIC-KAY AGE AND A MEASUREMENT OK THK ABUNDAKCK OF B USING A MAGNETIC SPECTROMETER * CD. Orth, A. BulTington, P.M. Lubln, T.S. Meat, and G.F. Saoot Space Sciences Laboratory and Lawrence Berkeley Laboratory University cf California, Berkeley Berkeley, California (USA) 94720 ABSTRACT The abundances of isotopes from Li to C have been measured using our superconducting magnetic spectrometer in a balloon-Ъогпе experiment flying at a residual atmosphere of about 6 g/cm^. In particular, we report the measurement of the isotopes of beryllium from 1.5 to 3 GV/c (i.e. 200-600 MeV/nucleon for Be10) with a mass resolution ranging fr^m 0.25 to O.k anru and with high efficiency. The detection technique, which utilizes dS/dx measured in a 2.5 cm plastic scintillator versus rigidity measured in the spectrometer, is described in detail along with an accelerator calibration of the scintillator. Mean ages for cosmic rays can be inferred from these Be observations and calculated abundances using several propagation models. I. INTRODUCTION Radioactive isotopes have long been recognized as clocks which could measure the cosmic ray propagation time. Be is practical for such a measurement since it is the most abundant radioactive isotope and because its mean life cf 2.2 (half life = 1.5) million years approaches the expected range of cosmic ray ages. In this paper we describe the construction and calibration of an instrument capable of separating the isotopes of Beryllium. Previous experiments (Preszler, et al., 1975» Garcia-Munoz/" et al., 1975i and Hagen, et al., 1977) have separated Be10 within an energy range from kO to 290 MeV/nucleon; the present experiment covers the energy range from 200 to 600 MeV/nucleon. At these higher energies, interpretation of the results is less subject to systematic uncertainties from solar modulation and energy depend ence in the Be production cross sections. The apparatus separates isotopes by using a superconducting magnetic spectrometer to measure the rigidity of each event and a large plastic scintillator to measure the specific energy deposition. * This work supported by N.A.S.A. grant NGR-05-003-553, and by the Lawrence Berkeley Laboratory. Л! KAKATU" Л schi-mntic »!' tin- upp:initu:. Ь ;.hown in f-'i^ur- I. , (••• t-.lK t.-l .•':. ::ti,5n<-tic spectrometer la described iti Umool, el ui . {1 .) i ). V ..r modifications hiive been made :>ince then in convert th'- чррщтли:. ;'or !-«• measurement: (l) ft scintillation counter call-м th- ":..i ] •• c uiit'-i" tin- beer. Introduced between scintillation trigger counters o\ чп i :.;ч , (;) the middle ьрагк chamber SC-2 was replaced with n thin-!' 11 uult , (j) the trigger counter Sp was replaced by an anticoincidence counter hole to define the region of high magnetic field integral, arid ('J) the trigger .••.-quirements have been changed as described below. Changes (2) mid (j) were made t minimize the material for multiple Coulomb scattering in the spectrometer. The superconducting magnet, cameras, and optics remain unchanged. The 2.5 cm thick isotope counter scintillator was placed near the Up of an aluminum box (l6 x 51 x Bl cm ) painted insid" with white diffusive BaSO^ paint. The reflectivity averaged over the box is 97jE. The box is viewed by four 5-inch photomultipliers (RCA C701j,j) through air light pipes (15 cm long) constructed from aluminized plexiglas. The total collection and conversion efficiency is about if. The photomultipliers are in a region of inhomogeneous magnetic field that averages about 1 kilogauss. The severe shielding problem WP.S solved with a combination of mu-metal and soft iron shields, and bucking coils around both the light pipes and the phototube necks. The coils typically have 100 to 200 ampere-turns. The weight of each photo tube, lightpipe, and shielding assembly is about 25 kilograms. The anode signal of each tube is individually digitized with 1024 bits and recorded by bit lights on the film. In addition, the last dynode signals are passively added, inverted, and digi^iz^n. Table I gives a summary of specifications for the apparatus. TABLE I: APPARATUS SPECIFICATIONS SI 1.3 x 89 x 69 cm Pilot Y S2 56 x 33 cm (hole) NE-110 Pilot Y Scintillators NE-110 { Pilot Y S3 0.9 x 76 x 46 cm Magnet: Isotope 2.5 x 76 x 46 cm AbovSit- e SI 1.3 x 86 x 6l cm Gondola Mat'l: WithiMagnetn Antspectrometei r76 cm diameter Tota1.6U-l xinstrumen ю5 ampere-turnst , yielding 5 kG-m mean field integral 0.8 g/cm2 3.4 x 10"3 rad. len. 8 g/cm2 ;98 The thretholds for the trigger counter pulae heights art- J<-riii<-'. :ti i- r—... of the «D«t probtfble lonliatlon of a reltitlvlatlt: i>l particle tmd arv art using ground level шиопв. A threshold of 11 times Z»l la usej t\ r Sj tuid .• to ellalnate the large flux of hydrogen and helium. A threshold o: 1 Mors Z"l is, used for Sb to guarantee full penetrutlon of the Isotope counter. A radio-coaeandable upper threshold on S^ 3et ut 7b times 2^1 excludes the oxygen and higher-Z nuclei. Figure 2 shows the region of dE/dx versus rigidity space that the trigger will accept. III. RESOLUTION The expected error on the final relative abundance of Be10 is a function of the total number of events observed and the mass resolution. The expected mass resolution 6(A/z) is plotted as a function of rigidity in figure 3. In the interval from 1.5'GV/c (geomagnetic cutoff) to 3 GV/c, the resolution varies from 0.25 to 0.4 amu for Beryllium. At lower rigidities the resolution'is dominated by the error in the rigidity measurement, but at higher rigidities variation in dE/dx is dominant. Using a flux of 0.3 events per m ster sec in the region of good resolution between 1.5 and 3 GV/c, and our geometry factor of 0.08 master, we estimate a 3 total of about 10 well-resolved Beryllium events will be recorded during a flight with 10 hours live time. IVk ACCELERATOR CALIBRATIONS The isotope counter was tested at the Berkeley Bevatron in beams of HlJ45.8 GV/c), С^г.о- GV/c), and ^(3.5 GV/c). In addition, the response to ground level cosmic ray muons was measured, giving typically 300 photoelectrons for each Z=l particle incident. For each of the beams, the peak and the full width at half maximum of the pulse area distribution were measured. The ratios of the peaks and the widths are consistent with the Simon-Landau theory, the known photoelectron statistics, and scintillation and electronic saturatior. During the Bevatron runs the counter was moved transverse to the baam to map the spatial uniformity. The peak of the pulse area distribution fcr each of the «four tubes was recorded for each beam position, and was found to vary by as much as a factor of two. However, the linear sum of the four tubes with an optimized .gain for each gave a signal which was quite uniform over the counter. The rms deviation about the mean of this signal was about 3$. A correction for this small non-uniformity can be made from the flight data usiiu 299 the abundant reletivistlc carbon nuclei events. The residual ooo- uniroraltie» will be le»a than, one-half percent and thue will ваке a negligible contribution to the mi resolution. V. PRELMDUBY RESULTS At the tine of eubedaaion (1 May, 1977), the Instrument vu ready for launch from Aberdeen, South Dakota. Hopefully, a successful flight will be achieved and, by the tlae of the conference, preliminary restate will be available. REFERENCES Preszler, A.M., Klsh, J.C., Lezniak, J.A., Simpson, G., and Webber, W.R., lUth International Coemic Ray Conference, 12, page 1*096 (1975. Munich). Garcia-Munoz, M., Mason, G.M., and Siapson, J.A., Astrophyalcal Journal Letters, 201, LlU-l and UA5 (1975)- Hagen, F.A., Fisher, A.J., and Orrnes, J.F., Astrophysical Journal, 212. 26> (1977)- Smoot, G.F., Buffington, A., Orth, CD., and Smith, L.H., 13th International Cosmic Ray Conference, 1, page 225 (1973» Denver). GONDOLA SMELL SI SCINTILLATOR SC-1 SPARK CHAMBER If MAGNET ANTI SC-2 sc-3 S3 ISOTOPE COUNTER FIGURE 1. A schematic of the apparatus. 300 ХР/ЯР 301 THE 1S0T0PIC COMPOSITION OF GALACTIC COSMIC RAY LITHIUM, BERYLLIUM ANO BORON M. Garcla-Munoz, G. M. Meson «nd J. A. Simpson Enrico Fermi Institute, University of Chicago Chicago, Illinois 60637 (USA) The 1sotop1c composition of galactic cosmic ray Li, Be, and В has been measured near 100 MeV/nucleon using the University of Chicago IMP-7 and IMP-8 cosmic ray telescopes during 1973 - 197S. The measured abundances allow us to make detailed checks of models of inter stellar propagation and solar modulation and to draw conclusions concerning the spectral forms at the source, and the minimum solar modulation level. For example, comparing these results to local interstellar spectra calculated using a "leaky box" model, 1t 1s found that if we Ignore solar modulation there is no unique leakage mean free path consistent with all the observations. However, by taking account of a sizeable level of residual solar modulation, excellent agreement is obtained between the calculated and measured abundances. Thus, these isotopic abundances confirm the old hypothesis that cosmic ray Li, Be, and В are produced as secondaries in interstellar space. 1. Introduction. The substantial amounts of the "light nuclei" (Li, Be, B) discovered in the cosmic rays was at once interpreted as the result of spallation of heavier nuclei in the interstellar gas (e.g. review by Peters 1952). As measurements of the cosmic ray abundances and the relevant spallation cross-sections became more precise, it became possible for the models to explain not only the total L1+Be+B abundance, but also the abundances of the separate species (e.g. review by Shapiro and Silberberg 1970). More recently, measurements of the isotopic abundances of cosmic ray Li, Be, and В have made 1t possible to refine the earlier studies, and confirm the spallation model of their origin in detail (e.g. Garcia-Munoz et ab_ 1975a, hereafter Paper I; Webber et al^ 1973; Preszler et al^ 1975T~Hagan et al. 1977). We now report a continuation of the sateTTlte measurements in Paper I over a longer time interval to increase the precision of the study. We find excellent agreement between our observations and model calculations for a 5 5 g/cm2 exponential path length distribution with source differential energy spectral forms like (T+400)*2-6 and with solar modulation of strength characterized by a "force-field" deceleration parameter (Urch and Gleeson 1973) Ф = 220. MeV/nucleon. This paper also introduces an analysis technique that uses these isotopic abundances to place constraints on various parameters in the models. For example, we find that if the errors in the propagation model calculations are taken to be ^ ± 10*", then our observations Impose the following limits on the models: the source spectra must be flatter than about (T+200)-2'6, the leakage mean free path must be £ 7 g/cm2, and the deceleration parameter must be Ф £ 100 MeV/nucleon. We note that the limit on the leakage mean free path deduced here shows that this parameter 1s consis tent with energy independence for observations below several GeV per nucleon. 302 1MP-T.|w»-1 Uo.t Histogram» Measurements. The measurements reported here were carried out 1n InterplaneU-y space using the University of Chicago IMP-7 and 1MP-8 cosmic ray telescopes, ard, unlike bal loon measurements, are free from corrections for atmospheric or geomagnetic effects. The data were collected under near solar minimum conditions during 27.900 hours of recording time between 1973 and 1975, Identically the same time Interval we used for our study of the Be™ abundance (Garcia-Munoz e£ al. 1977a). Data analysis techniques were Atomic Mass Units essentially the same afc described in Paper I, except Figure 1 that (1) the Be isotope separa tions in this work make use of a detailed accelerator calibration of the back-up instrument as described b_y, Garcia-Munoz et ajN, (1977a), and (2) the energy interval for boron i so topic analysis has Been narrowed in order to eliminate a region of low-gain in the instrument pulse-height-analyzers where the isotopic resolution was slightly degraded. Figure 1 shows the Li, Be and В mass histograms resulting from the analysis, and compares them with the peak shape for carbon. The number of events of each isotope Was determined from the histograms of Figure 1 by fit ting gaussian peak shapes to the distributions using a maximum likelihood technique. The results presented here are compared in Table 1 with balloon measurements at somewhat higher energies by the University of New Hampshire COMPARISON OF NE/.JUREHENTS OF THE ISOTOPIC COMPOSITION OF W: LIGHT NUCLEI IN THE LOU ENERGY GALACTIC C0SH1C RAYS EXPERIMENTERS University of Goddard Space University of New Hampshire Flight Center Chicago YEAR 1972 1974 1973 1973 - 1975 ENERGY INTERVAL 127 - 249 145 - 245 161 - 416 31 - 151 (KeV/n) Table 1 6 L1 /Lf о.га 0.55 10.05 0.51 +0.04 7 U /L1 0.62 0.45 * 0.04 0.49 + 0.04 7 o Be /8e 0.57 0.60 + 0.07 0.63 + 0.07 0.63 + 0.03 S 9 u Be7Be 0.34 0.32 + 0.06 0.27 + 0.05 0.34 + 0.03 ,0 1 B. /B. 0.09 0.08 + 0.03 0.10+ 0.04 0.03 + 0.01 S 10 B /B 0.25 0.31 +0.03 0.33 + 0.04 0.29 + 0.02 »"/« 0.75 0.59 +0.03 0.67 + 0.05 0.71 + 0.02 303 отмия OF «EASUHO m UUCUUTD KLMI« «мвмсп OF nc нит ISOTOKS i« тж uutcnc asnic r*n FM A- S.S s/o**. i„ • «o r*v/ujcno». ?• o.j /mrc/m5 Iklitbi 10 c" • CU • 103. l»-l • »-« ku uibimi litUH (Mr0J> Intorttl U'MUIM I»l4tl4 H»Hllll (ПгГ/О*!».) •f frf* L.' »- и « . 11- I.I « O.J 1.1 • 1.1 1.1» 0.1 M' П - 111 >M . 11 3.1 • 0.3 1.1 • 0.1 u:i.i M' 34 - 111 ISI . 14 I.l-O.l I.I I 0.4 0.1 1 0.2 h10 31 - 111 12» t 0.11 I.0.0» 0.11 • 0.01 0.011 • O.0M I10 41 • 121 til «13 1.1» 0.1 «.; • I.I l.<*l.t' •" ii - 112 10S0 «42 11.3 » 0.1 li.S »1.1 11.0 i 2.2 c" • c» IS - 111 1U0O . ISO 100 100 100 «s - in 7Н7.Ю» ICO 100 1» * ПМ 44U MljT group (Webber et al^ 1973; Preszler et al^ 1975) and the Goddard Space Flight Center group (Яадеп et al. 1977). THe" present results are In good agreement with the 1974 UNH ano^tKe-GSFC measurements except for the Be10 abundance (see discussion 1n Garda-Munoz et al. 1977a, b). 3. Model Calculations. We have used a steady-state, homogeneous model of cosmic ray propagation, which gives rise to an exponential distribu tion of path lengths through the interstellar gas, which 1s assumed to be of solar system composition. Details of the calculation, which Includes effects due to solar modulation, are given 1n another paper at this conference (Garcia- Munoz et al. 1977b). Table 2 compares the measured and calculated isotoplc abundances for the IMP-7/8 measurements and our assumed Tiodel. It can be seen that the model fits the Individual 1sotopic abundances quite well provided solar modulation effects are Included. Note, however, that there 1s a significant disagreement if we ignore modulation effects and attempt to compare directly the data at Earth orbit with the calculated local Interstellar abundances. 4. Model Constraints Imposed by Light Isotope Abundances. Earlier studies of model constraints imposed b.v the н and He isotope measurements (e.g. Mewaldt et al. 1976 and references therein) have been limited primarily by the experlmentaTTrrors in determining the rare H2 and He3 abundances. In this section we carry out a similar.study in which we examine the extent to which there are significant constraints Imposed by our relatively more accurate measurements of the Li, Be, and В Isotope abundances on the following three parameters In the model: (a) the leakage mean free path. Ле, (b) the level of solar modulation, which we characterize in this paper by the force field para meter Ф, and (c) the assumed differential energy spectra at the source, J = A(T+To)-2«6, characterized by the parameter T0 (MeV/nucleon) where A is a constant and T 1s the'particle kinetic energy (MeV/nuc). The spirit of this discussion 1s to show the trends and limitations revealed by varying the para meters mentioned above and to demonstrate a new approach to the selection of 304 Modulation paraaw/tdra atwM possible by Isotmpfc mm •ants of tM I1«|i^miel«t. we start iritt sourca spactra _jf6tha for. (T-MOO)"2*0. whore T 1» ki netic energy par nucleoli. This for* was chosen because after propagation with a leakage nan free path At ft 6 g/cm* It agrees rather closely with the local Inter stellar carbon spectral form derived In a study of 1973 spectra observed at 1 AU (Garcla-Hunoz at al. 1975b). Then, Figure 2~sbows the "growth"curves'1 for the light Isotopes as the leakage mean free path 1$ Increased for the case (part A of the fig ure) of no solar-aedulatlon, and the case (part ») of our adopted 197S level of solar modulation. The IMP observa tions are the horizontal hatched strips, while the growth curves (shown with 1 only a i 10X error) Increase Ltokogt Mton Fi«t Polh (9/cm ) with Increasing grammige. Note that for the assumption Figure 2 of no. solar modulation (part A).., no single valuvoiwe оoтf thtne leakc^leanc^e mean freтгеeс pathn 1s"tbns1stanis consistent witwmh aall'thi e observations. However, Part В of Figure Z shows that the.growth curves after IneTSJlmg a modulation characterized by a force-field decoloration Ф - 220 MeV/nucleon are a good fit to the data for i^ • 5-6 g/cm*. Figure 3 shows the variation 1n the Isotope ratios as a function pf leak age mean free path. The shaded areas In the calculated ratios represent a + 20% uncertainty which arises from errors In the spallation cross-sections. It 1s seen that the ratios change very little over a wide range of leakage mean free paths, but that 1n the cases of 11 and Be the direct comparison of calculated local Interstellar ratios disagrees rather strongly with the' IMP observations. On the other hand, Including solar modulation with * - 220 HaV/ nucleon brings the calculations Into excellent agreement with the 1 AU obser vations, as shown by the arrows 1n the figure which show size of the downward shift 1n the Interstellar,ratios under solar modulation as the particles propagate In to 1 AU. These changes arise due to the different A/Z ratios of the particles as well as their differing spectral forms 1n the local inter stellar space. We now consider variations In the source spectral forms as well as In the leakage mean free path and solar modulation strength. Figure 4(a) shows the leakage mean free path required to fit the IMP observations for spectral 2,6 forms (T+TQ)" when T0 Is changed as shown-In the figure. Notice that for з 305 modulation level • • 220 MtV/nucleon, that for Bt>B* Tp £ 200 MtV/nucleon there Is no slnglt value of tht leakage mn frtt path consistent with «11 the IMP observations. We conclude from Figure 4(a) that the source spectral forms must be flatter 2 6 than (T+To)" - where T0- 200 MeV/nucleon. A a-t««MMi> rigidity dependent source spectrum would fit this MtoMIAJUl criterion. II S Laokop-*- Ми+я Frt * Pgtti (fl/cm*)-in* Figures 4(b) and (c) Figure 3 examine the variation of the modulation strength allowed by the observations. As the Modulation strength 1s decreased from Ф « 220 MeV/nucleon to • - 0 (no modulation), the values of the leakage mean free paths consistent with the, various Isotope abundances spread out more and more. Below about » » 100 MeV/nucleon, the spread 1s great enough so that no single escape mean free path value fits all the date, and from this we con- elude that the lower limit on the modulation strength Ф Is about 100 MeV/ nucleon. Because these observations were made near the minimum phase of the current solar cycle, this constraint applies to the residual modulation near solar minimum. It can also be seen from Figure 4 that for a value of T0 » 200 MeV/nucleon, with Ф - 220 MeV/nucleon, or if >3" &*"> W" n *"* for the case where T0 • if . = сз =• а Nltfr/Мм. 400 MeV/nucleon with Ш/ а c= • o i Г . Itt to «-ОО MM VI MM* Ф - 100 MeV/nucleon W => сз c=a e= (that 1s, In the range f o ca а c=> a of parameters at the а * L^J 1'i'rT'i'>'a>»'' Hm1t of the allewed •+ЧГ Л,ЛЛт> values) that the maxi mum value of the leak pB I 4*BO I _>-l* i • iqb • age mean free path if allowed 1s about 7 if g/cmZ, since for higher If Ле values the mean, free - f path Intervals required to fit the observations ili't'1t'r''r»'i'''t'»4'i'a'l'ft'>,iralalfr'»'fA' are not consistent with Л.. «Ля* a common value. lei •46 I ' ••**' I ' *H* I __ ••» I " if We must emphasize if that Figure 4 Includes no uncertainties 1n the calculations. Include Ing a +,201 uncertainty ft'»'lf*'''t'Vt'J"l»'i'tla,ll*'*'«'>'»'f»,*l*lB In the calculations «•.ft»1 spreads out th» ranges, of escape meaifctfree Figure *' fiattis consistent with 306 the observations, *nd mikes the T0 and * Kelts less stringent. Finally» ме conclude from Figures 2 and 4 that In the case of a source spectrum of the form (T+400)_z-6 and a level of modulation • • 220 MeV/nucleoa the best fit to the data Is given by an escape mean free path of 5.5 *_ 1 g/cm< 5. Discussion. The detailed IMP-7/8 1sotop1c abundance measurements of LI, Be and В are seen to be 1n excellent agreement with a steady-state model of cosmic ray propagation with a 5.5 g/cm2 leakage mean free path, after Including effects of solar modulation. Using a solar modulation level deter- mlned by the 1973 electron observations at 1 AU, and taking account of obser vations of the low energy proton and helium spectra as well аь the Pioneer-10 gradient measurements (see discussion 1n Garcla-Munoz et al. 1975b) yields a modulation which brings the calculated local Inters telTar~ab~undances as well as the Individual Isotope ratios Into good agreement with the observations at Earth orbit. Varying parameters In the models show that these observations can place limits on the source spectral forms, the modulation level,' and the leakage mean free path. The present limits of this technique arise from uncertainties In the spallation cross-sections, and therefore future improve ments await more experimental determinations of these cross-sections. 6. Acknowledgements. We are grateful to the staff of the Laboratory for Astrophysics and Space Research for the cons'truction and data processing for the IMP-7 and IHP-8 instruments. This work was supported 1n part by NASA contract NAS 5-11067 and grant NGL 14-001-006 and NSF grant ATM 75-20407. References Garcla-Munoz, M., Mason, G. M. and Simpson, J. A. 1975a, Ap. J. (Lett.) 201. L145. 1975b, Ap. J. 202, 265. 1977a, submitted-to Ap. J. 1977b, this conference, paper OG 84. Hagen, F. A., Fisher, A. J. and Ormes, J. F. 1977, Ap. J. 212, 262. Mewaldt, R. A., Stone, E. C. and Vogt, R. E. 1976, Ap. J. Ж, 616. Peters, B. 1952, Progress In Cosmic Ray Physics, (Amsterdai5r"North Holland Pub. Co.), 1, 191. Preszler, A. M., Klsh, J. C, Leznlak, J. A., Simpson, G. and Webber, W. R. 1975, 14th Int'l. Cosmic Ray Conf. 12, 4096. Shapiro, M. M. and Sllberberg, R. 1970, АТГП. Rev. Hue. Scl. 20. 323. Urch, I. H. and Gleeson, L. J. 1973, As trophy s. and 5расе~5сТГ 20. 177. Webber, W. R., Leznlak, J. A., K1sh, J. and Damle. S. V. 1973. Astrophys. and Space Sc1. 24, 17. 307 THE IS0T0P1C COMPOSITION OF GALACTIC COSMIC-RAY BERYLLIUM AND THE COSMIC RAY AGE M. Garcia-Munoz, G. M. Mason and 0. A. Simpson Enrico Fermi Institute, University of Chicago Chicago", Illinois 60637 (USA) We report a continuation of measurements of the ^lactic cosmic ray beryllium 1sotop1c composition near 100 MeV per nucleon, using data collected through 1976 by the University of Chicago Instruments on the IMP-7 and IMP-8 earth satellites. It 1s found that Bel°/Be * (3 + 1.4)*. much less than the ratio expected on the basis of Inter stellar propagation calculations where 1t 1s assumed that the Interstellar gas density is •*. 1 atom/cm3. These results extend and confirm our earlier measurements, and carry increased statistical precision. For a homogeneous, steady-state model of cosmic ray propagation, we find a cosmic ray lifetime for escape of (16 +_ 2J) x 106 years, and an average interstellar density of (0.19 +_ $• J]) atoms/cnw. 1. Introduction. The presence of relatively large amounts of secondary species such as Li, Be and В in the galactic cosmic rays has long been recognized as a measure of the amount of material through which heavier parent cosmic rays have passed — whether in or near their sources, or in interstellar space. Be1", with a half-life of 1.5 x 106 years, is a well- known "clock" which could provide a measure of the time required for the heavier cosmic rays to traverse the material in their path, and thus a number of studies have been made to experimentally determine its abundance (e.g. Garcia-Munoz et al. 1975, 1977b, hereafter called Papers I and II; Hagen et al. 1977;.Pr"eszTer et al. 1975). We reported in Paper I independent measurements on two satellites of~the Be'0 abundance using measurements taken during 1973-74, and extended this work in Paper II to Include 1975 data as well as the results of a calibration of back-up instrumentation using high energy Be beams accelerated at the Lawrence Berkeley Laboratory Bevatron. In the present report, these measurements have been carried through 1976, thereby nearly doubling the number of analyzed events of the first report. 2. Instrumentation. Figure 1 shows a cross-section of the IMP-7 and IMP-8 cosmic ray telescopes, which are of Identical design. In the figure, Dl, D2, and 03 are lithium-drifted S1 detectors of thickness 750, 1450, and 800 microns, respectively. The 04 1s a thick (11.5 g/cm2) Csl (Tl) scintillator viewed by 4 photodiodes. D5 Is a 3.98 g/cmZ sapphire Cerenkov radiator which also has a significant scintillation output, and which is viewed by a photomultlpHer tube. D6 is a plastic scintillation guard counter which defines the telescope geometry and can detect nuclear interaction events 1n the telescope. Asterisks denote detectors whose output 1s pulse-height- analyzed. 308 The events considered in this paper passed through the Dl, D2, and D3 UNIVERSITY OF CHICAGO detectors and stopped in the 04 detector ШР-7/8 (that is, did not trigger D5 or D6). For such events, the Instruments perform two Independent dE/dx measurements in the curved Dl and D2 detectors, and a residual energy measurement in the D4. By requiring that the two dE/dx measure ments be consistent with each other, particles passing thrcigh the detector edges, and certain interacting events are easily eliminated. For the retraining events, the Dl and 02 detector outputs are summed together and plotted versus the D4 detector signal, yielding dE/dx versus residual energy matrices in which the Be' and the Be9+Be10 tracks are clearly separated (see Papers I and II). At tms stage in the analysis, the determination of the distance between the Be7, the Be9 and the Be™ track centers becomes the key step in inter preting the data. This is not a straightforward task, since the Csl scintillator has a non-linear light Figure 1 output which cannot be calculated with a high level of confidence. We therefore took the identical back-up IMP-7/8 instrument to the Berkeley Bevatron, and exposed it to beams of high energy BBe'e , Be9, and Be10, using the Heckman spectrometer facility to separate each of these isotopes before directing the beam into the instrument. In this manner, the location of the isotope tracks in the back-up instrument matrix was unambiguously determined. Under the reasonable assumption that the relative separation of the Be tracks is the same in the back-up and the flight instruments, the measured track separations then formed the basis for analyzing the IMP-7 and IMP-8 flight data, wherein the Be' track centers were located, and relative distances to the Be9 and Be'0 tracks used were those measured at the Bevatron calibration. After analyzing the IMP-7 and IMP-8 Be data separately, and determining that the two resulting data sets were entirely consistent with each other, we combined them for further analysis. Figure 2(a) shows the sulfation of the IMP-7/8 flight data in which each event has been assigned a mass on the basis of its distance from the Be7 track center. We emphasize that the calibration of the mass scale in the figure is based on the Bevatron calibration results. A sample of the calibration data is shown for comparison in figure 2(b). 3. 1976 Data. Due to the very long collection times required for the Be flux measurements, instrument stability is a necessary requirement. The stability in both IMP instruments was monitored continuously in this study through (a) the use of an on-board electronic pulser, (b) the measurement of the heavily-populated carbon track center in successive 2-month periods, and 309 (c) the measurement of the beryllium track centers over \ 6-month periods. Etch Isotope Analysts Mass technique Independently Histogram of Cosmic Ruy B« demonstrated that Instrument drifts that could affect the (IMP 7 + IMP в 19731976 dalo) beryllium track location were less than +0.1 AMU over the entire 3-year collection (o) period. Therefore, 1n the analysis of 1976 data, we used the same track center locations and spacings as in the studies In Paper II. We emphasize that np_ correc tions have been applied to the data for arjy_ change in instrument response over the course of these measure ments. цтта i m и п fintiut Ю II 12 •S70 Figure 3 shows the iqass Be Calibration of the IMP 7/8 Back-up Telescope 120 histogram of the new data •geo ol IheLBL П Bevatron includpd in this study, 100 collected from 3965 hours of IMP-8 operation between 80 December 18, 1975 and October 15, 1976, and from 60 2636 hours of IMP-7 opera tion between December 20, 40 1975 and July 5, 1976. (The IMP-7 data collection stops 20 at a relatively early date due to a failure of the D5 7 8 9 " !0~ PMT high-voltage after the {> Atomic Mass Units spacecraft passed through an Earth shadow on July 9, 1976.) Figure 2 Due to the obviously low flux level of Be10 in thepresenc e of a large flux of Be9, as shown in figure 2, a distinctive Be10 peak could not be observed and therefore maximum likelihood fitting techniques were used to determine numbers of events of each isotope in the mass histograms. Figure 4 shows the maximum likelihood fit to the combined 1973 - 1976 data. Note that the distributions give an excellent fit to the observations. In particular, the fitting procedure yielded a separation between the Be' and Be9 peaks of 2.01 +_ 0.04 AMU—which strikingly shows the degree of agreement between the IMP-7/8 flight data and the calibration of the back-up Instrument at the Bevatron. This justifies our original assumption that the back-up instrument has the same relative response as the flight instruments. A summary of the number of Be events measured during 1976, and during the entire 34800 hours of collection between 1973 1976 is given in Table 1. The quoted errors include uncertainties in the background correction as well as the statistical uncertainties. Note that 310 • TTTTHTTT WTTT'WTI I W in'WII I I W I HWH ШЛ IfeMftMC MtMm MkttttTMn (IMPT+tMM tt7*Owo) i Jl i. 1 I» A I' I o I •/.„V luiiiulnuniuill"'"I r •> » L« 111 мт i 11 Hi i I » к> и a is Atomic MOM Unit» Atomic IteM UMN Figure 3 Figure 4 the background level of ^ 6 events/AMU Is extremely low. The beryllium fluxes and isotopic composition given in Table 1 are in good agreement with the previous reports in Papers I and II. The Be'/Be ratio measured here is also in good agreement with the reports of Hagen et a]^ (1977) and Preszler e_t al. (1975), while the Be'O/Be ratio given here is less than half of that measured by Preszler et aj (1975) and less than a third of the measurement of Hagen et_ al. (19777. "n view of the size of the quoted errors of these measurements, however, such discrepancies may not be very meaningful. TULE 1 SUMMARY OF IMP-7/8 IEHYLLIUM MEASUREMENTS Flux (x 105) (Pert/n2 ifster Number of Events Number of Events x sec x HeV/nuc) I* Isotoptc Energy Interval Observed Observed (at 80 MeV/nuc) Composition (MeV/nuc) (1976 Ditl) (1973-1976 Data) (1973-1976 Data) (1973-197» Oata) IMP-7 IHP-B (IHP-7 • 1MP-8) (IMP-7 • 39-132 90 + 10 425 + 21 17.6 • l.S 16.S • 1.4 0.617 1 0.01» *Be 34-131 34-119 215 + 16 8.4 + 1.1 10.8 + 1.2 0.352 + 0.020 "Be 31-123 31-108 5 + 3 17 + 7 1.2 + 0.7 0.6 + 0.05 0.03O + 0.01» Interval (AMU) Background 15 + 6 57 + 12 311 4; Node! Calculations. In order to interpret the Be10 abundance measurement, we ysed a homogeneous steady-state model of cosmic ray propa gation with an exponential path length distribution with a 6 g/ся2 escape яеап free path. Details on our calculations and assumptions are contained In Paper II. Energy-dependent cross-sections for the reactions were calculated using the algorithms of Sllberberg and Tsao (1973a), although for key production reactions for LI, Be and В Isotopes we used experimental results summarized by Sllberberg and Tsao (1973b). In the particular case of the light Isotopes, many of the cross-sections were measured by the Orsay group (e.g. Ralsbeck and Ylou 1973; Fontes 1976). Differential energy spectra at the'source were assumed to be of the form J « A(T+400)-2-°, where A 1s a constant and T Is the kinetic energy In MeV per nucleon. The calcu lations were carried out using a computer program written by J. D. Anglin. All cosmic ray isotopes from He* through S32 were Included In the calculation, which took account of energy loss by Ionization, secondary, tertiary and high-order production processes and radioactive decay including electron- capture effects. A series of such calculations was made, varying the assumed interstellar gas density over a wide range. Because ofrthe low energy of our measurements it was necessary to include effects of solar modulation, including convection, diffusion and energy loss. The local interstellar cosmic ray spectra calculated above were modulated using the Parker (1965) Fokker-Planck equation solved numerically using the technique of Fisk (1971). The overall level of modulation was deduced from the 1973 electron observations of Caldwell et aJL (1975) at earth orbit, taken together with the local interstellar electron spectrum calculated by Cunnings et al. (1973). For the Be data, we used modulation levels that fit the low energy proton and helium spectra measured at 1 AU during 1975 (Garcla-Munoz et aJL_ 1977a). The Be abundances for different Interstellar densities calculated with the above Interstellar propagation model were thus modulated to 1 AU, and then compared with the IMP-7/8 observations for both the 1976 and the 1973 - 1976 data. The-densities which fit our observations are summarized in Table 2. The deduced cosmic ray lifetimes in the table are simply the time required for particles moving at the speed of light to traverse 6 g/cm2 of Interstellar material of solar system composition of the deduced density. TAKE 2 DEDUCED MTEKTELUW DENSITIES AND С05ШС-Ш LIFETIMES Bt Isotoplc Composition (%) InttrsttlUr Density Cosmic Ray Ltfitla» OJti is! as! Ssl! (Atoms cm ) (VMM) 1976 Datl 0.673 0.397 0.030 +0 18 + 6 •0.041 • 0.041 + 0.016 0 17 - 18 ?£ x 10 "•''-o.io 0 19 1973-1976 D»U 6.617 0.352 0.030 0 IS* - 16-f| x 106 • 0.18 • 0.020 • 0.014 Jl. 5. Discussion. The average Interstellar density found 1n this work is well below the value of 1 atom/cm3 for the galactic disk, which ha* been confirmed by a number of recent measurements (e.g. Jenkins 1976). Thus, the present results could be consistent with cosmic ray propagation in a galactic halo, or In regions of very low «utter density within the disc (e.g. Cox and Smith 1976; Scott 197S). It Is also Interesting to note that the long cosmic ray lifetime derived here Is consistent with the measurement and Interpretation of high-energy electron spectra by Meegan »nd Earl (1975) •nd Hartmann et aU. (1977). 6; Ac know 1 cdgements. The University of Chicago IMP-7/8 experiments were constructed, and the data processed by the staff of the Laboratory for Astrophysics and Space Research. We are grateful to J. D. Anglln for writing the maximum-likelihood fitting program used here, and to T. G. Guzlk and C. DeGrazia for computational help. We particularly thank the Lawrence Berkeley Laboratory and the members of the H. Heckman group for their assis tance In the calibration of our back-up Instrument. This work was supported in part by NASA contract NAS 5-11067 and grant NGL 14-001-006 and NSF grant ATM 75-20407. References. Caldwell, 0., Evenson, P., Jondan, S. and Meyer, P. 1975, 14th Int'l. Cosmic Ray Conference, 3, 1000. Cox, D. P. and Smith,~B. W. 1976, Ap. J., 203, 361. Cummings, A. C, Stone, E. C. and Vogt, R. E. 1973, 13th Int'l. Cosmic Ray Conference, ]_, 340. Fisk, L. A. 1971, J. Geophys. Res. Letters. ]£, 221. Fontes, P. 1976, submitted to Phys. Rev. Garcia-Munoz, M., Mason, G. M. and Simpson, J. A. 1975, Ap. J. (Letters), 201, LI 41. 1977a, Ap. J.. 213, 263. 1977b, suBmTtteTTo Ap. J. , Hagen, F. A., Fisher, A. J. and Ormes, J. F. 1977, Ap. J., 212, 262. Hartmann, G., MUller, D. and Prince T. 1977, submitted to PTjyiT. Rev. Let. Jenkins, E. B. 1976, Goddard Space Flight Center publicationX-662-76-154, p. 239. Meegan, C. A. and Earl, J. A. 1975, Ap. J., 197, 219. Parker, E. N. 1965, Planet, and Space Sci., T17 9. Preszler, A. M., Kish, J. C, Leznlak, J. A., Simpson, G. and Webber, W. R. 1975, 14th Int'l. Cosmic Ray Conference, J2_, 4096. Raisbeck, G. and Yiou, F. 1973, 13th Int'l. Cosmic Ray Conference, X» 4M. Scott. J. S. 1975, Nature. 258, 58. Silberberg, R. and Tsao, C.~H7 1973*, Ap. J. (Suppl.). 25, 315. 1973b, Naval Research Lab., Washington, D.cT Report #7593. 313 Measurement of the isotopic composition of the primary cosmic radiation for the elements B-Nc C. Bjarle, N-Y. HerrstrOm, L. Jacobsson, G. Jonsson and K. Kristiansson. Department of Physics, University of Lund, Sttlvegatan 14, S-223 62 LUND, Sweden. The results are given from an investigation of the isotopic composition of primary cosmic ray B, C, N and 0. Preliminary re sult is also given from an investigation of Ne. The mass mea surements are made in nuclear emulsions exposed at about 3 g/cm2 atmospheric depth. The results for B-0 represented as quotients extrapolated to the top of the atmosphere, are: '"B/B - 0.61 ± 0.10; ,3C/C = 0.06 ± 0.03 15N/N = 0.33 ± 0.09; 170/0 = 0.05 ± 0.03; "0/0 = 0.08 ± 0.03. The preliminary result from the Ne-measurements shows that nuclei with masses larger than 20 exist among the primary neon nuclei. 1. Introduction. In this report we discuss the final results of an investigation of the isotopic composition of primary cosmic ray B, C, N and 0. We also give a preliminary result of a Ne-me- asurement. As most other investigations, on the mass content of cosmic ray nuclei are made with counter telescope,it is necessaTy to make measurements in independent detector media. Our detector, an Ilford G5 nuclear emulsion stack, was exposed 1970 over Canada. The energies of the nuclei fall in the interval 200-500 MeV/n. 2. Mass determination. The masses of the nuclei have been deter mined from measurements of the relation between track width and residual range on the last 12 mm of the stopping nuclei. The mea surements are made with a slit photometer. A detailed description of the present measurements of В and C is given by Bjarle and Herrstrdm 1976 and of the N and 0 measure ments by Jacchsison 1977. The measuring technique has been exten sively described by Jdnsson et al. 1970 in connection with a pre vious study of the isotopic composition of carbon. The mass re solution is 0.42 AMU for the carbon measurement and 0.50 AMU for the oxygen measurements. 314 5. Results. Our results for B-0 «re shown in Figures 1-4. The Figures show the mass spectra at the exposure depth in the atmo sphere. Figure 1. Boron spectrim t.N RELATIVE MASS 12 ISOTOPE MASS Figure -?. Carbon spectrum 1.10 RELATIVE MASS 13 ISOTOPE MASS 315 Figure J. Nitrogen spectrum Figure 4. Oxygen spectrum H IS Figure 5. Preliminary- neon .spectrum ' 316 Table 1 shows the estimated number of nuclides in the spectra of the Figures 1-4 and the isotopic quotients extrapolated to the top of the atmosphere. Table 1 of Extrapo luted Mass Number particl es quotient Boron 10 14 : l 11 19 B • 0.61 t 0,.1 0 12 1 Carbon 11 3 13c 12 74 c = 0.06 ± 0.,0 3 13 6 Nitrogen 14 16 15N = 0.33 ± 0. 09 15 9 Oxygen 14+15 4 i7n 16 66 -^=0.05 ±0.03 17 4 .,0 18 6 -gS-0.08 ±0.03 The result for boron indicates that the boron quotient for the radiation is somewhat smaller than the corresponding quotient for the solar system. An extrapolation of the carbon quotient to the source (Tsao, Shapiro and Silberberg 1973) shows that the abundance of l3C in the source must be very small - less than 54 on the 951 confiden ce level. The nitrogen quotient supports the assumption that nitrogen is present in the source. The nitrogen abundance in the CNO group is about 5* in the source. Our l70+1J0 abundance is somewhat larger than expected from the calculations by Tsao, Shapiro and Silberberg (1973). The difference is, however, not statistically significant. -The Figure 5 shows the preliminary neon spectrum at the point of measurement. We have fixed the mass scale on the bases of the knowledge of - 1) The distance in the spectrum between the different isotopes 317 Z) The mass resolution. 3) That only small amounts of ' *Ne and "Ne can be produced in the atmosphere. A preliminary extrapolation to the top of the atmosphere gives the quotient "Ne/Ne • 0.87. 4.. Acknowledgment. We are grateful to the Office of Naval Re search, Washington D.C. for the balloon exposure of the stack and to the Swedish Natural Science Research Council for financial sup port. 5. References. Bjarle, C. and HerrstrOm, N-Y.: Cosmic Ray Physics Report LUIP-CR-76-01, Lund, Sweden (1976). Jacobsson, L.: Cosmic and Subatomic Physics Report LUNFD6/NFFK-3001(1977). Jonsson, G., Kristiansson, K. and Malmqvist, L.: Nucl. Instr. Meth. 8±, 282 (1970). Tsao, C.H., Shapiro, M.M. and Silberberg, R.: 13th Int. Cosmic Ray Conf., Conf. Papers, Denver, Col. USA, Vol. 1, p.107 (1973). 318 NUN MASS OF COSMIC RAY Nt. Kg. SI at 1.2 G*V/«u Robert Owyer and Peter Htyer Enrico Fermi Institute and DepartmenCof Physics University of Chicago Chicago, Illinois 60637 (USA) We nave measured the mean «ass of comic ray He, Mg, and SI In the energy range of 800-1800 HeV/amu using a technique employing the effect of the geomagnetic field on the cosmic ray fluxes. The values for the neutron excess, X - 2Z, at the top of the atmosphere, are 0.45 t .10, 0.32 ± .11, 0.26 i .16 for Ne, Mg and SI, respectively. These results, when account Is taken of the effects of galactic propagation, imply values for the neutron excess 1n the cosmic ray source which are In agreement with solar system values. 1. Introduction. One of the crucial experimental tests of models for .nucleosynthesis, In particular explosive nucleosynthesis, 1s the isotoplc composition of the elements that constitute the cosmic radiation. These elements have different origins. Some are produced by the cosmic ray sources and, after modification through Interactions in interstellar space, are detected at earth. Others are of secondary origin and owe their existence In the cosmic rays almost entirely to spallation processes in the interstellar medium. Me have previously reported the isotoplc coroositlon of nitrogen (Owyer and Meyer 1975), an element of predominantly secondary origin. In this paper we focus on the isotoplc composition of some of the source nuclei, and present measurements of the mean mass of the elements Ne, Mg and Si at energies of about 1.2 GeV/amu. The method that we use for this work has been described by Peters (1974). It employs an accurate velocity measurement in the penumbra, the region of the geomagnetic cutoff where the earth's field 1s partially transmitting. A sample of nuclei filtered by the geomagnetic field (according to particle rigidity) will show differences 1n their velocity spectra depending on the mean A/Z of the element. Elements having isotopes with larger neutron excess (A - 2Z) will show lower velocities allowed by the earth's field. 2. Instrumentation, Balloon Flights and Data Analysis. We have built and flown an experiment (Figure 1) designed to use this differential geo magnetic method for the determination of the mean mass of elements in the cosmic rays. Our results were obtained in two balloon flights from Palestine, Texas. The exposure factor for the flights was 20 m* sr hr under a residual atmosphere of 3-5 g/cm2. In this work the velocity measurement 1s achieved by a liquid Cerenkov counter (T2) having an effective index of refraction of 1.29 and hence sensitive to energies in the range 600 MeV/amu to several GeV/amu. In addition, the counter telescope (Figure 1) contains two Identical plastic scintillators, Tl and T3, for Identifying the charge of the particle and rejecting nuclear Interactions 1n tht Instrument. A multlwlre proportional chamber hodoscope (MWPC A through C) 1s used to Identify the particle trajectory and reject background events. The guard counter G serves as an 319 additional «Id to reduce the background. To the m diti, three selection criteria tre applied to generate the sample of non- Interactlng nuclei used In this analysis: 1) a unique straight line In the MUPC hodoscope 2) agreement between top and bottom sc'^illator pulse GUARD COUHTCK heights wi » 8-91' for Ne, G 2" P.M. TUBES Mg, SI. 0 25cm. 3) low guard counter signal. Figure 1 A prerequisite for the deter mination of Isotoplc composi Schematic cross-section of the instru tion 1s a complete separation ment. MUPC is multlwlre proportional of the elements to such a chamber. degree that each event can unambiguously be assigned to a specific nuclear charge. That this 1s accomplished after the application of the three selection criteria is shown in Figure 2 which gives the charge spectrum obtained for the elements of interest here. Figure 3 shows the Cerenkov signal spectrum of magnesium and oxygen. The spectrum of the reference element oxygen (assumed for the moment to be all 1Б0, A/Z = 2.0) as measured in the Cerenkov counter is used as a tem plate for any desired mean mass. Thus, also shown in Figure 3 Is the cal culated spectrum, derived from the oxygen distribution, which represents what would be seen for Mg if this element consisted entirely of ">Mg (the original oxygen distribution corresponding to A/Z = 2.0 represents pure 24Mg). The average mass of Mg over the energy range oi the penumbra (800 - 1400 MeV/amu in the figure) 1s obtained by a weighted average of the posi tion of the actual Mg points relative to the two reference spectra for each channel in the 'erenkov signal distribution, using the method of maximum likelihood. "T- i i i 150 While the details of the т analysis will be published Uni t i elsewhere, we wish to summarize &I00 • 1 the corrections that need to be t applied to the data and the « \ tests that have been made to 1• Sm f X - ensure the absence of system UJ ii atic errors. An important i „L {} correction involves the slowing ,1n, ^lillillUnj- 0 1 JiiiXiL J i F Ne No Nlg 1II L .Al jlLiSi i . down of the cosmic rays 1n the atmosphere; however, this cor rection can be made very accu rately. Since silicon slows Figure 2 down more than oxygen, the mean mass of Si relative to 0, Charge histogram of a portion of measured at balloon altitude is flight data after selection criteria 0.5 amu higher than after cor were applied but without atmospheric rection to the top of the correction. 320 Energy. MeV/amu 640 800 1080 1750 T 1 1 T 1 1 T I I 1 - o Oxygen, A/Z«2.0 5000 500 " • Mognesium . | . • • • - * Oxygen. A/Z« » 1 " -Ю00 100 26/12 A * • • c А» в • *> A o юо| 10 A t - A 0 § m A 1 • 10 | O Д o 0 A o o Q Д - 1 1 1 1 1 1 1 1 1 1 02 0.4 0.6 0.8 1.0 1.2 Cerenkov Signal/Z2 M067 Figure 3 Cerenkov signal spectra for Mg and 0 observed In the first flight. These spectra are Influenced by the geomagnetic field for signals C/Zz £ 0.7. Note the logarithmic vertical scale. atmosphere. For Mg and He, this effect Is oroportlonately less. The un certainty 1n the cror recti on 1s less than 10X (0.05 ami) since It depends on the measurement of residual pressure and only slightly on the measurement of velocity» bo^h accurately known. The accuracy of the velocity measurement depends on the resolution of the Cerenkov counter which 1s limited by the number of photoelectrons detected. Since In our Cerenkov counter we measure 48 liberated photo- electrons for a Z - 1, е • 1 particle, there 1s no deconvolutlon correction required for signals (C/ZZ) below 90X of that of a relat1v1st1c nucleus for the elements considered here. Since our .results would be sensitive to any charge dependent velocity bias, we have carefully checked that our three selection criteria do not Involve such a bias. By repeating the Isotoplc analysis over a range of values for the selection criteria, we have verified that even 1f the criteria are tightened to such a degree as to eliminate almost 50* of the sample of nuclei, the resulting mean mass changes by only .07 ami for SI and less, than .04 for Ne and Mg. Also-we make a correction for spallation 321 In the atmosphere, although for these abundant cv«n Z elements, the correc tion Itself is at most .04 amu and the uncertainty smaller. Finally In the analysis of our results ме have assumed the Isotoplc abundance ratio «60: Щ: »8o - *6:2:2 (Shapiro et ajLj_ 1975 and 1976) at the top of the atmosphere. In Figure 3 the Ng Cerenkov spectrum has been normalized by the Mg/0 ratio observed In our balloon flights measureo et an energy above the effect! of the geoaagnetlc ftfld (above about 1.5 Gev/amu for the data 1n Figure 3). That this Is the correct ratio to us» within the penumbra (800-1400 MeV/anu where ме cannot measure the Interplanetary ratio) Is not Immediately obvious. In our first balloorr.fHght, the gondola drifted from Palestine, Texas (vertical cutoff rigidity 4.5 GV) northeast to Arkansas (cutoff of 3.3'fV) so that» using data from the low cutoff region we may obtain the chemical ratios of N», Mg, SI relative to 0 down to t 1250 MeV/ amu. (In the second flight, the balloon drifted southwest of Palestine and the penumbra1 /range was from approximately 1200 to 1800 HeV/amu.) Other recent experiments using plastic Cerenkov detectors In counter telescopes have measured these energy spectra In the range 0.5 to greater than 2 GeV/ amu (Jullusson and Meyer 1975, Haehl et aU. 1976). Although the Ne/0 ratio seems to be constant, the Mg/0 and S170* ratios are possibly increasing with Increasing energy wltfvin this energy range. We have used the results of Jullusson and Meyer (1975) which show a slight variation in the Mg/0 ratio and a somewhat greater Increase In the S1/0 ratio, to account for these ratios in the 800-1250 MeV/anu region which Is Inaccessible to our measure ment. The results of Maehl et^aJL^ (1976) are consistent with anything between a variation slightly stronger than that of Jullusson and Meyer and no variation at all. It should be noted that the difference in the results for the neutron excess between the case we have chosen and the assumption of energy independent chemical abundance ratios is .04 amu for Mg and .09 amu for Si. 3,j Results. In Table 1 we show the results for the mean neutron excess, ff - 2Z, measured in this experiment and corrected to the top of the atmosphere. Me nave made a calculation to see what these measurements imply for the Isotoplc composition of the cosmic ray sources. The model of galactic propagation and confinement we have assumed Is a steady-state model with an exponential distribution of particle pathlengths, and with a mean pathlength of 6.0 g/cm2 of interstellar material (composition: H with 10* He). This calculation uses the partial fragmentation cross-sections calculated witH the semi-empirical formula of Sllberberg arid Tsao (1973). The elemental composition assumed for the source Is that calculated by Shapiro et al_i (1975). For the abundant even Z nuclei considered here the arriving flux is dominated by the surviving source nuclei relative to the secondary compo nent produced by spallation en route. We wish to note that even in the case of neon (which has the largest secondary component In the arriving flux of the three elements) the fragmentation of parent nuclei (Mg, Si,... Fe) produces each of the three stable nuclei with essentially equal proba bility and does so almost Independent of the parent element's isotopic composition. In Table 1 we compare our results for the Implied neutron excess of the elements Ne, Mg and SI at the cosmic ray sources with the solar system abundances of Cameron (1973). There is no Indication that for any of these 322 TABLE 1 - Results for the mean neutron excess, X - 21, observed tn thti experiment and the Implied results tor the cosmic ray source composition compared with the values for solar system material. This Experiment Extrapolated to the Solar System (top of the atmosphere) Cosmic Ray Source (Cameron, 1973) Neon 0.45 t .10 0.20 • .14 0.22 Magnesium n.32 t .11 0.21 i .13 0.33 Silicon 0.26 ± .16 0.19 • .17 0.11 nuclear species the neutron excess is different from that found in the solar system. This conclusion differs from that of Fisher et a_l_^ (1976) who observed a substantial neutron excess for Ne (0.8 ± 0.lT~at about 450 •MeV/amu which they interpreted as requiring the cosmic ray source neon to be considerably more neutron rich than solar system material. In prelim inary results of Preszler et а_1_^ (1975) also at lower energy (150-450 MeV/ amu), the neutron excess oT~neon was measured to be 0.60 ± .14. The results of Fisher et aK_ (1976) for Mg and Si (0.3 ± .1 and 0.1 ± .1, respectively) are Tn agreement with our higher energy results and show no need for a source composition different from solar system values. Thus even though cosmic rays are a sample of matter which has under gone a unique history resulting in acceleration to high energies, our results show that, after account was taken of the effect of galactic propagation, the neutron excess of Ne, Mg and Si is remarkably similar to solar system values. This work demonstrates the value of using the geo magnetic field in studies of the cosmic ray isotopic composition at ener gies of 1 GeV/amu or higher. Further progress in identifying individual isotopes will give more information on the nature of the sources that produce the cosmic radiation. 4. Acknowledgements. We are grateful to Messrs. G. Kelderhouse, W. C. Johnson, G. Mlnagawa, W. Hollis, L. Glennie, L. Littleton, N. Beck for various important contributions to this experiment, and to J. Caldwell for many useful discussions. The program to evaluate the effects of cosmic ray propagation was kindly made available to us by Dr. J. D. Anglin. We wish to thank Dr. T. Parnell for generously sharing with us his experience with the design and construction of multiwire proportional chambers adapted for balloon borne experiments. Research was supported, in part, by the National Aeronautics and Space Administration under Grant #NGL 14-001-005. 5. References Cameron, A. G. W. 1973, Sp. Sci. Rev., 15, 121. Dwyer, R. and Meyer, P. 1975, Phys. Rev. Lett., 35_, 601. Fisher, A. J., Hagen, F. A., Maehl, R. C, Ormes, J. F., Arens, J. F. 1976, Ap. J.. 205, 938. Jullusson, E. and Meyer, P. 1975, Proc. 14th Int'l. Cosmic Ray Conference, Munich, 1, 256, and private communication. Maehl, R. C,"Ormes, J. F., Fisher, A. J., Hagen, F. A. 1976, Goddard Space Flight Center preprint Xt661-76-132. 323 Peters. В. 1974. Nucl. Instrum. Methods. 121, 206. Preszler, A. N.. Klsh, J. C, Leznlak, J.~JC, Stepson, G. and Webber, U. R. 1975, Proc. 14th Int'1. Cosnlc Ray Conference, Munich, 12. 4096. Shapiro, N. N.. Sllberberg, R. and Tsao, C. H. 1975. Proc. 14Th Int'l. Cosalc Ray Conference, Munich, 2, 532. Shapiro, N. N.. Sllberberg, R. and T?ao, C. H. 1976, private conwnlcatlon. Sllberberg, R. and Tsao, C. H. 1973, Ap. J. (Suppl.). 25. 315 and 335. 324 A ftmdy ©Г Оош1о Ray Chart* * I ectopic Composition Using « Charged Partial* Telescope •atya Dev Tama Department of Physics №1 varsity School of Solano*a Oujarat University Ahmedabad 380 00» XXDU A charged particle talaaoopa la used to atady low and medium chare*« primary ooamio ray alaaantal and laotopic ooapo - sltlon. 1h* apparataa vaa flown in upper atmosphere, aboard a balloon flight oondu- etad over Port Churchill, Canada* In 1973. The iw topic raaolntlon of the appa ratus vaa evaluated by calibrating It by (1) Monoenergetie relatlviatlc 0» ions and (ID cosale ray amons. Fart of the data Is analysed and isotopic composition of Be Is discussed, though with poor statistics. 1. Introduction. The elemental abundance of low' Z ale - mental U, Be, В In cosmic ray nuclei are assumed to have been prodqaed In the nuclear collisions of their parent nuclei with the .Interstellar matter. To obtain the galactic confi nement and propagation life time of cosmic rays, knowledge of average density of Interstellar matter, which is not directly measurable, ls 9{. B-rparl^ant. Details of the apparatus used are given in earlier papers £Verma and Herzo, 3973} Verma ft Norokauer, 1975). The apparatus is essentially a charged particle telescope,-designed and fabricated at LSU, Baton Rouge. It 32S Is briefly described her* *nd shown scheartleally In Щ.1 The chert» of the Incident particles Is measured In the plastic scintillation dete ctors Dlt F1 and a Luclte oerenkor detector Fg. Thin SPARK CHIMKRS plastic detectors Dg, D3 4 0^ war* used as rang* counters. I -1 ] 0, Total energy of the particles, P, |- stopping in Csl (TI) crystal ~f" 1 detectors» was determined in 0 ? ? « INCHES F3 and P4. Both Csl (Tl) Si (Li) K>M «let» d»»««loc (mcMKh P| detectors were one inch thick. инчцич lucita Five 3000 channel pulse height rmrnm Plottic KinlilMion dtMCteri D,- 0. analysers were used to record C«l (Tl) cr»» rence Papers)Vol.lt p.331 (Presented at the 14th Int.Conf, on Cosmic Rays, Munich) 1975). Verma,S. D., and Негво, D. 1973, Conference Papers Vol.4 p.2883 (Presented at the 18th Int.Conf. on Cosmic Rays, Denver 1973). Verna.S.D., and Norckauer, Н.Л, 1975, Conference Paper Vol.9, p.3193. (Presented at the 14th Int.Conf.on Cosmic Rays, Munich, 1975). Webber V.R, Lezniak.JUL, Kish, J. and Cattle, S.Vn 1973 Astrophys. and Space Sol. £4., 17, * Some time after the balloon flights were performed, -author moved.from LSD, Baton Rouge, USA to Gujarat University, Ahmedabad, India, leaving part of the data behind. This data is yet to be obtained aad analysed. ii.irnowi•daamanta. Authors give thanks to Dr.D.Herzo, Stener Kleve, Don Riddle, K.Rost, Don Babcock and graduate students H.R.Norckaner and G.Seab for the help in carrying out the experiment. Dr.David Cheshire for PDP-8/K programming assistance and Dr.L.B.Levit for the help in the Monte Carlo Calculations. Thanks are due to the following agencies for the financial support. The research Corporation, cne LSU Council of Research, Office of Naval Research, National Aeronautics and Space Administration under grant NGR 19-001-012 and National Science Foundation under the Science 321 Dtvelopa*nt Prograaat through th* dapnrtaant of rnjrale» and Astronomy, LSD, Baton Root*. Thanks art also dut to Bastaan Kodak for donating file for photographing spark eh «bar and Ttxaa Instruments for donating alaotronlca ooaponanta for tha •zp*rlaant. 3» ISOTONIC COMPOSITION OF LOW ENERGY COSMIC RAY PARTICLES WITH CHARGES Z - 5 - в R. Beaujeen, H. Sageblel and V. Enge Inetltut fur Rein* und Angevandte Kernphyalk University of Kiel, 23 KIEL, Weat-Germeny A stack of 250 um Daleel oelluloae nitrate ahaata waa flown from Ft. Churchill under 2.7 g/cmz ataoapherlc depth for 13.25 hours. The charge and mass determination on 647 stopping particlea waa dona by means of the cona length тегаиа realdual' range method ualng the REL criterion. The isotoplc ratios are Measured In the top part of the stack and under additional 2 g/em| In the bottom part. Averaged numbers for C"/C «nd 01'+1в/0 are 7.4*1.9 and 15X1.6 respectively [%]. 1. Introduction We have shown that cellulose nitrate nuclear track detector ahaata can be used to measure the iaotopic composition of low energy cosmic ray particles even in the absence of a clear separation of adjacent isotopes (1,2). Hence we carried on our early measurements with an Improved evaluation technique on an enlarged number of particles. The analysis of more than 1000 stopping particles in the energy region 100-200 MeV/nuc was performed in a stack of Daleal cellulose nitrate which was flown from Ft. Churchill under 2.7 g/cm for 13.25 hours. The stack consists of 100 sheets of a size 250 /um x 20x40 cm . The upper 56 sheets (equivalent to 2 g/cm ) were etched at 40 °C In a 6n NaOH solution with 0.05 % Benax added for 240 min. The residual range and cone length measurements along the tracks were done under a Lejltz -optical microscope with a 100x objective. 5 So ?. Data evaluation The «canned atack volume was divided into two parte. On* part Includes particles •topping In eheeta 1 to 15 (further refered to aa the top part), the aecond part Include* those particles stopping between sheeta 41 to 56 (refered to aa the bottom part). These two aeta are aeparated by a mean matter of 9 ft 2 (/cm for particlea with Ui zenith angle. It waa deteoted during the analysis that the response of the detector sheets varies fro» the edge towarda the centre of the stack. The reason for thla behaviour la not cleared but may be due to a temperature effect since the edge was not completely protected against heating or cooling at celling level during the balloon flight. The main change in response occurs within 1 cm from the edge whereas the inner part shows a reasonable constant response. As a compromise between a large number of particles and a somewhat constant response on a single detector sheet all particles were rejected, from the following analysis with tracks in the area 2 cm or closer to the edge. Thus 75 % of the particles remained within an area of 16x56 cm2. Furthermore a different response was detected for the top and bottom parte which can be easily seen in the cone length versus residual range diagrams of the accepted particles (Fig. 1+2). Further studies are in progress to clear whether a temperature gradient or a diffusion process causes such spatial changes in the response. J» Fig. 1+2. Cone length versus residual range measurements in the top (left) and bottom part of the stack. ; 3. Results | The charge and mass identification of the accepted particles | with nuclear charge Z=5-8 was done following the procedure | described in ref. (3). The track etching rates for the top ; and bottom parts were calibrated using the predominant peaks of C and 0 . The mass histograms in Pig. 3 and U are based mainly on single cone events. Prom some available multi-cone events a fluctuation of 0.6 amu for the single cones was deduced. To extract the isotopic abundances from Fig. 3 and U the following methods were applied. The amount of secondary 11 15 C and 0 produced in the matter above and in the stack was calculated (4,5) and substracted on the left aide of the corresponding distributions. Assuming symmetric Gaussian : distributions for C12 and 016 the amount of C13 and 017+18 is thus obtained as the excess on the right side. For Boron and 332 Nitrogen the aesn «ass of the distribution yields the portion of B10, B11, N1** and N15 (the calculated- number of N1' was taken Into consideration ). J 1 ijn ijj AJ.,..I Fig. 3+4. Mass histogram for the top (left) and bottom part. The number of particles per 0.2 amu is plotted. Table I a) calculated numbers (#) C11/C N13/N 015/0 top 3 i 1.8 0.7 - 0.7 2.2 * 0.9 bottom U - 1.7 1.3 - 1.3 3.7 - 1.3 b) measured numbers (90 B11/B C13/C N15/N 017+18/0 top 70 i 26 7.7 i 3 40.2 t 8.9 16.9 * 2.5 bottom 78^ 22 7.2 ±2.3 47 ± 8.5 13.1 ± 2.5 4. Dlacuaalon 3)9 Table I euanarisee the calculated number* of the aecondarlea and the resulting numbera of the coaalo ray particles. All number* sr* valid st detector level. The eosale ray abundaacee sre In good agreement vlth results rerleved by J.P. Meyer (6). The amount of C* can be explained by interstellar and ataoapherio fragnantations and thua any evidence for a dlaturbed CMO-cycle eannot be deduced. However, the neaaured 017+18/0 ratio is larger than expected froa pur* fragnentatlon processes. But this does not contradict the conclusion drawn froa the C15 result as 018 aay originate froa • different burning. Acknowledgment We like to thank Prof. Dr. S. Bagge for hie support and the staff of Raven Ind. for performing the balloon flight. Thla work was supported by the Deutsche Porachungsgeaeinachaft. References 1) R. Beaujean and W. Enge, Zeitachr. Phyalk 356. (1972),416 2) K. Pukui, V. Enge and R. BeauJean, Z. Physlk A227 (1976), 99 3) R. Beaujean and W. Enge, Fragmentation and isotope measure ments on accelerator Neon and Argon particles, Nuc. Track Detection ±, (1977), 19 4) R. Silberberg and C.H. Tsao, Aatroph. J. Suppl. Ho.220(1), 25. (1973) 315 5) P.J. Lindstrom, D.E. Creiner, H.H. Heckman, B. Cork and F.S. Bieser, Lawrence Berkeley Lab. report LBL-3650 (1975) 6) J.P. Meyer, Rapp. paper 14th ICRC MUnchen 1975, p. 3698 3J4 Fa - Xaetepee la Ceemic Hay a G. Siegmea, K.P. Bartheleaa end W. «ме laatitut fllr Belae uad Aagevaadte KeraBayalk, Omlverelty of Kiel, 23 KIEL, Fad. ftepwblle ef German? Abatract: Uelng tha ultraviolet irradiation af Ьежаа pleycarbenete plaatlo far tha eahancemaat ef мм resolution tba laeteplc ceapoeition of eoaale ray auolal waa meaaured. An UV- expo- aura equipment vaa built, vbieh anablaa aa ta irradiate 12 600 em2 of plaatic ia ana ram. A u« method of particla identification vaa evaluated. Tha accuracy af ma** determl- natian will ba diaeuaaad. la thla experiment we eatlmat* a raaalutioa of ±0.7 aam for tha iraa region. Baaed upon the aeaeureaent of about 530 particle* vlth anergica of 150 - 500 N*V/nuc tha laotopic distribution af Pa-auclai will be preaeated. 1. Introduction Much verk baa been devoted to tha atudy of tha char«a coapoaitien of coaaic radiation. Apart froa a lot af phyai- cal proeaaaaa that aay tranafora eoaale ray coapoeltien during» their propagation in apaee the elementally determin ed abundance pattern of the eleaeata baa conatituted the baae for theoriee of nueleoayntheeia. In thla regard leoto- pic ratioa are of particular importance becauee theT are inaenaitive to chaaleal aeparation effect*. Ve atuc ad the iaotopea of iron in aore detail bacauae thia 1* an element vhcae modification by aecondary proceaaea of heavier eleaenta baa been email and which 1* heavy enough to have very aimilar value* «f A/Z. 2. The Method To Improve the reaolution capability of the detector ve used the UV-enhancement effect (Siegmon et al., 1975). Therefore ve built a UV-equipaent which allowe the aenaita- tien of aa area of 12 600 cm2 plaatie in one run vlth a homogeneity better than 0.5*>. The etching vaa don* in an ultra*onically agitated, mechanically atirrad eolution of 355 6.0*0.05 а ИаОМ vita O.05*>*oax aarfaotaat at 70*0.03*C. Ve ae*d the amltlble-dK/dx-R aethed by «•«••rime *»• etehed oot* leagtb. L •• • fmaotlea ef tb* roaldaal rang* • *f 0Y- •ahaac*d track*. Dopaadlag •• the u|lt, oharge aad atopping position *f the laeldeat particle up t» It e«a«a oeald be •eaaured. Th«a each •••at wa* repreaeated by ••••ral data palnta la the wall known LH-plat. la ardar to reduce the track data t* a alagl* paraaeter for «aeh event tba following pre- cadnra waa perforated! 1. The aata *f LB-data ef aaoh •••at wara fitted by вшава ef *ultabl* analytical function*. (Ta dlacrlalaata againat nuclear iataraotiena twa a*laction criteria wara waadt (a) value* of L which darlato by aora thaa 3/#a («2 ataadard darlatlaaa) fraa the fit war* alialnatad aad (b) at laaat k eon** aad 75)f ot th* ariginal con** araat b* good to accept th* track.) 2. Th* resulting track-fit* ware thaa integrated to get on* aingle paraaatar. A* only cane length* ware actually aeaaured th* integration wa* per forated over a fixed L-int*rval fL1i L2l which reaaina naelum(td for th* whole analysis. For practical raaaona thia -ralu* waa replaced by a a*an rang* R. Thua, by uaing th* LR-aaaaureaanta a* input data th* application of the described procedure lead* t* a aingle paraaeter for each event. Apart froa the adopted aelection criteria only the** •••nt* could be accepted whoме data aeta centaia cenea with length L>Lt and L<^2* In tot»\l 335 particle* ware l*ft| 1/3 wa* excluded froa analy*i*, aapaclally th* ahort track* of" nuclei with Z <16. To calibrate the R-histogram we aad* the following aasuaptiona: I. The highest charge nuclei of greatest abuadaac* ar* of lrion <«nd furtheraore th* a**t abundant, iron iaotope i* th* aueleeu* "bF*. II. The ми of the Si-aa**-dlatribution la I «28. 336 III. TIM nodal of the Reatrlatod Snorcy !«••• (»L) -« prepeaed by Bentea (l***) 1« feea* ta he an ad- eanate theory whieh deeerlbee the aedlfled energy leaa aa being relevant far the traek feral** »re- eeaaee. V* eheae til* adjaetabla parameter w# of vhe ML Medal to 1 OOO ••, • value waleh 1* eenaletent with the aaauaatleaa" I. and II. The preoeduro hew to fix the dlaeret aet of uirpi Z and aaaeea A at the eerroot I-pealtiena we deaeribe 1* Slecnen et nl. (197*). 1. Reenlta The hietofraa of 335 particle» of galaetle eeenio rays •oaaorod In this experiaent la proaoatad 1» ti.g.1. Tba nua- bare of tba aeala give fi in^a and the charges ara aaeigned aceerdlng to tha calibration. Ля ITCMIM i i itv MM M • t • • t • m • * t H • • • II n 10 41 iO U it IS unto 91 м it II U «I V11 H 1•• It M SI s Ar K Cct Sc » V Cr Mt ff» Co МГ Fig.1 : Spoctrua af SI (Z>1%) up to HI (Za28) far energies I - 150 - 500 HaV/nue. In Tab.I «a auaaariae tba eoaalo ray Iron naaa diatribu- J tlan absorbed undar 1.8 «/on atnaapbara. Tha results ara cam-/ parad with athar axparlaantara and with thaoratieal та1паа. Although thara ara aoaa diaerapaneiaa whloh any bo af ata- tiatlcal origin tha raanlta of tha diffarant experinenta agree within JOjt on an average. 337 Kiel Vetoer Clapfcam Meake •>•» Tk**r«tir*l 1.8 t/c»? 5 t/<«* 5.5 t/«"S »«/-«" T.ltf" l->-laolop« ***at* * •«•at* 1 * < <5?.55)-54 56 >в 15 2t ltt 2« 1* 50 Л.5 е (55)«5*»57 71 46 25 45 27* 4t 52 «0 50.T *М?9 SO) 2? t; t« 50 155 25 < 10 19*9 O.J Table 1. Jb_ There are several touroti of trror 1» the mass identlfl- cation procedure. The most obvious error source is duo to fluctuationa la tho cono leag-ths. Tha x- and s-projoctlona which woro dotormlaod with an aecuraey of 1>m yield aa error In tho Pe-region of 0.6 asm. Tho error of tho track-fit, in tho leaat square sense, ia ahown in fig. 2 . Tho expectation -value of the distribu tion loads to an error of pass determination which ia close to that derived from measurement errors (0.56 asm for Fe). Another aouroe of error originates in tho location of tho mass linos themselves. The uncertainty in tho cali bration at tho 5 Fe peak is believed to give an absolute value of 0.1 asm in the loca tion of the mass seals. This uncertainty will lead to a systematic displacement of the observed masa distribution relative to the calculated scale. The remaining- systema tic errors dopendon the adopt 1 a»** ed track theory RIL including PiC« 2 lta adjustible parameter w . Track-fit errors To eetlatate the established reeelutlea tbe lr»» pert »f the histogram waa fitted by a leaat-aquare technique ta 3 Gaussian dlatrlbutlona caatarad at atnti л»5*,5б,}$. The beat fit lndlcataa а mass raaolutlan ef 0.7 aam. 5. Concluelon Uaually coastc ray abundancea ara coapared with aalar abundancaa aa elaborated by Caaeron (1973)* I» order ta atudy differencea between aolar ayatea abundancea and aeesureetenta, a flrat approxlaation atarta with a aolar coapoaitlen at the • ource of galactic coaalc raya which ia than tranaforaed during it* propagation in apace before being detected. In the case of iron which ahowa coaparabla abundances of Fa and Fe there ia obviously a significant difference to a aolar aource coapoaition. Although the data have not been extrapolated to the top of tba ataoaphera it ia iapoaaible to explain the discrepancy with ataoapheric fragmentation*. Thus the Fa-aeaauraaent accounts for the existence of large fractions of. the iaotope Fe and, leas significant of Fa, in the coaaic ray aourcea theaaelves. Following the aaauapt- ion that super nova events are the aource of galactic coaaic rays, the aource coapoaition aay be regarded as a aignature of nucleoaynthetic processes In such eventa. Although there is a variety of poaaibla iron isotopic configurations in cosmic rays which have bee extensively diacuaaed by Voosley 54 (1976) an excess of Fe would imply that cosaic ray iron originates in an environment that was more neutronrich than the aource of solar iron.. The interpretation of Fe requiring 58 the ejection of an inner Fe-rich zone with a very high neutron excess ia ambiguous as well as difficult to measure. Acknowledgements The authors are grateful to Prof.E.Bagge for his support of the work. This work was supported by the Deutsche Forschunga- gemeinschaft. В.Гегапсе» 539 Bentea Е.V., A Study ef Charged Particle Track, a la CelltUeee Nitrate, U.S.Naval Itadleleclcal ОеГеа- aa Laberatery, Saa rraaclaoe, (19(a) Caaeroa A.O.V., Spaoa Sel. Кат. Ц, 121, (1973) Claphaa V.M.. Fowler P.H., O'Ceallalc* C.. 0>Sullivan 0., ТЬомеоп A., Hth CRC Munchea, eOO, (1975) Haoka R.P., Benton E.V., 14th СЯС MUnehen, 393, (1973) Siasmon G.-Г., Bartnoloaa K.-P., Enfe V. Nucl.Inetr.Math. I£8,a6l, (1973) Slacaon G.-F., Koahnan H.J., Bartholoaa K.-P., Knee V., The dapandanca of tha вааа-ldantiricatlan-acala on different track formation mod»lm, 9th Int. Conf. on Solid State Nucl. Track Detectore, MUnchan, in prap. (1976) Taao C.H., Shapiro м.М., Silberber*- H., 13th Int. Coaalc Ray Conf., Vol.1, 107, (1973) Vebbez V.R., Laznlak J.A., Kiah J., Nucl.Inatr.Math. 111. 301, (1973) Vooaley S.E., Aatrophya. Spaca Sci. 22» 103i (1976) 540 COSKIC RAY IRON MKASURFHEHTS MITH FRAGMENTATION CONTROLLED R. Scherrer, V. Rnge, R. BeauJ«an Inatitut fUr Rain* und Angevandt» Karnphyelk, Unlveralty of Kiel, 23 KIKT., Weat-Qermany S. Hertzman, K. Krlatianaaon, K. sederatroa Department of Phyalca, Uttivaraity of Lund, LUND, Sweden Theoretic»! • Exparim-m Q »° A datactor ayatem for mass determination of Iron group nuclai in tha energy interval 500 - 600 MeV/nuc. waa eon- atructad eonaiatlng of ealluloae nitrate plaatlc and nuclear aauleiona. One advantage of this detector system is the high rejection of Interacting particles. In two balloon flights It was exposed to cosmic ray. During tha evaluation range ami track width in emulsion and cone- lengths and ranges in plastics were measured. The identi fication parameters of the particles are determined inde pendently in emulsion and plastic. Results'of mass determination will be d iscussad. Cooniinatei: OG 1.6 Milling »ddre»: Dr. W. Enge Inatitut fUr Reine und Angewandte Kernphyslk Uhiversity of Kiel 2300 KIEL West-Germany MWIMMJ. Ml A STUDT OP THE BE/B RATIO AS A FUNCTION OF ENERGY AND ITS INrXI CATIONS FOR THE ACE OF COSMIC HATS V. K. Webber, J. A. Umiak, J. C. Klah and G. A. Slap*» Physic* Department university of New Hampshire Durham, N. H. 0382* Haasureaenta of tha Be/B ratio have bean made In the energy range 200 MeV/nuc to t20 OeV/nuc using a coaalc-ray taleecope which contalna a solid and a gas Cerenkov detector. Thaae •aasuraannts are compared with tha predicted le/I ratio based on current galactic propagation modela. The data agreea with the predictions for complete 10Be decay at energlea below -v4 GeV/nuc. At higher energies time dilation effecta lead to a larger fraction of 10B* surviving, thus resulting In an In creasing Be/B ratio. This may be used to estimate the coaalc- ray llfetlae. The measured Be/B ratio Indicates a coaalc-ray lifetime that la £l0 x Ю6 years. Introduction. In a recent paper (Webber tt at. , 1977) we reported a new aeas- ureaent of the abundance of the radioactive cosmlc-rai' v isotope 10, These l0 results Indicate that the surviving fraction, f, of. Bhe Is 0.20±0.09 In the energy range 145-245 MeV/nuc. Garcia Ktrioz eX at. (1975a, b) have obtained a very slailar value for this fraction at an average energy tlOO MeV/nuc. These measurements clearly Indicate that most of the '"Be has decayed; however, the Interpretation of this decay In terms of a cosmic-ray age is very dependent on the characteristics of the solar modulation. If no modulation whatsoever la assumed, the data is consistent with a mean age of %10хЮ6 years. Modulation effecta will tend to Increase the deduced age. The modulation effects are best illustrated by use of Figure 1, which summarizes the current *°Be data In terms of the cosmic-ray age. The values of f are thoae measured at earth and are not corrected for solar modulation effects. For the solar modulation we use the theory of Gleeson and Axford (1968). Two aspects of this modulation are of importance. One is the total modu lation as described by the parameter ф (in Mv) which represents the adia- batic energy loss suffered by the cosmic rays in the modulation region. The cosmic rays observed at earth EiwgylMV/MC) come from higher energy interstellar cosmic rays. At these higher ener Щиле. 1. Katationifup bvbw&n thi 10 gies the expected fraction of 10Be тгалилеЛ iuAviving {tuuMjan \ок 8e jtn the galactic radiation will be and tkt. амглще. age. oi co&mlc щ/л. larger because the cross sections Pointi- indicated by * OAZ loom ШгЬЬел for 104e production axe larger; hence tt at., 1977; o GaJicJba. Uunoz vt at., a given measured abundance fraction 1975b; S3 1ндт Hagen vt at., 1976. for 10Be will lead to a smaller value ш of f, the surviving fraction. Ла • result the correction (or the total modu lation will cause the experimental aetata In Figure 1 to move to the rljtK •» tha amount of lnterplemetary energy loaa (tn NaV/n const (Е+Е0)*-ь It la now fairly well established thst the amount of Interstellar matter traveraed by cosmic rays at high energlea la lea» than at low energies and that above *4 GeV/nuc References Cesar sky, C. J., and Audouze, J., 1974, Astron. & Astrophys., 30_, 119. Fontes, P., 1975, Ph.D. Thesis, Universlte de Paria-Sud Centre D'Orsay, France. Garcia Munoz, M., Mason, G. H., and Simpson, J. A., 1975a, Ap. J., 201, L145. Garcia Munoz, M., Mason, G. M., and Simpson, J. A., 1975b, Ap. J., 201, L141. Gleeson, L. J., and Axford, W. I., 1968, Ap. J... 154, 111. Hagen, F. A., Fischer, A. J., and Ormes, J. F., 1976, GSFC Preprint X-661-76-105. Lezniak, J. A., 1975, Nuc. Inst. & Methods, 126, 129. O'Dell, F. W., Shapiro, M. M. , Silberberg, R., and Tsao, C. H. , 1975, 14th International Conf. on Cosmic Rays, Munich 2_, 526. Shapiro, M. M., 1972, Rapporteur Paper, 12th International Conference on Cosmic Rays, Hob art, (J, 422. Shapiro, M. M., and Silberberg, R., 1975, Phil. Transactions Royal Society (London), A277, 317. Shapiro, M. M. , Silberberg, R., and Tsao, C. H., 1976, private correspondence with J. A. Lezniak. Webber, W. R. , Lezniak, J. A., Kish, J. C., and Simpson, G. A., 1977, Astro- physical Letters (in press). Westergaard, N., Ramaty, R. R., and Rasmussen, I. L., 1975, unpublished manu script, Danish Research Inst. 5-S7 Isotopic Composition of Low-liner 9 >• Cosmic-Ray Nuclei with 2-10-H by G. A. Simpson, J. Barbary, J. Kish, J. A. Lezniak and W. R. Webber Physics Dept. University of New Hampshire Durham, N. H. 03B24 1 lieorctical Q lixpcrinienlil fiTI Both D We report measurements of the isotopic composition of cos mic-ray nuclei with Z-10-14 made in 1974 with an improved Ceren- kov x total-energy telescope. These observations cover the energy range from'approximately 450 to 600 MeV/nuo. Mass reso lutions o -v.0.4 AMU are achieved, which allows us to make state ments about the abundance of individual isotopes of these nuclei as well as mean mass estimates for individual charges. The de tails of the experimental technique and its advantages and limi tations are discussed. Within the selected energy intervals, the following mass distributions are observed at the instrument: 20+21Ne=83%, 22НелЛ7%, 2*+25Mg=85%, 26МдлЛ5%, 2,+2»Si*95%, 3eSi- 5%. The mean of the composite mass distributions of these charges are determined to be Ne=20.38±0.18, Mg=24.20±0.15, and Si=28.20±0.18. An interpretation of these abundances in terms of the source composition and interstellar fragmentation will be made. Coordinates: OG Origin 1.6 Mailing address: Dr. G. A. Simpson Physics Dept. University of New Hampshire DeMeritt Hall Durham, N. H. 03824 54« A MKASUUMCNT OF TO ISOTOTIC COMPOS ГПОХ 0Г COSMIC-HAT Ft AND ОТШЖ ШКХ11 Will S-20-2S C. A. Simpson, J. Kiah, J. AN Lesnlak and U. 1. Webber Physics Department University of Hew Hampshire Durham, N. R. 03824 We report Kara measurements of tha Isotoplc composition of coemlc-ray Fa nuclat and othar nuclei with Zlt20. The energy range of analysis la reatrietad to • relatively narrow Inter val of *130 NeV/nuc for Fa, thus permitting a meaa reaolutloa - 0.40*0.06 AMD. A total of *80 fa nuclei are analysed for their laotoplc composition, approximately 67X of which are s*Fe; algnlfleant fractlona of 5*Fe and s^Fe ere alao observed. The aource abundancea of Fa laotopea ara found to be 5*Fe » 65±10X, seFe - 28±8Z and "Fe - 11±9X. Kuclal with 2 - 20-24 generally have an laotoplc abundance consistent vlth е frag mentation origin. In £his paper we present new results on the Isotoplc composition of Fe group nuclei using an Improved version of our original Carenkov x total-energy telescope (Webber tt at., 1973). From this new dsta we Infer a mass resolu tion o£0.4 AMD for Fe over a limited range of energies—a factor >2 better than earlier measurements. The Experiment. The instrument used in this study Is Illustrated in Figure 1. It.is a multipurpose experiment, designed to study the charge composition and energy spectra of coemle-ray nuclei as well as the isotoplc composition In the dE/dx x Е and C * Е modes. It was flown three times In the summer of 1974. The total collection factor was 8450 n2 ater-sec, and the average atmospheric depth was 2.8 gm/cm2. Isotoplc analysis la achieved in this experiment by combining measure ments of: 1) the particle charge fro» the dE/dx x Е and dE/dx x C counter combi nations 2) the particle velocity from the UVT luclte Cerenkov counter C, and 3) the kinetic energy as measured In the thick scintillators El and E2, which bring the particle to rest. The operation of the C x Е method for isotope resolution has been dis cussed in our previous publications (Webber it at., 1973a, b). Basically one can show that for particles which are just above the Cerenkov threshold, but which stop in the total-energy counters, the fractional masa aaparatlon is greatly enhanced over the simple mass-difference fraction UJL This arises because, juat above the Cerenkov threshold, the magnitude of the Cerenkov light emission is a rapidly varying funcf jn of velocity; hence, if the Cer enkov output is measured to a certain precision, the velocity will be meas ured to a much greater precision. Two significant technical innovations present In this experiment are the use of diffuse white-paint precompenaation of detectors to remove spatial non- unlformltiea associated with the light collection process, and a simple path- length error correction scheme based on spherically curved counters and a 54» е Ingle MMurtMoC of the radius of the particle trajectory fro* the detector axis (Stsmson, If77). With the path-length correction, the Ha lting resolution of the Oareakov de tector la «45 for Fa part Idea. Although all partlclea that atop In the total energy countera Й?^' and are above the Cerenkov threah- It old can be ueed In the C " I aode кол ye»» of analyala, the actual amaa reso lution la a atrong taction of en KXIM ergy, which la maximised over a » V**4" •ИС Jf *L-t fairly narrow band of energlea. Wa 1КХЯУИГ» llluatrate this effect for Fa nu .EI'Z4*XM*,HL-V clei In our teleacope In Figure 2. IPO* »*ИЛ) Thla figure above the maa» separa- tlon In percent/AMD aa a function I0NC fPM> of energy In both the C dimenaion and the total Е dlmanalon, aa pre км; wee . dicted by a detailed calculation (TMCE t МЧ1 of each detector's reaponae. Alao shown are the actual reeolutione of the Cerenlcov and total-energy counters. Щши. 1. Schematic, йяалшф о{ 1Ч14 iKAtion o{ coimic-Jiai/ iiotopt txpvumuit. This figure shows that: 1) The resolution In the El dimension IRON RESOLUTION-SEMMXMN is fat better than ia neceaaary to i i i i ni| i i i i i ni| i i i i i HI resolve adjacent Fe and other heav ier nuclei isotopes over the whole range of energies. 2) In the C di mension the 1 AMU mass aeparation for Fe Is greater than the 2o* reso lution over a range from 540 to 640 MeV/nuc. This energy Interval cor responds to particles stopping in the El counter—by the time the I I 1 8 I> particles have reached E2, the maaa Ui i i • 11111) n 'i itinij J i fi'ni > separation has deteriorated. For this reason we have utilized only El stopping data px thia study. This "tuning" for maximum mass res Iff El olution in the C * Е node by «elec ting the energy range of-analysis has been recognized In the theoret 4r -1ft ical discussions (e.g., Webber e£ (С/Сщ^Х at., 1973b) but has not previoualy Figwu. 2. Fe nuclei man Auotution end been used in practice. man iepanation In tht Ce/ienfeov and Abpects of the Data Analysis and totaZ-гтлад aounWit. Calibration. The rejection of nu clear interactions in the telescope is of great Importance In thia technique since the particles are brought to reat and typically *l/2 of them Interact. 550 In addition the light output In the plaatlc aclnillletora and la the Qsrvtxkov counter auat be known accurately as a function of Z and • ao that the location of aaea llnaa «ay ba put on any two-dimensional matrix of events. Thaaa procaduraa ага dlacuaaed In a aeparata paper on experimental tech niques (Webber et *£., thla conference). The Experimental Pata. The first atap In the aaaa analyala la to claaalfy events according to chert,* using a two-dimensional charge analyala technique (dK/dx vs. Е and dE/dx vs. C). Next a matrix of the Cerenkov versus total en ergy output la made for each charge for particles stopping In El. This matrix la shown for Fa nuclei In Figure 3. Note that the aaaa resolution appears to be good enough to define a slngla mass line between tha two endpolnte. A pKioKL, the actual mass corres ponding to thla mase line la unknown and must be established from tha Instrumental calibra tion. The llnaa drawn In Figure 3 ahow the response lines for 5HFe, 56Ie, and seFe pre dicted from the El response funct.lon which le consistent with the endpolnt data for all charges. The predicted mass line for 56Fe is a good lit to the principal mass line observ able In Figure 3. We may thus infer that the data in this figure represents dominance by a single Fe Isotope. Although the shape of a mass line in C- El space is accurately defined, there is an error on the absolute value of this mass line. Based on the uncertainties in determining the endpoints for various Z particles in the dif ferent scintillators and the calibration Tigwie. 3. Fe nuctzi using Scandium as noted below, we estimate Cvienkov x total гпелду the absolute mass error a-±0.3 AW. Thus the mxt/Ux. data in Figure 3 corresponds to a dominant mass of 56.0±0.3 AMD. We now construct series of mass histograms by summing events in C-El space perpendicular to the Inferred mass lines for each charge. This data is presented In Figure 4. Note, in particular, the mass distribution for Sc nu clei. This charge has only one stable isotope lfSSc, and the mean mase that we measure is 45;17±0.21 AMU and is consistent with a single isotope. This forms a very important confirmation of the calibration of the mass scale for all heavier nuclei. Next consider the mass histogram for Fa. Assuming that the observed Fe distribution la dominated by a alngle Isotope 56Fe, and tha 56Fa distribution can be represented by a Gaussian function, the inferred mass resolution for Fe la o»0.40*0.06 AMD. This resolution ia a factor $2 x batter than the eati- mated mass reaolutlon crv0.9 AMU for Fe In our earlier experiment (Webber Vt at., 1973) or the value of otl.l AKU inferred from tile data of Fischer et at., 1975. With this resolution it should be possible to resolve adjacent isotopes of equal abundance; therefore, for moat heavier nuclei not dominated by a utp- gle isotope, the isotopic composition derived by summing events in mass -inter vals of A-±0.5 AMD should be meaningful. Ml Havre 5 екове a *-A MtrU for «11 zbo amelel—la effect a matrix o{ the awelldaa. la tale flgare, кожи ceatered oa inte gral аам awaken wltk • width of ±0.5 AMD are afcowa. The leotoplc composition of all aaclel wltk Ж-20-2С la ohtalaed fro» etmple cowntlag of events la each Vox, with the exception of Fa where the number of events la found by fit ting Geuaaian distributions to the aaaa hlatograv ahowa In Figure 4. Thla data la ahowa In colum 4 of Table I. *—*—4 figuxt 4. Uau hi&togtiam {ол 1*10-16 MUJUU.. Table I Z-20-26 Isotoplc Coapoaltlon Corrected For Predicted Energy Inst. Corrected Aa- Iaotope Spe- Mean Inter- Ob- Frag- For Atao- Fraction «timed Fraction cles Energy val served men- sphere & Z of of Source of Z, A (MGV) (MeV) Events tatIon Spectrum Element Fe-56 Abund. Fe-56 26-58 582 127 14 54.4 65.1 20.3±5.7 27.8±7.6 29 28 26-56 597 133 53 198 247 73.0±10.2 100 100 100 26-54 613 139 8 17.1 20.1 5.914.0 8.115.4 2.3 24-52 573 112 4 13.6 13.0 51126 5.912.9 11.0 24-50 589 116 4 11.1 12.3 49126 5.613.0 _ 9.6 22-48 539 93 6 19.5 21.4 74130 12.214.9 8.1 22-47 547 95 2 3.8 4.1 14115 2.312.5 - 6.7 22-46 555 97 2 3.3 3.5 12114 2.012.3 6.6 21-45 531 89 7 18.4 18.6 100 11.215.1 4.5 21-43 548 93 2 _ 20-44 511 86 2 5.5 5.6 13.919.7 3.512.4 5.3 20-43 520 89 2 5.6 5.8 14.119.9 3.612.5 - 4.5 20-42 530 92 8 22.0 23.9 56.Ш9.6 14.114.9 14.0* 20-41 539 95 3 6.4 7.1 15.91113 4.012.6 1.3 *Croe« action 26-56 to 20-42 - 27 a». «tlon*. To obtain the partial frajjacmtatlo i encountered In the Instrument and the atmosphere, target factor» presented by Llndstroa et at., 1975 to multiply the (hydreasa) cross sections. The proton-nucleus croaa sections ea emperlaeatal maaaareaenta where poaslble (e.g.. Perron, 1975) and irlcal fosaala Of Sllberherg and Taao, 1973. For total-reaction мм ths foraala of Meyer and Casse, 1975. ЗЪ2 The Instrumental correction la «ad* by eatlmatlng the flux of »*r- •irt+ririrtrlrint- tlclee which, tmcldent on the to», of tha instrument, produce the observed number of events corrected for neu гУМ— Nm иатр|Лч tron «tripping reactions. Ma do *o-r Isotope £>:>о thla by multiplying each abundance value by the factor exp /X, where г * го X, tha Interaction mean free path, 1« given by X - 1.67 * КГ* A/e-r Ч where o? la the total interaction cross section, and A- la the average atomic uelght of the material tra т- versed. To correct for atmospheric In teractions we use a propagation code, which Includes the complete network of atmospheric nuclear frag mentation reactions In a one-dimen sional model. A mean atmospheric depth of 3.0 g/cm corresponding to the telescope's average zenith angle of 18* has been used in making this correction. For the Isotopes in question, the magnitude of the cor rection ranges from +33Z for 56Fe to -8X for 52Cr, which Is built up in Пдиле. 5. 2-А pZot $o* 2*20-26 nuateJ.. the atmosphere by fragmentation from Fe. These numbers refer to particles in slightly different energy intervals and therefore are not directly comparable. We use the measured spectrum of Fe presented by Meyer tt at., 1974 to normalize our results to the same energy (600 MeV/nuc) for all charges and isotopes. It should also be noted that even though the overall instrumental and atmospheric corrections are large—for example, for 56Fe—the total correction Is 4.43, the difference between the total corrections for 56Fe and 58Fe Is only 32. Therefore our lsotopic comparisons are essentially unaffected by these corrections. Isotoplc Abundances Above the Atmosphere. Of special note in the iaotoplc ra tios are the following: 1) Fe has a substantial abundance of 58Fe (20±6Z). This component can not be produced by propagation effects and is the first clear signal that the source regions that produce cosmic rays are different from the regions that produce material typical of our solar system. The abundance of 5h?e at the top of the atmosphere is 6 t 4%, only mar ginally above the criterion of significance. The mean mass of Fe is S6.31±0.24 AMU. 2) Cr is observed with roughly equal abundances of the principal iso tope* 5QCr and 52cr. The mean mass of this element Is 50.9±0.55 AMU. 3) Ti has a dominant Isotope 1,eTi, which comprises 71+2831 of all Ti. The mean mass of this element is 47.5±0.4 AMU. 4) Sc was used as a response function calibration point assuming it to be 100% t,5Sc. The mean mass of events with masses between 44 and 46 AMU is 45.17 AMU. 353 5) Са la largely **Се. Other Isotope* coaatltate half or leae •( Cal- clwa. *"Ce, «kick la ether «ark kaa b»aa Miasms' ta be. preeeac tn the eoexce and therefore la the eomtaaat laotoae, la aat araaaat. The aaaa aaaa of Ca ta 4I.lltO.J7 AND. Extrapolation ta tha Cosmic-Bar «oarce* —4 a Coamarlooa with Itaclaaaratbaala Theorloe the dlffualoa model «kick kaa beee moat widely ueed la ex»lata lag arav: oua coamtc-ray t-imm aettloa maeaarsmaate hae beam the "leaky-hex" aoaal (Cavab IMS). Thle aoaal learns ca am axeoaemttal dletrtautloa of Ufotlmea or mate- rial path lamgthe la tha lataratallar medium. The pravloee charge coamoeltlon data ha«a aacabllahad Char tha aaan path length la >-• g/cm* of Utaratallar hydrogen for energies <1 OeV/nac. The lncaratallar calculation la performed by setting up aa laltlal eourcc abundance dlatrlbotloo and integrating atepvlee tha coamlc-ray transport aqua tion. At each atep the fragsmntatloa, radlo-aucllda decay and energy loms of the partlclaa are calculated. Ш have assumed a aat of source absmdancse (аае Tabla I) and calculated a «at of laotoplc abundancea at tha boundary of tha helioephere which provide a cloae agreeaant with tha meaaured laotoplc frac- tlona. He alao have modified tha fragmentation parametera for Ca Isotopes aa noted below. The aourca abundancaa require tha preaance of 21±(Z of 5*Fe at the aource, and an upper limit of tl4Z for the abundance of SHFe. Except for poaalbly Ca, the obeerved laotoplc abundance of tha other elements la coaala- tent with their origin aa Interatallar fragmentation prodocta of Pa. In the caae of Ca tha charge abundance predicted ualng tha aamlemplrlcal croae-sectlon formula is only about 0.5 of that actually measured. In our ex periment only 18 Ca nuclei are analysed for their laotoplc abundance, eo that it la poaalble that one la observing an extreme atatietlcal 'variation. How- етег, our overall charge abundance ratio ^/Fe » 0.28 la very alxdlar to that obtained by other worker» baaed on much greater atatlatlca. In tha peat tha difference between tha predicted and obaarved amount* of Ca relative to Fe waa ueually attributed to a aourca abundance of ч0Са » 15* that of Fa (e.g., Sha piro and Silberberg, 1975); however, that poaalblllty la not allowed by our data ainca we *ee no <*°Ca. we obaarve that, ualng tha semlemplrical formula of Silberberg and Taao (1973), tha total croaa section for Fa fragmenting to Ca la t<48 mb. There are no direct measurement* to confirm thia value and It is poaalble that it la low. We have carried out a aaparata calculation using a total crosd section of 70 mb (**2Ca • 27 mb). The** predictions are in much closer agreement with our observations. References Cowsik.R., Pal.Yash, Tendon,S.N., 6 Verma.R.P., 1965, Phye.Rwv., 158. 1238. Fischer,A.J., Hagen,F.A., Maehl,R.C., Oraea.J.R., a Arena,J.F., 1976, Astro- phys. J., 205, 938. Lindatroa, P. J., Grelnar, D. E., Heckman, H. H., Cork, B., and Buaer, F. S., 1975, Preprint, CBL-3650. Meyer,J.P., Caasa.M., & Waatargaard.N., 1975, Proc.l3th Int.Cosmic lay Coof., 12, 4144. Meyer, P., Ramaty, R., and Webber, W. R., 1974, Phyalca Today, 27_, 10. Perron, Claude, 1975, Ph.D. Thesis, Centra-d'Oraay Univareite Parla,-Sud. Shapiro,M.M., & Silberberg^R. ,1975,Proc. 13th Int.Cosmic Ray Coof.J4unlch,2, 538 Silberberg, R., and Taao, C. H., 1973, Ap. J. Suppl., 25,, 315. Simpson, G. A., 1977, Ph.D. Thesis, University of Haw Hampshire. Webber, W. R., Lezniak, J. A., and Kish, J., 1973a, Ap. J., 183. L81. Webber, W. R., Lesniak, J. A., and Kish, J., 1973b, Hue. Inst. a Math,, 111. 301. iu А ОЕЯСЯЖСГ СОвСС NOT lBJTUni, Mx P. S. Freier, C. M. Gilman, И. R. Scarlett and C. J. Waddington School of Physics and Astronomy, University of Minnesota Minnaapolia Kinnaaota 55455, U.S.A. Theoretic*) 13 Empsrimwt»! [xj •<*» Q An array of particla datactors has baan developed to study the isotopic composition of cosmic ray nuclei having charges between those of neon and nickel. Design details of this array were discussed at the Munich Conference, Vol. 9, p. 3166. Since then the array has been calibrated with argon ions accelerated at the Bevalac and has been exposed to the cosmic rays on a balloon flight. The calibration data have been used to study the uniformity of response and charge resolution obtainable from the electronic detector elements of the array. The balloon flight vai not entirely successful due to a pressure leak, and only about 1,000 2 > 10 nuclei were completely detected. As a consequence the nuclear emulsions have not yet been developed and hence mass determinations have not yet been made. We will report on the charge and energy spectra obtained from these data. A second balloon flight is scheduled for Spring, 1977. CoonJirutes: OG 1.6 Isotopic Composition of Cosmic Rays Mailing address: C. J. Waddington Professor, Department of Physics/Astronomy University of Minnesota .375 Physics Building 116 Church Street Minneapolis, Minnesota 55455 555 ISOTONIC CQNPoimoN or IMM titour MUCUI AT J20-S00 HtV/w S.P. Milan. I.S. Cartwrlght. f.D. Price end C. Tarl* University,of California, Berkeley Theoretical [3 ВЖЦНИМШ (3 Both Q W« report oo preliminary analysis of data raturnad by the I^IS-II Iron Isotope, experiment, flown On a high altitude balloon fro*) Watertown, South Dakota In September 1976. The Instrument consists of eight spark chambers, two scintillators, and a pre cision Cherenkov counter to measure particle trajectory, charge and speed, and a passive polycarbonate plastic stack to measure range. In the range 320-500 MeV/emu, Isotopes are resolved using the Cherenkov vs. range mode of analysis. Charge Is measured by both Cherenkov vs. dE/dX and dE/dX vs. range, with redundant charge information provided by cone etch rates in the passive plastic detectors. Hon t.'ian 1600 nuclei of the Iron group were collected In 2k hours under 3 g/cm* residual atmosphere. Our present esti mate Is that the mass resolution will prove to be a ..35 amu and that much of the analysis will be completed in time for presentation at the meeting. jfbffingaddn»: Prof p Buft)r- pr(ce Department of Physics University of California Berkeley, California 94720 A STUDT ОГ ТЖ ISOTOPES OP COSMIC «AT Г. KOCUtt EM LEXAM POLYCAUOKATI EXPOSES OR STYLAl P. B..Price and t. K. Shirk University of California •ark*lay C. O'Cea Heigh, D. 0*Sullivan *nd А. Thompson Dublin Inatituta for Advanced Studies Ireland Theoretical • Expariwaatal Q) Both Q One of the Lexan Polycarbonate atacka expoaed during the Skylab Ultraheavy Coamic Kay Experiment waa uaed to atudy the iaotopic coapoaition of Iron Group nuclei.- The atack, which conaiated of 32 slieeta of Lexan, each 250 ua'thick, waa expoaed for 253 days inside the spacecraft. Results of measurements on approximately 400 nuclei that stopped in the stack, with energiea in the interval fro» = 130 to = 200 MeV/N will be presented. Coordinates: OG 1.6 (Isotopic Composition of Cosmic Rays) Mailing address: Drs. D. O'Sullivan and A. Thompson, Dublin Institute for Advanced Studies, 5 Herrion Square, Dublin 2, Ireland. >57 ТЖ SKCTMM OP COSMC MJCTHOe wlTM ENERGIES t AND 100 C*V Charles A. Mw|u Marshall Space Plight Center. Humtsvllle. М*кш (ША) 5M1I James A. Earl Department of Physics and Astronomy, University of Nmrylamd Collet* Park. Maryland (USA) 20742 This paper reports final results on cosmic electrons observed during a balloon exposure of 3S00 t 60 в1sec sterad at an average depth of 4.1 g/ся2. Within the detector, the development of showers in 33.7 rl of lead was sampled by 10 scintillation counters and 216 (M tubes. The identification of electrons during flights was confirmed by extensive analysis of data from a ground level exposure, and from a calibration at the Cornell Electron Synchrotron. The spectral intensity of primary cosmic ray electrons in particles/s^sec sterad GeV was found to have the following power law dependence .upon the electron energy Е in GeV: 2. The Experiment. The detector was a calorimeter amde up of 10 scintillation counters, в trays of Geiger - Mueller tubes, and IS lead plates. This instru ment was similar in configuration and operation to the one employed by Earl, Neely and Rygg (1972), but its geometric factor of SS2 cm2 star, «as %4S times larger, and its calorimeter thickness of 55 rl. was %1.8 times greater. The experiment was flown successfully on balloons three times: twice in 1969 from Sioux Falls, South Dakota, and once in 1973 from Palestine. Texas. During a total time at ceiling of 48.3 hrs an exposure of 3SO0 * 60 в2 sec. sterad was accumulated at an average depth of 4.8 g/cm2. In addition, an exposure of 13,600 m2 sec. sterad was performed outdoors at College Park, MX)., under flight conditions. The ground-level spectrum of secondary electrons was measured during this run. Because the detailed response of a lead-plastic calorimeter сал not be predicted reliably on the basis of shower theory alone, a calibra tion was performed at the Cornell Electron Synchrotron. 3. Results. Combined data on the primary spectrum between 6 and 100 GeV, which are given in Table 1 along with the ground level data, can be described by the power law ^» (800 ± 60) Е-3-**0"1 ah which represents а least squares fit to the data. In this equation, the spectral intensity dJ/dE is expressed in electrons/m2sec sterad GeV and the energy Е is expressed in GeV. This spectrun is in excellent agreement with the results of Earl, Neely and Rygg (1972). Table 1 Spectral Intensities (electrons m"2 s_1 sr"1 GeV1 Combined Flights Ground Level Energy Intensity Energy Intensity 6.42 (1.3 ± 0.3) 0.56 (4.8 ± 1.0) 11.4 (2.2 ± 0.3) x 10-1 1.00 (1.3 ± 0.3) 20.3 (3.2 ± 0.S) x 10"2 1.78 (2.7 ± 0.3) x 10"1 36.1 (S ± 1) x 10-3 3.16 (6.6 + 0.7) x 10-2 64.1 (6 ± 2) x lO"1» 5.62 (8.9 ± 0.8) x 10"3 114 (3+2) x 10"5 10.0 (1.4 ± 0.2) x 10"3 17.8 (2.2 ± 0.6) x 10"ц 559 4- Discussion. We interpret the observed spectrum, which it tigmiflcaaUy (ttfptrthMtbt nuclear spectrum, *t • fully steepened o*e reflecting •» equilibrium between Injection and energy lo»«. The electron spectrum 1* steepened only above • break mingy which evidently Hit be less them t GeY. This implies that т » 12 Hrr which 1» not compatible with the widely «voted value v • 3 Hyr (Shapiro and Silberberg 1970) obtained fro* the мм path 1 «fifth of ~-i g/cmz travened by cosmic rayt near 3 CeV. To resolve thi» contradiction, we iu||eit that the confinement volume for cosmic rays extend* above and below the galactic disc to include a region of reduced gas demsity where the Compton-synchrotron mechanism is operative but where the probability of nuclear interactions is small. The results of Garcia-Mumot, Mason and Simpson (1977) confirm the interpretation reached here. Prom their observations of the Be10 isotope, these authors conclude that t > 10 Nyr. This agreement between the leakage time measured by a clock based on nuclear radioactivity with that measured by a clock based upon electron energy loss constitutes one more piece of evidence that fits into a halo model of galactic cosmic-ray con finement . This research was supported by the National Aeronautics and Space Administration under grant NGR 21-002-066. References Brecht, J., Univ. of MD, Dept. of Phys. and Astr., Tech. Rep. No. 955 (1969). Earl, J.A., D. Neely, and T. Rygg, J. Geophys. Res., 77, 1087 (1972). Garcia-Munoz, M., G.M. Mason and J.A. Simpson, APS Bulletin, 28, 567 (1977). Meegan, C.A., Ph.D. thesis, Univ. of MD (1973). Meegan, C.A. and Earl, J.A., Conference Papers, 13th International Conference on Cosmic Rays, Vol. 5, p. 3067 (1973). Meegan, C.A. and J.A. Earl, Conference Papers, 14th?International Conference on Cosmic Rays, Munich, Vol. I, p. 419 (1975). Meegan, C.A., and J.A. Earl, Ap.J., 197, 219 (1975). Neely, O.E., Univ. of MD, Dept. of Phys. and Astr., Tech. Rep. No. 845, (1968). Prince, T., D. Mullerand G. Hartmann, Bull. Am. Phys. Soc, 28_, 601 (1977). Shapiro, M., and R. Silberberg, Ann. Rev. Nucl. Sci., 20, 323 (1970). Но COSMIC NAY ELECTIONS: A DISCUSSION OF RECENT OKEKOTIQ» Dietrich Wlltr end Thomas Prince Enrico Fermi Institute, University of Chicago Chicago. Illinois COC37 (USA) Our recent measurement of tht spectrum of court c ray electrons hat provided statistically significant evidence for a spectral shapt that Is wen steeper than that of protons. Tht electron spectrum dots not fit Mil to a single power 1«м, and tht abundance of electrons relative to that of protons decreases from % IX at 10 GeV to % 0.1S at 300 GeV. This result 1s consistent with a galactic escape HfttlM for electrons exceeding 10' years. Mt shall discuss our data In light of current models for the propagation of cosalc rays 1n the galaxy» and conclusions will be drawn concerning the consistency of various models with the observations. 1. Introduction. Significant Information can be obtained from an accurate Measurement of the high energy cosmic ray electron spectrum because of the Interaction of electrons with the magnetic and photon' fields 1n tht galaxy. Electrons lose energy by synchrotron radiation and Inverse Compton collisions (in the Thomson limit) at the rate: d«V«*fc a-k6* where Е , 1 1s the electron energy, k а »0" *CW»h * H (Hl/tx)] (GeV sec)" , H*. 1s the mean perpendicular magnetic field component (1n gauss), and m^h 1s the ambient photon energy density (In eV/cm3). Such energy losses define a radiative lifetime XT» a, («B)"' which can be compared with the time scale of other loss processes, such as leakage from the galaxy, which are expected to affect all cosmic ray species. In this paper, we will Interpret our new measurement of the electron spectrum above 10 GeV (Hartmann et", al. 1977) in the context of models for the storage and propagation of cosmic rays In the galaxy. The most Important feature of our data that must be considered Is a spectral Index 0b* 3.0 above 10 GeV (Figures 1-3). Such a spectral Index 1s considerably steeper than that of the primary nuclear component («.^ 2.7). This feature strongly suggests the influence of radiative energy losses on the shape of the electron spectrum and we shall discuss Its Implications In some detail. 2. Comparison of Electron Measurements. He wish to first compare and contrast our measurement of the high energy electron spectrum with the results of other experimenters. A mil tl decade logarithmic plot tends to mask the differences between the various results, so In Figure 1 we have plotted all measurements of the differential energy spectrum of electrons divided by a reference energy spectrum of E~3*°. Large discrepancies In absolute flux outside the quoted error bars are Immediately apparent, and are Indicative of systematic errors 1n at least some of the experiments. Besides these differences, significant discrepancies also exist In the quoted spectral Indices which range from*» 2.7 to at> 3.4. Figure 1 Indicates how difficult It 1s for most experiments to make a definitive statement about tht spectral Index above 40 GeV due to the size of the experimental errors. We believe that the combination of good statistics, good background rejection, and Work supported, in part, by NSF grant AST74-16310 and NASA grant NGL14-001-005. 541 юоо 1 Ч 1 1 vi 1 I I | 4 I I I va i I ш о Id юо MlU t » ANAND*t«,'73 * FOLK».*T9 o UATSUOrtal.TS Ui 4 MCEMN ft EARL.'TS | HOLLER.'73 o SILVERIER»,'7« O aUFFHMTON «t t.'TS + FREER •••I.TT • THIS, WORK Ю _i_l 10 Ю0 Ю00 ENERGY (6«V) Figure 1. Differential energy spectrum of electrons multiplied by E+3-0. The errors for our data are statistical only. extensive high energy calibration make our measurement very reliable. Comparing our results to those of other experiments (Figure 1), we find ourselves 1n good agreement with the results of Meegan and Earl (1975). Our data »re also qualitatively similar to those of Sllverberg (1976), although the differences are outside the quoted errors. We also note that our measurements below 40 GeV are consistent with the results of Fulks (1973). but disagree with the low energy data of NUller and Meyer (1973).* Although our data are consistent with at least one other experiment» they differ from previous results 1n one significant feature. Other experi ments have found their data to be consistent with a single power law spectrum. Me find that a single power law 1s not a good fit to our data. Further» our data suggest a gradual steepening In spectral Index from** 3.0 at 15 GeV to oc % 3.4 above 40 GeV. 3. Interpretation of Results. We now wish to discuss various models for the propagation of cosmic rays In the galaxy. It should be stressed that any Interpretation of the high energy electron results 1s extremely model dependent. One Is Immediately confronted with the problem of an abundance of free parameters In the models, many of which are only approximately known. The electron data themselves do not specify a unique model. Rather, given a * As noted by Milller and Meyer (1973), their data below >x<30 GeV may be too low due to an erroneous dead-time correction. It appears now that ti «s was the case, and these previous low energy results should be considered withdrawn. J6J «ode), the electron data can be used to put constraints on the parameters of the «ode!. We therefore wish to concentrate on the simplest models with the fewest free parameters. We first discuss the homogenous Model and later the disk-halo models. A. Hwy*nus Model. The homogeneous or "le*k^ box* model Is governed by the equation: 3§* •&<-«'««)-*«> where N(E) Is the density of electrons of energy E, 'C(E) is the «scape life time from the confinement volume, k 1s the energy loss coefficient defined In the Introduction, and Q(E) 1s the source function for electrons. All quanti ties are assumed tc be position Independent. We also assume the escape life- . t1me*C(E) to -be energy dependent In the form of a power law: *U(E)a'C.fffe^ (E0 • 1 GeV). This assumption with fw 0.5 would be In agreement with *•' measurements of the nuclear cosmic ray composition (e.g. Caldwell 1977). We further assume that the injection spectrum of electrons follows a power law: ftCE>- AE"r The solution to equation 1 1s discussed by Ramaty and Silverberg (1974). Up to a normalization factor M« , N(E) is determined by 3 free parameters: Г i This yields а solution for z • 0 and "Vi_ -* °* of (4) i where N, Is an overall normalization constant Qualitatively, this model leads to the following Interpretation. An electron of energy Е has a lifetime "CR • (W6) • against radiative losses. During this time It propagates an average distance f which Is proportional to the square roots of diffusion coefficient and lifetime: J • (гоГц)*'»' • (I-P/KE)* . For large energies, ? 1s smaller than the thickness of the disk, I.e. J*u , and a fully steepened spectrum will be observed: N(e> <*» Е'«**'> . For sufficiently small energies, ? exceeds the dimension of the source disk and thus defines the dimensions of a halo whose volume increases with the Inverse square root of the electron energy. As a consequence, the spectrum observed within the disk will be proportional to e" Gemot Hartmann*. Dietrich Mffller, end Thomas Prince Enrico Fermi Institute, University of Chicago Chicago, Illinois 60637 (USA) We have measured the spectrum of cosmic ray electrons with a new instrument that combines a transition radiation detector with a shower detector. The transition radiation detector provider, unique Identification of Individual electrons and good discrimination against protons. At the same time, 1t allows the construction of a large area Instrument (0.48 m2 ster) and consequently makes possible a measurement of Improved statistical accuracy. The Instrument has been calibrated with electron beams of 5-300 GeV at Fermilab, thus eliminating energy dependent biases. A first balloon flight yielded 30 hours of data at an altitude of 5 g/cm2. We shall describe the design of the Instrument, the accelerator, calibrations, and the analysis of the balloon flight data. The spectrum of electrons is found to be significantly steeper than that of protons over the whole energy range. IN Introduction. The shape of the spectrum of high energy cosmic ray electrons has remained controversial -for many years. The important question 1s the existence of a possible steepening of the spectrum which may occur as a result of radiative energy losses of electrons in the Interstellar magnetic and photon fields. Such a steepening would provide significant Information about the propagation and lifetime of cosmic rays in the galaxy. Most of the electron data until the early 1970's (including, for Instance, the measurements of Anand et al. (1973) and the previous results of Miiller and Meyer (1973)) supported an electron spectrum that has the same slope as the proton spectrum (differential spectral index a % 2.7). Some of the more recent results are consistent with steeper spectra: Silverberg (1976) and Matsuo et a]. (1975) have reported a % 3.1-3.2, and Meegan and Earl (1975) found a spectral index a x 3.4. Significant discrepancies exist not only with respect to the slope of the spectra but also with respect to the absolute electron fluxes. This situation is undoubtedly caused by experimental difficulties, most importantly the following: (1) In the shower detectors used In previous measurements, interacting protons may masquerade as electron showers, leading to a substantial back ground that is difficult to correct for. With shower detectors, individual electrons cannot be identified, unless the interaction is made visible in a nuclear emulsion. (2) Until recently, detectors could be calibrated at accelerators only < up to electron energies of 20 GeV. The energy response and shower character- '. istics had to be empirically extrapolated to higher energies. (3) The small flux and limited detector size made large statistical errors unavoidable. * Present address: Max Planck Institute, Garching, Germany 367 To overcome thtst ГМ1ШП difficulties, м have constructed а мм dttector of vtry large im (geometric factor 0.48*7 ster). that 1s able, with tht aid of transition radiation detectors, to Identify Individual electrons, and that also has been calibrated at Fermilab with electron-beams covering the energy range 5 to 300 fieV. This Instrument has been exposed In a first balloon flight 1n October 1975, for 30 hours at 5 g/cm2 of residual atmosphere above Palestine, Texas. In the following, we shall present pre liminary results from this experiment. 2. Description of the Instrument. A cross" E523 SCMTILLATOIt N€110 .— POLYETHYLENE ЛММ section of the Instrument •mm LEAD вюа SOL» POLYETHYLENE Is shown In Figure 1. Figure 1 The main components are: Schematic cross section of the Instrument (1) A scintillator telescope consisting of plastic scintillators Tl, T2, T3, and T4. (2) A transition radiation detector of 6 plastic foam radiators, each followed by a multlwlre proportional chamber. (3) A shallow shower detector that 1s formed by lead plates above scintillators T2, T3, and T4. All counter elements are pulse height analyzed. A time of flight mea surement between Tl and T4 Identifies those particles that traverse the detector 1n the downward direction. The wire directions of consecutive proportional chambers are orthogonal to each other, and groups of 5 adjacent wires are connected to a common amplifier. All wire groups are Individually pulse height analyzed, and therefore trajectory and pulse height Information 1s measured (with a spatial resolution of 5 cm) for each particle traversing the Instrument. The shower counter employs a total of 8 radiation lengths of lead (the shower Is sampled after 4, 6 and 8 r.l.). The photomultipller tubes were Individually calibrated with a nanosecond light source In order to Insure linear response beyond the largest expected shower signals. In order to accept a particle as an electron, the following criteria 36fi have to be met: (1) A particle. Identi fied as singly charged by the pulse height In Л, must traverse the detector In the downward direction. The tra jectory of the particle eust be uniquely defined from the signals of the multlwlre chambers. (2) The pulse heights In the shower counter (T2, T3, T4) must be consistent with the profile of an electro magnetic cascade. The shower signals also measure the electron energy. (3) X-ray transition radiation must be detected X 4 • in the transition radiation Я0ТИ И LUO (MOUnON LCWTM) detector; Only particles Figure 2 with Lorentz factors Comparison of measured and expected Y • E/mc2 Jfc 10* will produce shower signals transition radiation 1n saturation. Therefore, the transition radiation detector acts as an efficient discriminator between protons and electrons, and can positively Identify those Interacting protons that lead to shower signals that might otherwise be indistinguishable from electron showers. The major new element of this instrument 1s the transition radiation detector. Its design Is based on results of extensive accelerator studies of transition radiation (Cherry et al. 1974b and Prince et al. 1975). We have learned 1n these studies thafpolyethylene foam (Dow TEnTfoam") can be used as a particularly simple but efficient radiator material, and that the com bination chosen In this experiment, 15 cm thick radiators with 2 cm thick xenon-filled proportional chambers with rather thin windows, yields an optimum x-ray signal under the geometric constraints of a balloon borne. Instrument. We shall discuss at this conference some further details of such detectors 1n the paper by Cherry et al. (1977).. 3. Accelerator Calibrations'. Realizing the crucial Importance of accelerator calibrations of any new cosmic ray detector, we have exposed our Instrument at Fermi lab to beams of electrons, pi oris, and protons covering the entire energy range frcm 5 to 300 GaV. These calibrations yield the following Information: * (1) Shower signals: The average shower signals observed 1n the shower counters T2, T3, and T4 were measured as a function of the electron energy. It 1s not possible to analytically calculate these signals (I.e. the number of secondary electrons) with the required accuracy. However, we were pleased to observe that our measured shower signals are well represented by the results of an earlier extrapolation (Hffller 1972) that was based on measure ments below 15 GeV. This Is «hewn 1n Figure 2. The "shower sum'", I.e. the 569 sum of the pulse heights froa T2, T3, and T4, Is * aonotonlc function of the tltctron energy» tht energy FLIGHT EVENTS resolution being constant EtrOGcV tt «bout 30* FWHN over tht 300- whole tntrgy range. «л X'$4 (2) Transition radiation signals: An iltctron, accompanied by 200- transition radiation x-rays, yields pulse heights 1n the proportion»! chambers that are larger than the Ion 100 ization signal alone by a factor of about 2. The transition radiation signal reaches saturation for electron energies above 5 ICf» 1СГ5 10° GeV (Y % 10*). and renins Ю» constant up to the highest LIKELIHOOD RATIO electron energies covered Figure 3 (300 SeV, Y *6 x 10»). Details of this behavior Likelihood ratio histogram of a sample of are discussed by Cherry night data. The likelihood ratio Is a et al. (1977). measure of the transition radiation signal. The events In this histogram are due to singly charged particles with unique (3) Backscatter: A trajectories which have a total shower signal certain fraction of electron greater than 10 GeV electrons and fairly showers leads to ambiguities 2 1n the determination of the good shower fits (X < 4). electron trajectory due to additional tracks that are most likely generated by particles scattered backwards from the shower detector. Knowledge of this effect 1s very Important, since 1t could lead to an energy-dependent effi ciency correction 1n the data analysis. He measured that the fraction of electrons with unique trajectories decreases slowly with Increasing energy, from 55* at 10 GeV to 21* at 200 GeV. 4. Data Analysis. The analysis of the flight data has been per- formed along the following steps: (1) All events are rejected that do not exhibit a unique trajectory or that do not traverse the Instrument in the downward direction. (2) All pulse height readings are normalized, using In flight calibrations due to penetrating protons,«.-particles, and heavier nuclei. (3) The response of the shower counter Is analyzed. Each set of pulse heights t< (1 * 1...3) measured In T2, T3, and T4, Is compared with expected pulse heights nj(E) which are known, together with the expected variations6j(E), from the accelerator calibrations. Minimizing the parameter 2 2 X •^x£[(t{-ni(E))^/6'1(E) ]determines the energy E, and the minimum value of xz measures the goodness of fit. Obviously, only events with sufficiently small values of x2 can be due to electrons. 370 (4) Likelihood techniques (Cherry et a!.. 1974a) «re used ТЯАМГПОМ RAOUTON «JKCLMOOO ЛОТО to evaluate the transition radi ation signals that are measured . n:: i r;:: 11 j; j j 11:;: l»: t»t!: In the six proportional chambers. The probability that a particle without transition radiation produces a pulse height x due to Ionization in the 1-th chamber Is given by a Landau- : III!!.:;"T::I::;'«: Vavllov distribution Р^Чх). 1 ':K:.:!:\IS.S: -Я! The signal of an electron, ac companied by transition radia 1 ^.:i::i:vi.",i,:' tion x-rays, follows a diffei- ent distribution P w(X.). I e These distributions are known from the accelerator calibra ••• • I. II i| U I t , I tions. For each event, we m ::s-:.! :•. :i :!:.'.•:!•..•: • • I S-o 2» measure six pulse heights ! ЗДЧ Xi(1»l...6) along the tra jectory of the particle. Me •'1 then compute the likelihood ratio L-lPeOHxtyJTPpllJtx,) as a measure of the transition radiation signal, and, there fore, as a means to Identify J each Individual particle. L»l indicates equal likelihood for a proton and an electron, Figure 4 while L»l 1s expected for an Correlation matrix of transition radiation electron (with transition ra signal (likelihood ratio) versus deviation diation), and L«l for a proton from expected shower signal. Electrons (without transition radiation). should have a large likelihood ratio and a small deviation from the expected shower In Figure 3 we show some signal. results of this analysis pro cedure. We have plotted a likelihood distribution of those events that exhibit fairly good shower fits (x2 - 4). The clear separation between protons and electrons Is well demonstrated. He also notice that the number of proton in duced showers Is larger than the number of electron showers. Without the transition radiation detector, a comparably clean identification of Individual electrons 1s not possible. The good separation of protons and electrons 1s also illustrated in the scatter plot of Fig. 4. Here we correlate the linear deviation in the signals of one of the shower counters (at a depth of 6 r.l.) from the expected signal of "ideal" showers, with the transition radiation parameter L. Clearly, a subset of the events is well separated and distin guishes itself as electrons c"ue to large L-values and sma'l deviations from the expected shower signals. iL » Results. We have determined the differential energy spectrum of the electrons identified in lis fashion, after taking all selection effi ciencies properly into account. The resulting electron fluxes are corrected for the amount of residual atm- there, and are*plotted in Figure 5. 'Results from previous Investigation? .... c shown for comparison. We notice Immediately that our spectrum 1s rater steep, significantly steeper than the spectrum of protons. Our result is in good agreement with the data of Heegan and Earl 571 (1975). and 1n qualitative a- I I I III! I I I llll I I I I Ml greement with the spectrum of Sllverberg (1976). It Is testing to Interpret the slope of our spectrum In the 10' context of a steepening due to radiative energy losses. These aspects will be further discussed 1ir the accompanying paper (№11 er and Prince, 1S77). KTr 6. Acknowledge ments: We acknowledge the o help of Messrs. D. Bonasera, a> M. Cherry, E. Drag, M. John son, and L. Littleton, and 10"' of Mrs. N. Beck and Mrs. L. Glemie. Vie also appreciate the assistance of the staff O of Fermi lab and the services of the NCAR Balloon Facility. 4 « Anond et ol,l973 This work was supported in I0 » Fulks, 1975 part by NSF Grant No. AST 74- 16310, and NASA Grant NGL 14- ° Matsuo et ol, 1975 001-005. *Meegon а Earl, 1975 BMuller.1973 References 5 ° Silverberg.1976 ftiand, K. C., Daniel R. R. Ю « Butfington et o\, 1975 and Stephens, S. A. 1973, 13th Int'l Conf. • This work on Cosmic Rays, Denver, 1, 355. Buffington, A., Orth, CD. I06 and Smoot, G. F. 1975, 10 100 1000 Ap. J., 199, 669. Cherry, M. L.,"Muner, D. ENERGY, GeV and Prince, T. A. Figure 5: Differential energy spectrum 1974a, Nucl. Instr. of cosmic ray electrons. and Meth. 115, 141. CherryT'M. L.\ Hartmann, G., Miiller, D. and Prince, T. A. 1974b, Phys. Rev, DIP, 3594. 1977, Plovdiv Conference. Freier, P., Gilman, C. and Haddington, C. «J. 1977, Ap. J., 213., 588. Fulks, G. J. 1975, J. Geophys. Res., 80, 1701. Matsuo, M., Nishimura, J., Kobayashi,T., Niu, K., Aizu, E., Hlralwa, H. and Taira, T. 1975, 14th Int'l Conf. on Cosmic Rays, Munich, 12., 4132. Meegan, C. A. and Earl, J. A. 1975, Ap. J., 197, 219. Miiller, D. 1972, Phys. Rev. 05, 2677. Miiller, D. 1973, 13th Int'l Conf. on Cosmic Rays, Denver, 1_, 361. Miiller, D. and Meyer, P. 1973, Ap. J., 186, 841. MUller, D., and Prince, T. A. 1977, Plovdiv Conference. Prince, T. A., Muller, 0., Hartmann, G. and Cherry, M. L. 1975, Nucl. Instr. and Meth., 123, 231. Silverberg, R. F. 1976, J. Geophys. Res., 81_, 3944. 372 BIGB ЕИЖВСТ MIMAJtT 1ХЕСЛ0И SPECTRUM. OBSERVED BY TBI EMULSIOH СПАЖИ K.Al«u*.l.Hlralw»».ll.FuJll»».J.HUhl»ur>»»,T.T«lr»»»*,T.Kob»y—hl»*«*tK.lllu***— J.J.Lord#,t.J.Wllkee# and *.L.Golded«. * Tanagava Praf.Colledga.Kanagawa ** Institute of Space and- Aeronautical Science, Univ.of Tokyo, Tokyo *** Dap.of Physics,- Kanagava University, Kanagava **** Dap.of 'Phyaics.Aoyama-unlveraity, Tokyo ***** Dap.of "Physic*, Nagoya University, Nagoya # Dap.of Physics, University of Washington,'Seattle ## Cosmic Kay Lab. Johnaon Space Center, NASA Houston Theoretical • Experimental Q Both Q He analyse the high energy primary electrons observed with two emulsion cjhaabera of each area 40caX50cm exposed at an average altitude of A mb for 25 hr from pleatine Texas, Sep.23, 1976. The electrons are selected to have zenith angle less than 60*, thus the otal exposure is amount, to be SflT - 8.4- 10* m2 sec str, being about three dimes of the exposure ever made by our group. I Performing the energy calibration up to 300 GeV by FNAL electron beam for the emulsion chamber of the same design, accurate analysis is made on the primary electron spectrum at several hundred GeV region. Coordinates: 4.3. Mailing address: Jun Nishimura : Insitute of Space and Aeronautical Science, University of Tokyo,Komaba,Tokyo,Japan. Hi GALACTIC RADIO EMISSION, THE EMISSIVITY IN THE DIRECTION OF HiI REGIONS AND THE LOCAL INTERSTELLAR ELECTRON SPECTRUM W. R. Webber Physics Dept. University of New Hampshire Durham, New Hampshire 03824 The relationship between galactic radio emission and the interstellar electron spectrum is re-examined utilizing determinations of radio emissivlty in the direction of HJJ regions. In view of possible large spatial gradients of electrons as implied by the Нтт data and uncertainties in the galactic magnetic-field distribution, a new procedure is presented for determining the local interstellar elec tron spectpum. This procedure involves the normalization of the radio spectrum to the electron data at high ener gies where the solar modulation is small and leads to a lower local interstellar electron intensity than previ ously deduced. Nonthermal radio emission from the galaxy is usually attributed to syn chrotron emission from energetic electrons moving in galactic magnetic fields. In this case the radio-wave emissivity e^)^* i- ' -N where N is the coeffi cient of the cosmic-ray electron spectrum which has a spectral index q and where В designates the r.m.s. I component of the interstellar field. Thus a study of the nonthermal radio spectrum will lead to: (1) information on the cosmic-ray electron distribution in the galaxy. Also, since the cosmic-ray electron spectrum is directly measured at earth, it should In theory be pos sible to (2) compare this with the synchrotron emissivity observed for our local spiral arm near the sun to deduce the magnitude of the residual cosmic- ray modulation occurring near the sun. In this paper we plan to investigate both points 1 and 2. These relationships have been discussed numerous times in the literature in the past, but usually from a rather restricted point of view. Consider the usual procedure for estimating the local interstellar elec tron spectrum using the nonthermal radio emission. If one considers frequen cies above 10-20 MHz, corresponding to electron energies above a few hundred MeV, then the simple relation I(v) - /E(v)dr exists. I(v) is the measured nonthermal radio spectrum. One usually takes this spectrum as observed in the anticenter direction and assumes that the line of sight is ^5 Kpc and that e(v) is uniform throughout the emitting re gion. Then taking an average value of B, typically 5 pG, one arrives at an average electron spectrum. It is this electron spectrum that is assumed to be the "local" Interstellar electron spectrum to be compared with that measured at earth to estimate the effects of solar modulation (e.g., Goldstein Zt at., 1970; Cummings, 1973). We believe that the following arguments will show that this is a far too simplistic picture. 374 It is clear, for example, from measurements of the Y-ray distribution along the disk of the galaxy, that the cosmic-ray nuclei intensity (and by implication also the electron intensity) is a function of position (e.g., Fichtel eX at., 1976). Models in which the intensity of cosmic rays varies linearly with the interstellar gas density, which is In turn in equilibrium with the magnetic field, imply a cosmic-ray intensity positional dependence which depends strongly on the spiral field structure. In this paper we wish to draw attention ..to some> radio data which bears directly on the cosmic-ray electron distribution in the galaxy and apparently has been overlooked by cosmic-ray workers. In recent years several authors (e.g., Roger, 1969; Parrish, 1972; Caswell, 1976) have been able to obtain direct estimates of a more localized radio emissivity by studying nonthermal emission occurring in front of HJJ regions. The direction and distance of these Hyj regions are known and if it can be assumed that these regions are opaque to radiation incident on rnem from behind and that no radiation comes in around the sides into the beam of the observing telescope, then a measure of the foreground brightness can be converted into a measure of the average emissivity along a line of sight at that frequency. Beamwidths ^1-2* at fre quencies from 5-50 MHz have been used for this purpose and 12 Нц regions satisfying the criterion that they are opaque over a beamwidth have been studied. Several observations of the same HJJ- regions at different frequen cies have been made. The HJT regions generally lie in our own or nearby spiral arms at distances within 0.3- 3 Kpc of the sun. Figure 1 shows a map of the emissivities deduced as a function of galactic longitude. The following features are worth noting. (1) Most of the emission in the anticenter direction apparently oc curs within ъ2 Kpc of the sun (e.g., Parrish, 1972). This clearly inval- 27 idates assumptions that went into the derivation of a "local" inter stellar electron- spectrum using anticenter radio data, zs discussed earlier. (2) Emissivity variations of at least a factor of two are evi dent in the data—indicating similar variations in the electron density on a scale ^1 Kpc. In Figure 2 we show the emis Vlgaxe 1. Em'ii^v-iti&i deduced &lom sivities deduced from the radio data HJJ /1e.gj.0n. itudiei platted a& a in the direction of several HJJ- re function o& gcrfactic long-itade and gions. On this figure we also show distance. (1) the emissivity deduced directly from the electron intensity measured at a time of minimum modulation in 1965- 66. The only free parameter in this calculation is B, which is taken to be 5 pG. (2) The "average" emissivity deduced from the anticenter radio spectrum by Cummings (1973), making the assumptions noted earlier. There are several points of interest in a comparison of emissivities shown in this figure. 375 First of all, at energies <1 GeV the emissivity obtained from the elec IO-S* : ' ' ' 1 ' 1 1 » 1 I1T| T-'T tron spectrum measured at the earth к«е : Is a factor ^10 lover than that de Ф GwnNtbri. o» : duced from the anticenter data—the Ф^-, ф NGC M99 0** • difference is even greater "if the <Ь»-^^ ft ф IC «05 0»3 • ^^""*— •rtr— ф KW 113 I 30 • comparison Is made with emissivi- ~~*« ф 1С 1в05 22 tles deduced from Нц studies. At K540 r*Loco»"Em«tt*itj el^^ N — least part If not all of this dif ; (CunrnBei Ю731 ЩЩЩ?X " ference has been ascribed to solar X modulation effects and this assump I : X . tion allows one to estimate the 1 OtowvM Eteclrora magnitude of the modulation exist _ 1» 1965-6* ing In 1965-66. We also notice that w* INonnolliad 10 1972) N^ the emisslvlties deduced from Нц 3 го-' *ч. ', s studies are generally a factor of 2 ^ 44" higher than those deduced from the 02 0.3 04 0.6 0Л 10 IS 2 3 anticenter studies. This is under Ekclnti Емгду ie«V)'tor'B*5pO stood if we accept the HJJ results I0« Ю ю" K as indicating that most of the emis Frtqixncy (MHl) sion in the-anticenter direction comes from within 2 Kpc rather chin the 5 Kpc used In the calculations involving the anticenter spectrum. F-сдилг I. SynchfLO&ixm eml&i4.v-LtieA •in inteMteZtat ipace. de.dac.zd iiwm At energies >2 GeV, corres (I) Me.&&wiment& -in the diiection o$ ponding to frequencies >4X1C2 MHz, HJJ лед-iono, (21 Me.atLLizment& o( the the emissivitles from the radio tsstnZ. nadio in£e.Ablty -in the. anti- data are still a factor of 2-4* слпХял dUie.ation and {3) Мга&илетп£& higher than those deduced using the o& the еЛгеЛяоп &рге£>'шт cut easrth. cosmic-ray data. Here solar rodulation effects can no longer be invoked, be cause according to conventional modulation theory the residual modulation of electrons near the sun should only be 'Ч.О-гОХ at these energies. This dis crepancy has been pointed out In the past (Setti and Woltjer, 1971; Cowsik and Mittledorf, 1974) but has apparently been overlooked by those strictly Inter ested only in the modulation effects on electrons at lower energies. Recently this discrepancy has been pointed out again in papers by Badhwar Assuming that both the cosmic-ray and radio data are basically correct, there are two possible explanations of the discrepancy. One is that the cos mic-ray flux near the sun is low by a factor of from 2 to 4 as compared with the average flux within 1-2 Kpc. The other is that the magnitude of В is greater than the value t<5 pG typically used in the calculation. Since В en ters into the eaissivity formula *v>B2 a factor $2 is needed here to brfng the two seta of data Into agreement. It is also possible that a combination-of these two effects Is at work. t 376 Яик la no avldeac* from the Шц region data that ааагеу ratios* vlthU 0.5 K>c of the «am glva «miliar amlaalvltle* than thoaa further away *W- •var, thoaa 1ц raglooa >1 Kac fro» tha sun generally lta la othar apiral ana and therefore also preaumably lnclud* tha much lower aalaalvlty of a» late гага region. At the praaeaC tta» It la not poaatbla to retch a uaaful coeclualoa on thla point froa tha *JJ data. Turning to arguaenta Invoking the 1 field, we note that Cow*Ik and Mlt- tledorf (1974) hare already auggaated that If correlated lncreaaea In • are aaaoclated with lncraaaea In the relativlatlc electron flux, then tha average synchrotron —laalvlty nonllnearly lucre**** with B. According to thee* au thor* , thla affect could enhance tha radio emlaalon by a factor of up to 10 froa that expected ualng the uaual asauaptlona regarding B. Freler tt at. (1977) have pointed out that alnce c(v) la roughly \Ш2 only a aaall percentage of the line of eight In any direction needa to occur In relatively high В field* to provide the required Increase In eaiaalvlty, even assuming that the electron flux Is unrelated to the magnitude of B. They show that consistency between the measured emlsslvlty and that predicted froa the electron spectrum can be obtained If It Is assumed that dense cloud» In which, the value of Bx Is ^70 uG occupy t<2Z of the llne-of-slght distance. We suggest here a very simple explanation for the discrepancy which In volves the magnitude of В and how It is interpreted. Most recent reviews of the interstellar magnetic field seem to agree that the average magnitude of this field, at least In the spiral arms, is 2-3 yG (Whiteoak,Д874; Manches ter, 1974). This conclusion is based on studies using Faraday rotation and Zeeman splitting which measure the average field along the line of sight. In addition it Is likely that local Irregularities in the field are the same order of magnitude as the average field, so that the total average field could be t>5 uG (Joklpii and Lerche, 1959; Manchester, 1974). It is not precisely clear how these "average" magnetic-field characteristics are arrived at, how ever. For example, although each of the Individual measurements represents an average field, the so-called "average" quoted In the literature could just as well be the most probable value of the average field—rather than the average of the average. Since the synchrotron emission Is roughly MJ2, there is an Important difference between these two quantities as well as between either of these quantities and В , the appropriate value for synchrotron emission. It is reasonable to consider that the В field Is frozen Into the Interstellar gas and that therefore the В field shares the Inhomogenelty of the gas distribu tion. In other words, there should be a distribution of В fields In Inter stellar space with a most probable, average and rms value. We are not aware that such a distribution has been determined for the Interstellar field. The Interplanetary field exhibits a distribution of magnetic-field mag nitudes when averaged over yearly periods (Marlanl e£ aJL., 1975). The inter planetary field magnitude distribution for the year 1972 Is shown In Figure 3. This field distribution has a most probable value of 45 uG, an average value of 58 uG and an rms value of 72 |iG. For example, the rms value is 1.60 x the most probable value, so that the synchrotron emission would be 2.56 x that calculated if the most probable value was used. 577 it 1» likely that the distribution of В field» «loot a line of sight in interstellar space would be at least as broad as that «een for the Interplane tary field, and probably broader, elfce several localized • field distribu tions would contribute to the overall ltae-of-elght distribution. In Figure 3 we illustrate a possible interstel lar field distribution p(B) where Bras - 2 B_. It Is also possible that this distribution has an ex tended tall or а bump at higher field strengths, similar to that suggested by Freler tt at., 1977, due to the presence of dense clouds along the line of sight, thus fur ther enhancing the synchrotron emis sion. It sees* clear fro» these arguments that it is possible to ex plain the discrepancy between radio emission and cosmic-ray data In terms of a proper description of the В T6*ol I Arbitrary Units) В field or some combination of this with a smaller electron flux near the sun than the average In the Есдиле 3. Vii&Ubuti/jn. o{> &itld local spiral arm. amplitude*: (I) Хп&лрйилгХйЛц &ield, 1912; (Z) Voitible. inteJUtettax lield Given this explanation then vheM. B, • 2 8, the question Is how to provide a 'nmi "fr "correct" interstellar electron spectrin near the sun with which to compare the spectrum measured near the earth In order to deduce the magnitude of the solar-modulation effects. We believe it Is possible to determine this "cor rect" spectrum accurately If two well-founded and rather simple assumptions are made. The first relates to the fact that the galactic nonthermal radio spectrum is very accurately known In many directions along the galactic disk. This spectral shape Is also independent of galactic longitude (Cane, 1977). It Is therefore possible to take this spectrum In a particular direction, say s 2 the anticenter direction, and using the relationship v(MHz)-1.6xlO~ Beff E 2 where Е Is the electron energy in MeV, where Beff Is defined by ^IB p<В)]|в^^О (В «j Is the most probable value of B2 p(B)) to map the correct shape of the electron spectrum near the sun apart from a normalization constant. Beff, which for our sample Interstellar field is 2 В, is here taken to be 10 uG, but this value Is not fixed since reasonable variations in В will be absorbed in the normalization process. In Figure 4 we show the galactic radio spectrum in the anticenter direction recently obtained by Webber (1977) after a careful re-analysis of all existing data and radio maps. The proper normalization to the electron data at earth is accomplished by assuming that at energies t,4 GeV the solar modulation at sunspot minimum is small (<10Z) so that the two spec tra agree at this and higher energies. This normalization is shown in the bottom part of Figure 4. For the electron spectrum at earth wc have assumed an integral flux >4 GeV of 4.5 electrons/m2-ster-eec, which is a consensus value from many experiments and agrees with the measurement of Buffington it at. (1975), which we believe to be the most accurate. This corresponds to a differential flux at 4 GeV of 2.2x10-3 electrons/mZ-ster-sec MeV. The value of Вш corresponding to this normalization Is 10 uG. It is seen that the local Interstellar electron spectrum obtained using this normalization is a factor i»2 lower than the "best" spectrum derived by Cummings (1973) , using the 37в antlrenter data directly. It la In r i мин approximate agreement with Lhe lower-limit «pactrum quoted by dim ming* In hla work. The local spec trum obtained from the radio data la now aeen to agree very cloaely with the electron apectrun measured at earth over the energy range where they may be compared (4-15 GeV) and where the electron modulation la expected to be email. Thla lower Interstellar elec tron spectrum Impllea a smaller re sidual modulation In the 1965-66 time period than deduced earlier. In a separate paper we treat this problem In some dstail. In aunmary here we report that using recent 05 см oe o» щ а г measurements of the interplanetary Etocim EiwatKMV) diffusion coefficient and Its radial dependence from solar particle figunz 4. HontheJmal ladio ipecfium studies (Zwickl and Webber, 1977), in anticenteK diMLcXion (a^tex ШеЬЬел, within the framework of the solar- /977). Hoimatization o$ thii лргсХлит modulation theory as developed by to the ob&eAve.d еЛесХлоп лресХлит at eanth thnougk the. Kelatlan v(MKz) « Gleeson and Axford (1968), we arrive 2 at a residual modulation parameter 1.6 x !0'*s ве^ Е" {МеИ, to obtain ф - 140 ± 30 MV for 1965-66. the. local inteA&tellaA elect/ion ipec- tftum, ii ikown in the. bottom oi -tfte &igwie. References Badhwar, G. D. , Daniel, R. R.-, and Stephens, S. A., 1977, Nature, 265^ 424. Bufflngton, A., Orth, C. D. , and Smoot, G. F. , 1975, Ap. J., 199, 669. Cane, H.,"1977, Ph.D. Thesis, University of Tasmania. Caswell, J. L., 1976, MNRAS, 177, 601. Cowsik, R., and Mlttledorf, J., 1974, Ap. J., 189, 51. Cummlngs, A. C., 1974, Ph.D. Thesis, Cal. Tech. Fichtel, C. E., Hartman, R. C., Kniffen, D. A., Thompson, D. J., Bignomi, G. F., Ogelman, H., Ozel, M. E., and Turner, T., 1975 Astrophys. J., 198, 163. Freier, P. S., Gilman, C. and Waddlngton, C. J., 1977, Ap. J. (In press). Gleeson, L. J., and Axford, W. I., 1968, Ap. J., 154, 1011. Goldstein, M. L., Ramaty, R., and Ftsk, L. A., 1970, Phys. Rev. Lett., 24^, 1193. \ Jokipii, J. R., and Lerche, I., 1969, Ap. J., 157, 1137. Manchester, R. N., 1974, Ap. J. 188, 637. Mariani, F., Diodata, L., and Moreno, G., 1975, Solar Physics, 45, 241. Parrieh, A., 1972, Ap. J., 174, 33. Roger, R. S., 1969, Ap. J., 155. 831. Setti, G., and Woltjer, L., 1971, Ap. Letters, £, 1251. Webber, W. R., 1976, Proc. A.S.A. 3_, !• Hhiteoak, J. B., 1974, Galactic Radio Astronomy, publ. I.A.U., p. 137. Zwickl, R. D., and Webber, tf. R., 1977, Solar Physics (in press). I 379 o., 'ijii. CuUPuJlVIui; 0.- . Ll.CYiii-i. I-'.*. ^-'.,..'..'* up CuU.lC it'i'i-- jkC.iil._rt.TLb bV rULtMii-j V.V. Uaov Space ite search Institute, Academy of liciencea, USbR. Theoretical H -.xperimental D Both D j^ulsars are repeatedly considered ao a possible oourcea of cosmic raye. But cosmic rays generated by pulsars can be investigated only in the case of NP ^432 pulfiar which caus es the Crab Nebula activity. The analysis of observable data about the Crab's emission as well as the theoretical con siderations show that the electromacnctic field in Crab represents mainly intensive low-frequency waves. In order to avoid the contradiction between the expected value of circular polarization derree und the existent upper limit it is necessary to assume that there are relativistic elect rons. The difference in energy spectra of electrons and positrons responsible for the Crab's optical and radio emissions does not exceed several per cents. In this paper the following conclusions are obtained: 1. very "young" pulsars (age Z« 1U years) are sources of ralativistic electrons una nuclei; 2. "old" pulsars ( Coordinates: OG 1.7 (i-lectrons and Positrons) 00 1.3 ^Origin and Transport of Cosmic Bays). Mailing address: Jr. V.V. Lisov _•pacе 'itesearch Institute Academy of Sciences, UjjSR 8d>, Profsojuznaja str., 117810, Moscow, USSR. 580 RADIOEMISSIOI OP THE GALAXIES »S HA LOS AHD ELECTRCBS OP COSMIC BAYS V.A. Dogiel Space Research Institute, Academy of Sciences, USSR Theoretical SI , Experimental LI BothU The observed distribution of the radioemiasion of galaxies is studied. It is shown that the characteristics of this radio- emission are in a good agreement with those ofmdiospectra, cal culated for a diffusion model of the cosmic ray distribution. The parameters characterizing the particle motion are nu merically estimated. Coordinates: OG 1.7 (Electrons and Positrons) OG 1.8.5 (Propagation of Cosmic Rays) Mailing address: Dr. V.A. Dogiel Space Research Institute Acadeuy of Sciences, USSR Profsojuznaja etr., 88, Moscow, 117810, USSR. 381 ORIGIN OF COSMIC RAY POSITRONS M. Cilar, Phytic» Department, University of Lode, J. Wdovcsyk, Institute cf Nuclear Research, Lode, and A.W. Wolfandala, Physics Department, University of Durham. New calculations are reported of the predicted energy spectrua of primary positrons froa cosaic ray - I.S.M. gas nucleus interactions. Experiaantal data on the positron spectrua are considered and corrections are applied for the effect of solar modulation. The resulting interstellar spectrua of positrons is used, together with the results of our calculations, to derive the mean path length traversed by positrons in the I.S.H. Possible interpretations are given. 1. Introduction. It is well known that the energy «pectrue of cosaic ray positrons arriving at the top of the atmosphere provides useful information about the propagation of cosmic rays in the Galaxy. Basically, one predicts what the positron spectrum should be for various propagation models and compares it with observation after correcting the measured spectra for interplanetary modulation. The implications for the models are then considered. 2. Predicted Positron Spectrum. A basic prediction is the spectrum expected for the local interstellar spectrum (I.S.S.) of nuclei interacting with 1 g cm-2 of I.S.M. Somewhat discordant predictions have been made in the past. Our own predictions (Giler et al., 1977a, referred to as I) are given in Figure 1 where comparison is made with the results of other workers. In the simple leaky box model the positron I.S.S. is divided by .the spectrum of Figure 1 to give the mean path length for positrons. Xe+. 3. Measured Positron Spectrum. In I we analysed the available data and corrected the intensities appropriately for modulation. These corrections came from the work of Aitmuhambetov et al. (1975) and Charakchyan et al. (1973). The method yielded correction factors for 1965 (sunspot minimum) of 1.65. 1.28 and 1.10 at rigidities of 0.5, 1 and 2GV respectively. The I;S.S. intensities are given in Figure 2. И2 67 ' £(W) /о So Figure 1. The production spectrum of pobitrons per second per g of I.S.M. O—O Badhwar et al. (1975) -^—. Orth and Buffington (1976) .... Ramaty (1974) J( Daniel and Stephen! (1974) •—^^» Present work. Compariaon ahowi • significant spread in predictions, particularly in the important legion 1-10 GeV. The reasons include previous neglect of kaons in SOBC cases, approximations in decay kinematics and differences in adopted spectra of protons and nuclei. The uncertainties indicated on the positron intensities include estimates of the error in the modulation corrections. The spread in points is seen to be rather large but not excessive. 563 £(«eV) Figure 2. The measured interstellar positron spectrum (corrected for modulation by us). Experimental data: | Buffington et al. (1975). Q Fanselow et al. (1969) Q Daugherty et al. (1975) ф Hartnum et al. (1975) The middle* line is the best estimate and the dotted lines are I.S.D. limits. 4. Derived Mean Path Length. The value of Ae+ derived in the manner indicated in B2 give the results shown in Figure 3. There, the range of values (I.S.D. limits) is indicated. Also shown are values for nuclei, JJJ, taken from our earlier work (I). 5. Discussion. It is apparent from Figure 3 that the values of XJJ and Ae+ are inconsistent over quite a large range of energy. Unless the corrections for modulation and/or our analysis of positron production are incorrect then there seems to be evidence for the leaky box model of propagation being incorrect. 5в4 Of 'е(<М " Figure 3. Predicted dependence of X^ on energy from positron data (with I.S.D. limits and T - 3 x 10*y). The source of the X„values is: 0 Svunaary by Shapiro and Siblerberg (1975) jj Range quoted by Meyer (1975) It Summary by Julius son (1975). T(N)R relates to kinetic energy per nucleon of the same rigidity as positrons of energy Е and T(P)R relates to protons of the same rigidity. A number of explanations spring to mind. . The first is that protons (which give most of the positrons) and heavier nuclei (whose analysis gives XJJ) ara derived from different Mora likely is the possibility of ther.e being a significant cosmic ray gradient in the Galaxy. We showed (Giler et al., 1977b) that the gradient could, bajnaduced by a single continuous source at the Galactic Centre and one-dimensional diffusion along a spiral arm with a diffusion coefficient varying with energy.in a particular way. A mare reasonable model has been put forward more recently (Giler et al., 1977d) in which sources are distributed along the spire* arms,roughly in accordance with the variation of density of stars. Then, with a simple power law J«5 v: l.iion of the diffusion coefficient, D(E), the actual fore of XN(E) am) roughly *#+(*) follows. In another model (Gilar at al., 1977c) wt have examined the situation where Che aourcea have different apactral exponents ao that the required variat n of spectral shape with position in the Galaxy follows in a natural way. The model involves a hierarchy of sources with, as energy increases, the important production cosing from: main sequence atara (probably inportant only below 1G«V), flare atara (e.g. M- and R- atara), nov*e, supernova* and pulsars. Insofar as tha •volume densities of these •tars have different dependence on galactocentre radius the required dependence of spectral shape follows in a natural manner. 4>uch a nodal remains to be evaluated in detail. Acknowledgemanta. The authors are grateful to A.N. Charakchyan, J.L. Osborne and S. Silberberg for helpful discussions. Refereneces. Aitmuhambetov, A.A., Kolomeets, E.V. and Zusmanovich, A.G., 1975, Froc. 14th C.R. Conf., Munich, 990. Badhwar, CD. et al., 1975, Astroph. and Space Science, 37, 283. Buffington, A. Orth, CD. and Smoot, G.F., 1974, Phys. Rev. Lett., 33, 34; 1975, Ap. J., 199, 669. Charakchyan, A.N. et al., 1975, Proc. 14th Int. C.R. Conf., Munich, 3,1020. Daniel, R.R. and Stephens, S.A., 1975, Space Science Reviews, 17, .45. Daugherty, J.K. et al., 1975, Astroph. J., 198, 493. Fanselow, J.L. et al., 1969, Astroph. J., 15,8, 771. Giler, M., Wdowcryk, J. and Wolfendale, A.W., 1977a J. Phys. A. (in the» press), 1977b,c,d, Astron. and Astrophys. (submitted)' Hartman, F.R. and Pellerin, C.J., 1976, Astroph. J., 204, 927. Juliusson, E., 1975, Proc. of 14th Int. C.R. Conf., Munich, V8; 2689. Meyer, J.P., 1975, Proc. 14th Int. C.R. Conf., Munich, 11, 3698. Orth, CD. and Buffington, A., 1976, Astroph. J,, 206, 312. Ramaty. R., 1974, "High Energy Particles and Quanta in Astrophysics", (Ed. F.B. McDonald and C.A. Fichtel, MIT Press Camb. Mass. and London, England). Shapiro, M.M. and Silberberg, H- - 1574, Froc. Roy. Soc, 277, 319. 586 INTERPRETATION OF THE DISTRIBUTION OF SYNCHROTRCW RADIATIO" IN THE GALAXY USJNC A NEW MODEL OF GALACTIC SPIRAL STRUCTURE C. Brindie, D.K. French and J.L. Osborne, Department of Physici, University of Durban, Durham, U.K. A new composite map of galactic spiral structure has been obtained by combining the model of tieorgelin, based on the observations of HII regions, with the map of neutral hydrogen outside the solsr circle of Verschuur. .Using models relating the distribution of galactic maguetic field and relativistic electrons to this structure profiles of synchrotron emission along the galactic plane are predicted and compared with the observed profiles a" 150 MHz and 408 MHz. Evidence that the observed.magnitude of the galactic magnetic field in the neighbourhood of the sun is more typical of an interarm rather than an arm region is interpreted in the light of information on the local synchrotron emissivity derived from low frequency surveys. 1. Introduction. Most recent interpretations of the radio coufinuum -radiation from the Galaxy have been based upon semi-empirical regular two-armed spiral models. For example Price (1974) used the pattern of Lin and Shu (1967) which was derived from density-wave theory with the parameters chosen to give agreement with the positions of the inner (Sagittarius) and outer (Perseus) arms adjacent to the sun. This pattern represents a possible 'grand design' for the Galaxy but does not include detailed features. Here we use a new composite observational model of galactic structure to make a more detailed comparison between observed and predicted distributions of radio emission. We test this by taking some 'reasonable' models for the relating of electron density and magnetic field to galactic structure and investigating whether they can account for, in the first instance, the observed profile along the galactic plane. The present work is an extension of that of French and Osborne (1976) where a more detailed discussion of the observational data is given. 2. Galactic Structure and Magnetic Field. Figure 1 shows the galactic structure used in these calculations. For distances from the galactic centre, R < 10 kpc we have taken the pattern derived by Georgelin (197Б) using HII regions as spiral arm tracers. This has a four-armed structure. The pattern places the sun in a region between two major arms and the local Prion feature is regarded as an interarm spur. When the HII regions are weighted according to their excitation parameters it is seen that their concentration is weaker in the local feature than in the major arms. Further discussion" of the significance of this local feature is given below There is no information from HII regions on the galactic structure for R < 4 kpc. ' Rather than finding arbitrary distributions of electron density and magnetic field to account for the longitude profile in the region 387 FIG. I. Galactic Spiral Structure. Galactic spiral structure used in the Calculation o/ the longitude profiles based on the spiral structures of Georgelin and Verschuur = = = using H If regions and HI regions respectively. Local feature inserted such that tangenti"l directions from the Sun are at I" = 80° and 263°. ® Location of Ue Sun at R = 10 kpc. 30° > 4* > 330° we do not consider this part of the Galaxy. For R > 10 kpc we have taken the composite 21 cm HI map of Verschuur (1973). Where the two patterns overlap theie is quite good agreement. For the overall radial dependence of the magnitude of the magnetic field in the Galaxy we adopt a form 2 2 2 H(R) «c [l - exp (-R /i,)] [exp (-R /R0 )] with R in kpc. The first term ensures that the field falls to zero at the galactic centre as would be expected from dynamo theory of field genзгation. The precise form of this is not important as we are concerned only with R > 4 kpc. R determines the rate of radial fall off in the field strength. Situ:e there is no information on this, apart from the synchrotron radiation itself, we treat it as a free parameter. We then modulate the field strength according to the spiral structure. Spiral shock theory of Roberts, and Yuan (1970)-predicts that the interstellar gas is compressed when it overtakes the density wave. The magnetic field, being tied to the gas, will be compressed in.the same ratio. An assymmetric profile of density across the arm with a sharp rise on the inner edge followed by an approximately }89 exponential decrease' is predicted. At R • 10 kpc ihe compression profile- had the approximate form (4.1 exp (-13.7a/A) + 0.7). Where A is *he radia' separation of the two arras adjacent to the point being considerc-1 and a is the distance to the inner of these arms. We adopt this form lor the modulation of the field throughout the Calaxy. Theory in fact predicts a less sharply peaked but widening compression profile for R increasing bcyona 10 kpc,but for the over arms of the Calaxy the emission integrated over 'he line of sight is very insensitive to the form of the compression. The ridge lines of the arms will not follow the relatively smooth pattern we have adopted, as evidenced by irregularities in the arms of external galaxies. In our previous work (French and Osborne 1976) it wris argued that because of these irregularities a smoother (cos)2 profile might in fact be seen. . Computations have in fact shown that this is not the case unless the scale of irregularities is very small. We therefore retain the sharply peaked profile but recognise in advance that irregularities in the arms may shift the peaks through a few degrees or give multiple peaks. The galactic field in the neighbourhood of the sun has a regular component running roughly in the direction of galactic notation. There are irregularities in the field however which one can interpret as due to the superposition of an isotropically random field of approximately equal magnitude (see for example Osborne et al. 1973). VJe shall assume that over the Galaxy as a whole the magnetic field can be decomposed into two components: a regular component that runs parallel to the spiral anr<* and a random component. The magnitude of the regular component is assumed to be a constant factor F times that of the random component throughout the Galaxy. As we shall see, the resultant profile of synchrotron emission is very sensitive to the value of F so we regard this as a second free parameter. We normalise the field at the sun (R = 10) such that the regular component has.a strength of 3 v G as indicated by pulsar rotation measures (Manchester 197 3. Distribution of Relativistic Electrons. For a magnetic field strength of 3uG synchrotron emission between 100 and 400 MHz is produced by electrons in the energy range 1 to 10 GeV. We assume that the intensity of electrons of energy Б throughout the Galaxy can be separated into spatial and energy dependent parts. I (E,R)dE - IQ(R) E"Y dE We have noted that the strong similarity between the 150 MHz and 408 MHz profiles is evidence for this. From the summary spectrum of the electron density measured at the earth of Meyer (1974) we obtain у » 2.6 and IQ (R - 10 kpc) = 80+|g (m2 s sr Gev)"l. For the spatial dependence of the intensity we take two extreme cases in order to see if, given the optimum values of the two free parameters, RQ,F, describing the magnetic field, a comparison of the predicted and observed profiles can differentiate between them. As one case we take the electron intensity to be constant for 4 < R < 15 kpc; for R - 15 kpc it is zero. We call this model A. If, on_ the other hand, the cosmic rays are strongly scattered throughout the disc they will be tied to the gas so that as it is compressed by the density waves I0«c H, . It is also probable that the rate of input of cosmic rays from discrete sources is correlated with the rise in the gas density with the result that the electron intensity уч is proportional to some higher power of the magnetic field. For Model D, then, we take I0 •< H? again with zero intensity for R > 15 kpc. 4. Observations. Profiles of the continnum radio emission at 150 and 408 MHz at b1 0° are shown in Figures 2c and 2d. That at 150 MHz is fiom the surveys compiled by Landocker and Wielebinski (1970) which has for 13 resolutions (HPBW) of 2°2 x 2?2 J. - 5°-6° and гП. i85°-245°, 3?5 x 3?8 for I11- 245°-5° and 5°(dec) x 1?25 (R-A.) for J.1U 60°-]85°. That at 408 MHz is from the survey of Green (1974) for iIX= 200°-50°, effective beamwidth 2.86 x 2.86 averaged over |ЬЛ| < 3° and {,tv - 0?5, and from the survey of Seeger et al. (1965) for £П-5о"-300°, Hl'BW 2?2 x 1?7. We have corrected both profiles for thermal emission and galactic spurs (loops) as shown in Figure 2. An extragalactic background of 50 K at Galactic loncjilude J-'H;. ?.. (a) Contributions to the observed brightness temperature profiles at 150 Mils Ji'im the Galactic Spurs. Loop I 'North Polar Spur); Loop II (Of/is Arc); Loop III. (I:) Contribution from Thermal Emission at 150 ЛШхг. (r) Oh~ sirred continuum emission profile at 150 MHz uitli thermal and spur emissions subtracted. Partialis of the observed profile due to sources, Cas A, tile. Cygnus Complex and the Vela complex. (•'•"). The l.v.t Ivo ролкь have often brio lt> HrJ j« luni- Ici.'iiiM i>l tin- l.iijc seule (structure at thj Cai.i>.y but wc .at'end that >l.v <м.м:ч'оп it. mainly localised and only a «nail rnhanccoctu in the C.'il.iilir i .iJi.ilion is possible in these direction! S. The Predicted Profiles. In terms of brightness temperature the synchrotron omisr.ivity at frequency v varies as c (v)» J0lli('**l# 2 y-(Y* J)/', Koi an isotropic irregular field of magnitude H/F one has HX("I*1V2 • 0.5 Tii [Г((-у*5)/А)/Г((>+7>/М] (H/F)(>+1>'2. Thus substituting i»2.6 c(v) - 6.17 x 10* I C(!l sin 0)1,8 + 0.6B6 (ll/r)1,8]«"2'8. K kpc"1 where 6 is Che angle between the line of sight and the direction of th^ regular field running parallel to the spiral arms, H is in uG and v is in MHz. Predicted profiles are shown in Figs. 2 and 3 while table I summarises the variables giving best agreement with the obsc-rvat'ins. From tig. 3(a) it is immediately apparent that the spiral arms show up only in emission from the irregular component; for the regular component a long line of sight in a spiral arm is compensated by 6 being near zero. With ID and H fixed at the sun by direct observation, the values of RJj and F can be fixed by demand ing exact agreement of prediction and observation at two longitudes. We choose £=180° and 308° . The overall shape of the predicted profile can then be compared with that observed. 0I 1 1 1 1 180 90 360 270 180 Galactic longitude FiEurc 3. Calculated profile tor Hjdel A with local arm inserted at 0.5 kpc from sun and showing tl.e decomposition into contributions from regular and random fields. (Table 1 line 2), (b) Comparison of the calculated profile for Model В (Table I line 3) with tho ubscrved 150 Milz profile. 591 Tabic I Electron „ . . _2 ., 2. _ Local Enitiivicy tov*rd*f» _...... Position of tun R (kpc ) F ,D^n*« . -K Distribution o r 1вО°(К kpc l) Model A In interan». 320 0.89 52 No local feature. Model A In interarm. 172 0.92 51 Local feature 0.5 kpc diitant. Model В In weak local ana. 389 1.54 37 There remains one important unknown, i-nd thus adjustable, parameter, the degree of compression in the Orion feature, whose general direction is shown in fig. 1, and the precise position of the sun relative to it. The emission due to the regular field of the Galaxy as a whole is increased or decreased depending on whether the sun is respectively in an interarm or arm region since the absolute value of the regular field at the sun is fixed. For Model A the sun must be in an interarm position. This is a manifestation of the well-known fact that the average galactic field deduced from synchrotron emission is considerably larger than that obtained from pulsar rotation measures. For Model В an interarm position of the sun would imply too high an emissivity from the arms. The best fit, shown in fig. 3(b), is obtained if the sun is in a region of weak compression (the field strength relative to the interarm value is 1.2:0.7 compared with 4.8:0.7 for major arms). 6. Conclusions. Independent of the electron distribution, the galactic model of fig. 1 gives quite good agreement with the observed profiles. In the region 30° References. Casvell, J.L., 197G, UNRAS, Г77, 601-616. French, D.K. «nd Osborne, J.L., 1976, MNRAS, Д7, 571-8'. Ccoreelin, Y.M. nnd Ccorgelin, Y.r., 1976, Aatr. Astrophys.. £9. 57- Creen, A.J., 1974, Astr. Astrophys. Suppl. ТЛ, 267. Landecker, J.L. and Wiclcbinski, 1'., 1970, Aust.M. Pliys. Astrophys. Suppl., 1£, 1. Lin, <:.C. nnd Shu, F.H., 1967, 1AU Symp. 31, Noorduijk, p. 313. Manchester, R.H., 1974, Astrophys. J., IBS, 637. Meyer, P., 1975, Origin of Cosmic Rays, p. 233, 'cds. J.L. Osborne and A.U. Uolfendale, D. Reidcl Publ. Co. Dordrecht, Holland. Osborne, J.L., Roberts, E. and WolfcndaU, A.U., i973, J. Phys. A., £, 421-33. Price, R.M., 1974, Astron. Astrophys, _33. 33. Kobcrts, U.H. and Yuan. C, 1970, Astrophys., J., 161, B77. Sector, Ch.L. ct al., 1965, Bull. astr. Inst. Ncthcrl., J_§, 11. Verschuur, C.L., 1973, Astr. Astrophys., 2_7_, 73. 392 COSMIC ELECTRONS. GALACTIC RADIO BACKGROUND AND COSMIC RAY CONFINEMENT 6. 0. Badhwar, R. R. Daniel* and S. A. Stephens* NASA Johnson Space Center Houston. Texas 77058 USA It 1s shown from an analysis of cosmic rty electron measure ments and the radio background data thVt 1n some of the electron measurements the spectral indA Is Inconsistent with that derived from the radio data and In some others the observed Intensity s Incompatible with our present under standing of the galactic magnetic field and. the dimension of the radio emitting region. It Is found that the mtfe»|t1c field required to explain the radio flux has to be greater, than 2 uG. The observed difference 1n the radio -spectra*?' shape towards the Anticenter and Halo Minimum Is Interpreted as evidence for the magnetic field В decreasing with 2, the height above-the galactic plane, as В « Z"0,3. Finally» the radio and electron data are Interpreted within the framework of the Disk-Halo diffusion model with the diffusion coefficient D « Is Eu and В « Z"5(where 6, у and (. are constants) and shown that the most plausible Interpretation requires that (I) the electron Injection spectrum Is not a single power, law all the way down but has a flattening In the GeV range, (II) the observed steepening In the radio and electron data is a combined effect of the Injection spectrum and the first break due to continuous energy loss of electrons in space, (ill) the electron Injection spectrum in the range of roughly 5-30 GeV has an Index y0 % 2.1 - 2.3 and (1v) the three constants have values n % 0.4, f % 0.5-1 and ? % 0.3. 1. Introduction. In the present Investigation we take advantage of two recent developments which permit us to make interesting Inferences on galactic nwgrietlc fields, the confinement of cosmic rays In the Galaxy and related problems. The first of these 1s the availability of a new compilation by Webster1 of galactic background radio data from the directions of the Anti center (A) and HaTo Minimum (H). The second 1s a formulation by Bulanov and Dogiel2 of the propagation of electrons in the Disk-Halo diffusion model in which the diffusion coefficient and magnetic field are assumed to vary with Z, the distance perpendicular to the galactic plane, and the leakage time of cosmic rays to vary with energy; all of them are assumed to be power law dependences. In the first part of the paper we make use of the many experi mental determinations of the electron spectrum3"17 and the radio data toward A, to obtain bounds on the mean magnetic field; we are also able to demonstrati that some of the electron neasurements are not quite compatible with our know ledge of the. mean magnetic field and the extent of the galactic disk toward A. We then Interpret the difference In the steepening of the radio spectra toward A and H as prpviding evidence for a decrease In the magnetic field with Z. Finally, we apply the calculations of Bulanov and Dogiel2 to the radio and electron observations to deduce Information on the injection spectrum of electrons and the Z-dependence of the diffusion coefficient and the magnetic field. *NASA-NRC Sr. Postdoctoral Resident Research Associate on leave from Tata Institute of Fundamental Research, Bombay, India. 595 2. Present Status of the Electron Measurements end Their Impllotion. '*. 2.1 Radio Data Used. In Figure 1 we have displayed the completion by Webster1 of the background radio data toward A and H In the frequency range 10-8000 MHz; the data points given 1n arbitrary units by Webster have been normalized by us using the radio survey at 81.5 MHz1». It is -Interesting to note that In both directions there 1s spectral steepening and that It 1s possible to fit well powy law spectra of the type X ^V: I(v) * v before and after the steepen \ ,'• ing; the spectral Indices thus obtained are summarized 1n Table 1. In Figure 1 the power law spectra fitted at low and high frequencies are extended by the dotted lines until the points of their Fig, 1 - Radio data on the Meta intersection. In the observed back galactlc component (M) and toward ground radio spectra there also exists the Anticenter (A) and Halo (H). A' a weak component of metagalactlc origin and H' represent A-M and H-M, which 1s yet not determined accurately. respectively. Even so to assess Its effect on the galactic component, in which we are interested, we have made use of the data on the metagalactlc component (line M of Figure 1) compiled by Daniel and Stephens19 and obtained the galactic spectral shapes and Indices; the former are represented as Curves A' and h" 1n Figure 1 and the latter given 1n Table 1. 2.2 Hatching of Electron Measurements With the hadЬ Data Toward A. The radio spectral Index ад - 0.79 ± 0.08 in Table 1 for v > 300 MHz toward A implies that the electron spectrum, roughly between 4 and 20 GeV, should have a power law spectrum of the type j(E)*AE~Y where y»2.58 ± 0,16. It 1s encouraging to note that of the 15 experimental measurements Included 1n Figure 2, all but two11*»15 have values of у consistent with 2.58 within two standard deviations. Nevertheless, the values of the absolute DlitMM (пя Мм *м ht kac intensity still vary by a factor of as Fig. 2 - Calculated mean magnetic much as three at a typical energy of field plotted against the extent of 10 GeV. We therefore examine the the radio emitting region toward the implications of the latter on the basis Anticenter for various electron of the observed radio flux at 1400 MHz. measurements. In Figure 2, we have plotted for each 594 Table I Spectral indices fitted before and after the radio steepening In the direction of the Anticenter and Halo Minimum Frequency Spectral Index a ' Range 1n Without Correction for After Correction for MHz Metagalactlc Component Metagalactlc Component Anticenter (A) 10 - 200 0.3S t 0.02 0.27 s 0.02* j 400-8000 0.79 ± 0.08 0.79 t 0.08 1 Halo Minimum (H) 10 - 200 0.4S i 0.02 0.27 t 0.02* | 400-8000 0.79 ± 0.10 0.79 t 0.10 I •These errors do not include t he uncertainties in the metagalactlc component experimental measurements3"17 the extent of the гафНо emitting region for vary ing mean magnetic fields , Г (Y +2)(1-U) "1, C(Y *2)(1 ;Tg+V-l 0 0 'o 2 Yo f 2(1-2e)+5(1-u)+2(1-«)j2 V 2HO-U)+ 2ТГ6-) T (i) cosmic ray sources are uniformly distributed within the Disk of half thickness b and Inject electrons with a single power law spectrum j(E) « AE •^o. (il) the region of residence and propagation has a half thickness d where d>b; 6 v (iii) the diffusion coefficient is written as D(E,Z) = Dft(|) (f-) where D . 6 and u are constants; (iv) B(Z) = В (r) where В and £ are constants; and (v) x(E) is the mean distance traversed by electrons when they lose about half their energy due to synchrotron, and inverse Compton losses. The relations calculated from their formulation connecting the radio and electron spectral indices o and y, respectively, and the constants u, 6 and £ for various values of x(E) for the A and H directions are summarized in Table 2; these relations hold only when u 13. Hovtstadt. D.. Mayer, P., and Schmidt, P. J., Astrophyi. Litt.. 9. 165. 1971. 14. 311verberg, R., Ormes, J. F., and Balasubralimanyan, V. K., J. Geoptyi. Res., 78. 7165, 1973. 15. Neegan, C. A., and Earl, J. A., Astrophys. J., 197, 219, 197S. 16. Frcltr, P., Gllman, C, and Haddington, Г. J., Astrophys. J.. 21_3- 588 1977. 17. Hartmann, G., Hulltr, D., and Prlnct, T., preprint, 1977. 18. Purton, C. R., Hon. Not. Roy. Astron. Soc, 133, 463, 1966. 19. Denial, R. R. and Stephens, S. A., Sp?ce Scl77~Rev., 17, 45, 1975. 20. Stttl, 6., and Uoltjtr, L.. Nature Phys. Sc1.. 231. 577 1971. 21. Cowslk, R.. and Mltteldorf, J., Astrophys. J., W, 51, 1974. 22. Bulanov, S. V..and Dogltl. V. A.. Astrophys. агкГТра. Scl.. 29. 305, 1974. 39в SECONDARY POSITRONS ANO ELECTRONS IN THE COSMIC RADIATION G. P. B«dh»ar and S. A. Stephens* NASA Johnson Space Center Houston. Texas 77058 USA A ntw and Improved calculation of the secondary production and equilibrium spectrum of e1 In the Galaxy In the energy range of 1 MeV to 100 GeV has been performed. This has been done by obtaining an analytic representation of the accelerator data which describes accurately the Invariant cress-section P* t* and K* from threshold energy to about 1500 GeV. This calculation takes Into account the correct angular distribution of electrons In the decay of muons and the effect of nuclei-nuclei collisions. The contributions of beta decay е and knock-on e" have been Included. A comparison of the present calculations with earlier calculations and experiment 1s presented. 1. Introduction. The nuclear interaction of cosmic ray nuclei with 2 4 and (q - С/{~\+4т /*) and»C - CQ + C]p1 + C2p*) *NASA-NRC Sr. Postdoctoral Resident Research Associate on leave from Tata Institute of Fundamental Research, Bombay, India. >7* ' ft • «0 * ^)r [(i - ^З ?"« .XP[-B p,/(i • *i» (IB) for Z «. (1 - *£-) where Z - (1 - if) - [1 - (1 • ^-Л] * (p,/p*) end k - 3. The other consents are given In Table 1. At energies below 6.6 GeV, data on Invariant cross sections Is not available. Using the above expression, we have calculated da/dp* and this Is shown along with the data*-'.» In Figure 2. The agreement Table 1 Parameters for the Representation of Invariant Cross Section A C Cl 4 Particle mb/GeWc3 В (GeV/сП Co (GeV/c)-2 (GeV/c)"* + IT 153 5.55 0.3667 -3.5 0.8334 it 127 5.30 ... 7.0334 -4.5 1.667 K+ 8.85 4.05 2.5 K" 9.3 3.8 8.3 - Is good. We also find that the same expression applies to *~ and K~ also. The constants are given in Table 1, except that k»5. We also find that the total cross section for w and w" Is also In excellent agreement with the data. Thus we have an analytical representation of ** and K* which is valid up to -v.1500 GeV incident proton energy. In an earlier work9, we have shown that the Invariant cross section in p-nucleus collision 1s given by * ft>» m 1*"E ftC "f-{ (E 'MP - (E ft>»« (lc) where; лрр and xpN are the corresponding mean free paths, fn 1s the fraction of neutrons 1n the target nucleus and n, a constant, 1s » 0.195 ± 0.05. We have already shown9 that using the above representation, the sea level muon momentum spectrum and the charge ratio can be completely explained 1n the 1 to 5000 GeV/c range. We are thus confident that the production cross section for both ir* and K* can be accurately calculated as -£ • 2* / (Е -гЛ) р, dO, where 0 Is the ang**-ef amission. at ° P J- 3. Production Spectra, of Plons and Kaons. The production of plons or kaons per g cm-2 of hydrogen Is g1Ive\ n by do, R 1 J («"if-? (E.Ep) J (Ep) dEp (2) wher* J(E_) Is the Input proton spectrum.- Using the available experimental data, we rind tfsat the proton spectrum during the period of solar minimum In the neigh borhood of the Earth can be represented by three power law spectra 1q total 2 3 energy and art given as: 2.0 x W»Ep" '" for Ep > 70 GeV, 6.92 x 10 Ep *•* 400 for 3 GeV < En < 70 GtV and 2.81 x 10s -1 En -" for 1 < Ep < 3 GeV, tht flux wring expressed In particles (m2sr. sec. GeV)-1. Using the parameters It derived by Burger16, we have demodulated t the above spectra to obtain the proton / I spectrum 1n the Interstellar space. Using the Interstellar proton spectrum and equation (2), we first compute the a v± and K1 production spectrum 1n space and find that above about 15 GeV, the production spectra are power law with an Index of 2.75. 4. Production Spectra of e*. Electrons F1g. 1 - A plot of the * invariant arise from the decay of plons and kaons.. cross-section versus 1-x" at fixed pi The kinematics of the decay scheme has values. The lines are the calculations been discussed by a number of autKoVs. using Equation (1). Orth and Buffington2 were the first to point out that the energy distribution function used 1n all previous calcula tions was not correct. In our calcula tion we have used the energy distribu tion function *(E, Ep) due to Zatsepin and Kuz'mln11 which was derived from the V-A theory of weak interactions and is given by 32 «ИЕ, Ey) [3(1-в2)-2(3+02)(Е/Еу)](Е/Ем) gdE- for 0 i (E/Ev) £ (l-B)/2 *(E, Eu) - J x [1+ хщр а№»)-3(]п№Е»)гidE.ЩЪгH (l-S)/2 <. (E/Eu) < (l+e)/2 . CMS M0MEN1UM. p . WiiV'ci where е 1s the velocity of muon of energy Ey and Е 1s the energy of electrons from the decay of the muon. Fig. 2 - A plot of do/dp* versus p* •• we note that the expression derived by 2 for proton energies of 1.71, 2.5 and Orth and Bufflngton 1s valid only for 3 GeV. The histograms ara the В * 1. observed data and curves are calculated using equation (1). In order to obtain the production spectrum of e* in Interstellar space, we have to take Into account the chemical composition of cosmic rays 401 and of the interstellar gas. We neve assumed the Interstellar gas to be 751 hydrogen and 25% helium by mass. Tfm cosmic ray composition varies as л function of energy; the variation, however, is not strong enough to change the n/p+n significantly In the energy range of Interest here. We have used an n/p+n • .105 and have evaluated the effects of these neutrons using equation (1C). Orth and Buffington2 have shown that the ratio of e* production by all cosmic ray nuclei to just those produced in p-p collisions 1s 1.34 ± 0,09 for Interstellar space. The е produc tion spectrum 1s shorn 1n Figure 3 by curves, marked А, В, С, 0. Curve A (I and II) and Curve C (I and II) are for e+ from я+ and KT decay respectively, whereas Curve В and D are for e- from ir" and K" Fig. 3 - A plot of e* production spectrum respectively. Here I and II refer in interstellar space that includes effects to the input cosmic ray spectrum of due to cosmic ray and interstellar 2 (I) J(EP) * 2 x 10" En" -" for composition. Ep > 1 GeV and (II) for the demodulated proton spectrum. We 2 note that the maximum difference in iu' io io' HP IO flux resulting from the two spectra 1 "Г ^-TTTT ' ГЛТБ1]— > '—• M I l'| г*^ ГГГП is about 20% at -v- 35 MeV decreasing to * 10% at 1 GeV and 0% beyond 10 GeV. We also note that the e* pro duction spectrum above about 10 GeV is a power law with an index of 2.75 and the kaons contribute about 10% of the electrons. In the energy region below about 10 M| 1 IMI| MeV, knock-on electrons are the '«' dominant source of secondary nega- trons. Using the demodulated spectra i.0 and the method of Daniel and Stephens12, we have calculated the pro^ "T—ч 2.0 ySf • auction spectrum of knock-on e" which 1s shown by Curve F in Figure 3. fyy sc It was pointed out by Ramaty et al.13, 1.0 /ipi that е arising from the decay of radioactive nuclei produced In the III 1 interaction of cosmic ray nuclei with POSITION ENERGrlOVI interstellar gas would.contribute significantly to the е production below about 5 MeV. Following his Fig. 4A&B-Compar1s1on of the e-+T,/ V- and approach, but using the demodulated Q(E)xEe with Daniel and Stephens (B), spectrum, we have calculated this Ramaty (C), Perola et al. (D), Orth et contribution and it is shown in al, (E). In F1g. 5B Curve (B) 1s of Badhwar et al. 402 Figure 3 by Curve Е. There Is on* additional source that contributes equally to t* and a*. It 1s the decay of K*. whose contribution has been obtained by scaling from K- contributions. The total production spectrum, Q(E), of secondary ** -and e" can be obtained by simply adding up the various contribu tions •# e* and e" respectively. 5. Comparison With OtherCalculations. Instead of comparing the production spectrum of e* obtained above with thosi obtained 1n earlier calculations, It Is more useful 'and more sensitive to compare the ratio е /е" and the product Q(E) x E* where в 1s the spectral Index of the primary spectrum used by the respective authors. This Is because the ratio в /e" and Q(E) x EB are not sensitive to the form of the Input spectrum used. Figure 4 shows the results of various calculations. The solid lines refer to e1 from »'s and K's and the dashed lines Include knock-on electrons and e+ from radioactive decay. It Is clear that there are significant differences between the present calculator and those of others14»15. In particular, note that е /е" at energies it 5 GeV 1s not 1.27 as 1s generally assumed2'12,1Ц. We also note the e* spectra of Ramaty1'* Is very steep •compared to the present calculations. The differences between the present calculation and those of Orth and Buffington2 1n the '1.5 SeV range are due to their use of the CKP model for plon production. 6. The Equilibrium Electron Energy Spectrum. Having obtained improved production spectra of e* 1n the Galaxy which was the main goal of this calcula tion, we proceed to obtain the equilibrium spectra which depends on the model of propagation. The simplest model, which Is also consistent with the existing data on the chemical composition, is the leakage lifetime model. In such a model, the electron energy spectrum, N(E), can be obtained from the continuity equation f£ • Q(E) - {f [- (I + o£ + SE2) N] - J » 0 where т 1s the catastrophic leakage lifetime, I 1s the Ionization loss term, aE Is the continuous energy loss due to bremmstrahlung (a « 9.15xlO~6Nu,« where Nu. Is the Interstellar gas density) and eE2 1s loss due to synchroton and Inverse compton scattering, p Is numerically equal to (3.79xlO~ls H2 + 1.02 х10>~1вр) GeV-2 see'1, where H 1s the magnetic field 1n uG and p Is the total photon energy density of .7 eV cm"3. The leakage time, т, has not been deter mined with any degree of certainty. It has generally been taken to be 3.3 x 106. years, though recent observations suggest that It could be higher by as much as a factor of ten. Figure 5 presents our calculation. Note that for energies $ 1 GeV, the flux Is just proportional to the amount of matter traversed. We have also plotted the data of Stone et al.16, for e" 1n the 1 MeV to 3 MeV region and see that the source of these electrons can be understood as knock-on e". In Figure 6, we have plotted the е energy spectrum for 5 g cm'2 above 1 GeV for т • 3.3 x 106.years (Curve А), т » 3.3 x 107 (Curve C) and т(Е)»3.3 x 107 (E /E)v (Curve D) Here E„ » 3 GeV and л - 0.5 are values which fit the observed enerav variation of B/C+0 ratio well. The data points are taken from Buffing ton et "all7 and Fanselow et all8". We see that the results are consTstehl with about 5 g cm-2 of matter traversed at about 3-20 GeV (Curve B, modulated). 7. Conclusions. We have developed an elegant representation of the »* and K* invariant cross section that describes the observed data from -u 1 GeV to 403 •И 500 GeV and which has been Inde pendently shown to provide an adequate explanation of the vertical sea-level muon momentum spectra* and charge ratio u+/w". 1n the 1 to 5000 GeV/c. Using this representation, we have accurately calculated the production spectrum of e* 1n the 1 MeV to 150 GeV range. It 1s shown that the present e+ data 1n 3-20 GeV 1s consistent with 5 g cm-2 of hydrogen traversed by cosmic rays References T. G. D. Badhwar et al., Astrophys. & Spa. Scl., 37, 285 (1975). 2. C. Orth and A7"Buff1ngton, Ap. J., 206, 312 (1976). 3. Carey et al., Phys. Rev. Letts., 33. 327 & 330 (1974). 4. H. Bogglld and T. Ferbel, Annv. Rev. Nucl. Sc1., 24, 1 (1974). 5. T. Ferbel, In Proc. of the SLAC F1j. 5 - A plot of e* equilibrium Sunnier Institute on Particle spectrum. The data points-1n the 1-3 Phys., U. of Rochester, Report NeV region are the e" Cal Tech meesure- No. UR-500 (1974)(unpub11shed). ments. 6. W. B. Fowler et al., Phys. Rev., 103, 1479 (1956). 7. TTW. Morris et al., Phys. Rev., 103, 1472 (1956). 8. W. J. Flcklnger et al., Phys. Rev., 125, 2082 (1962). 9. G. D. Bacihwar et al., Phys. Rev. 015, 820 (1977). 10. 37~J. Burger, Ap. J., 166, 651 (1971). 11. G. T. Zatespln and V. A. Kuz'mln, JETP 14. 1294 (1962). . 12. R. R.~D~an1e1 and S. A. Stephens, Rev. Geophys. 4 Spa. Phys., 12, 233 (1974). ~ 13. R. Ramaty et al., J. Geophys. Re s., 75, 1141 (1970). 14. R. Ramaty, In High Energy Particles & Quanta 1n Astrophys., p.122 (1974) 15. G. Perola et al., Nuovo Clmento, »#Mfl 52B, 455 (1967). 16. ETC. Stone et al., Cal Tech, fig.' 6 - A plot of E2,75 x N(E) for Preprint (1976). positrons. A, C, D are unmodulated 17. Buffington et al., Ap. J., 199, and Curve В 1s the modulated spectre 669 (1975). for various leakage lifetimes. 18. J. I. Faneslow et al., Ap. J., 202. 265 (1975). 404 MEASUREMENT OF NEGATRON AND POSITRON SPECTRA TROM 5 TO 40 CV USING A MAGNETIC SPECTROMETER G. 0. Bidhwar, R. R. Daniel, T. Cleghorn. R. L. Golden J. L. Lacy, S. A. Stephens, J. E. Zips* NASA Johnson Space Center Houston, Texas 7705B USA Theoretical Experimental 1*3 Both D A superconducting magnet spectrometer combined with a gas Cerenkov detector and lead plate shower counter has been used to measure primary cosmic ray negatrons and positrons up to about 40 GeV. Proton contamination is reduced to less than a few percent of the expected positron flux by requiring a gas Cerenkov signal (YC * 4°) and an electron cascade of magnitude consistent with the (positive) magnetic deflection measurement. An exceedingly clean negatron separation is accomplished by selecting (i) negative curvature in the spectro meter and (ii) electron signature in the gas counter and lead plate shower counter. Each of these selections applied independently is capable of reducing contamination to less than 18 of the expected negatron flux in the energy range 5 to 40 GV. Preliminary data from a balloon flight from Palestine, Texas, in May 1976 will be presented. Coordinates: OG 1.7 Mailing address: Space Physics Branch, TC2 NASA Johnson Space Center Houston, TX 77058 USA 405 тнв НАТОНВ OP тнв OBSBBYBD «вито ВАТ вгвеянт. I. GUSSIFI01TI0I AMD TH1 RAT0BBS 01 РАЮЮХЖ PB0PAGATI0M IX SEAQB L.I.Doi I2CSAM, Moscow Region, USSR Abatraot. The oosnie ray spectrum observed la ,__ cant periods la broken Into lira intervales (1) В ~* 3-10° Ge7/nueleon, (2) 3 • 10' QeT/nncleoa S В «* 3 • 10 * QeT/aaalaen. (3) 3-10* GaY/*uelaaa\*Btf»30 MeT/nmoleon, («) 30 MeT/ana- 1еоп£В,<>1 MeY/nneleon, (5) B«*l MeY/aaelean. This elas- sifioation la baaad on aoaa phyeieal comsiderationa and observation data. Conaldaiation la given to the regalari- tiaa of propagation of partlelaa of various energies deter- •<"<«g the observed spectrum of ooamlo raja In texaa of the nodal of inhomogeneitlea with different stxuotorea and of the magnetic cloud nodal. 1. Classification of the observed coamlc ray epootrun. the cosmic ray epeetrun obserred in quiescent period* (i.e. in the absence of chronospherie flarea and Forbush-deereasea) near the Earth's orbit can be broken into five interrala (aae Pig.l): (1) kinetic energy lO^GeT^Ee 2=3-10tGeV, (2) 3*10* GeVaB* „LVCEJ гЗ-Ю2 Ge7, (3) 3*10a GeY^Bfc js-yt lfeV/nuoleon, (4) 30 MaT/nue- leon^Ev&l MeV/nucleon, (3) B*< 4l MeV/nucleon. Such divialon ia baaed on sone phyaieal conaidera- tiona and observation data. She upper boundary of interval (1) should be ~ 10" eT in case of ne- tagalaetic origin of ooanie rays (due to interactions with the 2.7 *K universal nierovava radia tion); the nunerous available ex perimental data, however, are in dicative of the existence of the particles with energies ~1080 eT in the primary cosmic rays. The boundary between intervals (1) and Fig.l The observed cosmic (2) ia characterized by the varia ray spectrum broken into tion in the power exponent-of the five energy ranges. The cosmic ray differential apaetxun shaded area shows the re from 3.2-3.З to 2.7 from interval gion subjected to solar (1) to interval (2) (this fact was modulation. first established' on the basis of EAS measurements in {!])• The boundary between intervals (1) and (2) has particular meaning in case of observations inside tbe solar system and cor responds to the upper energy boundary of cosmic ray modulation in interplanetary space established on the basis of the data of ma ny-year underground and ground- based observations /* 2j, 406 The chemical and iaotopio oompoaition and the regularities of the ooaaic ray modulation by aolar wind in the energy range (3) have been sufficiently studied and it la undoubted that in terval (3) la completely of galactic origin. Tha boundary between interval* (3) and (*») oorrapoanda to tba •'»<•« of the ooamle ray spectrum In kinetic energy/nuoleon and la probably somewhat variable «1th aolar activity. Thin boun dary separates the energy ranee of explicitly galactic origin (interral (3)) from range (4) «hoae origin la being extensively diacuaaad and ha* not bean olaar aa yet. Tba problem of commie ray origin for interval (4) la discussed in detail in /3**7 «here the following poaaible altematlvea are treated» the aolar, in terplanetary, and galactic origin. Tha boundary between energy range (4) and (5) la somewhat artificial, though it was aaauaed in the Abatract to be 1 meV/ nueleon. la tha aolar activity changea, thia boundary aay ahift to either side and the diaplaeemant nay be froa aeveral tenthw of MeV/nucleon to eaveral MeV/riueleon. The physical Meaning of thia boundary is that interval (S) ia markedly different in the ehemi- eal composition, form of energy spectrum, and mode of temporal variations from interval (4). This fact is undoubtedly indicative of different origin of cosmic rays in intervals (4) and (5). It is not excluded that the relative importance of various sources of interval (5) (low-energy cosmic ray generation in chromospherie flares and during the quiet Sun, acceleration by the interplanetary shock waves and other disturbances in solar wind, low-energy particle generation in the transient layer be tween the solar wind and galactic magnetic field) varies markedly in time thereby resulting in the shift of the boundary between intervale (4) and (5). The lower boundary of interval (5) extends, according to nu merous works (see, for example £b]}% up to energies of ~ 5.01 UeV/nucleon and, maybe, even lower, i.e. essentially coincide with the upper energy boundary of the solar wind particles. Thus, the observed cosmic ray spectrum is extended from ~10* eV/nuc- leon to ~1020 eV (since the superhigh-energy partuclea are ex clusively detected with EAS arrays, only the total energy can be determined), i.e. within ~-14 orders. 2. The features of cosmic ray propagation in the space. The problem of the cosmic ray origin in the above mentioned ener- gy ranges and the origin of the spectral form will be solved on the basis of the study of, first of all, the features of cosmic ray propagation in the space. These features are eventually deter mined by the mode of the dependence of the scattering transport path of the particles on particle rigidity. Considered below will be two models of inhomogeneities which seem to be close to those existing actually in the space. 3.2.1. Tka models of inhomogeneities with structures of dif ferent complexity. Tha. scattering of a charged particle moving A'lrmg *b» Bm-ln ГТаЛЛ «_=Q» ^Qj"! _llL *hft X-dlTeCtlOn ХГОШ-СО to °° by ah inhomoheneity тГ*(о.к(х),о)тва examined in 1964 in [b]. The value k(x) was set in the' form'b(x)=H0dF/dx, where F(x) is the flux of field h from - oo to x. If the particle scattering in 407 single interactloa with the aagnetle field lAhoaogeneity la ш- oonsidsrasls, the scattering angle at t—• •© nay be found fro* the relation of the partiele velooity components 6>e* 14/Wf - Presented im{6j was the analysis of particle soatterlng toy slac-la 1 shnnojgansltios of thro* typos t (1) C (i-2x*/J*Jtxp(3Z*/A) -typ*j щ3 Fio.2. Fig.2 Magnetic inhomogeneities with the field structure of different complexity, (a) A/vj „versus Jc/Я. ; (b)-(d) the form of the magnetic force lines in the inhomogeneities J=1-3 at hQ« HQ. The form of the field A^cjand the form of the inhomogenei- ties are shown in Pig.2. The mean scattering angle where Я, =R/300 H0 is the particle Larmor radius in the basic field H«. These results were generalized in £7] for the case of the enseable of iahomogeneities whose effective sise / are confined within \< £ Л * Ла and tiie values of h„ and the mean distance 1 between the inhomogeaeities were the following functions of the effective size Д of inhomogeneitiea: KahKwKf-, ш)*1*аЯгУ-. ™ where h02 and 1г are the corresponding parameters for the inhomo geneities of size t2 . If the space is filled by inhomogeneities of the вале size with, the characteristic distance between them 1, the trans port scattering path will be „ „ A-£'**(*/<&) > (4) where it is taken into account that the scattering is of statis tical nature and through, on the average, small angles. Substi tuting (2) in (4), considering the distributions (3), and integ rating over X from Лу to ^a , we shall obtain, according \лСЦ for the transport scattering path: Х4/Л2 Fig.3 shows the results of the calculations [Qj of the par ticle rigidity dependence of A.ei- expected accordipg to (5) for three types of inhomogeneities 0=1,2,3). The following important features of cosmic ray propagation may be inferred from (5) and from the analysis of Pig.3: (1) Ае^вс^,г«/ for the high-rigidity particles when ./З?? Д , . i.e. for R??300AaHo; (2) as R decreases, the path X&j. also decreases and becomes 1/г minimum at R=300AyM> (2J) ~ (J) as R decreases further, the path Xej- begins to increase and, in case of the particles with Po«X±% i.e. R«,3QQAi , the space is practically transparent ( Aef.—*o°). 2.2. Model of magnetic clouds. It was shown in /77 that the transport scattering path for the continuous spectrum of magnetic clouds of the form (3) would be determined by the expression (7) tyA, 409 - , jr 10' t ; , ~ ; to' "•"• ю* n' io' ю' ю' т' """A io* in' XJ* IL I u "^ ти т i i ^~ TIJ i i • • r^_ 10' **>• 10* И'3 10' Ю1 «' »'*"».*. 10"* й! 10* id1 ij' to' '0° «' *>"A 10' »V. to' «" J? Я"' Ю* «' "°Vi И* Ю* И"*' 10"1 10"' »' Я* 10' >»V« Fig.3 The expected dep«nd*Bees of At^ /G-* 1 3 oa E/300 tH0Xaf or j = 1, 2, 3; ty/l/lO"* , 10" , 10~Jj k . 4, 3, 2, 1, O, -1. 410 Fig.4 The dependence of Де/ /G on R/300 А.»Дг for the model of inhomogeneities of the type of magnetic clouds. The thick, dashed, dash-dotted, and thin lines corres pond to J5 =1.5» 1» 0.5» 0, -0.5 respectively. The results for oC=0.25» 0.5» 1» and 1.5 are presented 10 10 3 10 s а* Ях/Аг = "^ ~ • ' ч/Ь&те&**ш(1-Х^Хг)1г Да . The results of the calculations (8) on the basis of (7) are shown in Fig.4. It can be seen from the fi gure that: (i) for the particles with high rigidities RS.300A2/l2 irrespectively of the form of.inhomogeneity spectrum (i.e. at any oC andy3 ),Л^ ос Л2 ; (ii) at small R,Ae>—*-const. 5. Conclusion. The results presented above make it possible to understand the features of cosmic ray propagation in the Ga laxy and interplanetary space at high and low energies and to ex plain the behaviour of the spectrum in the various energy ranges. References 1. G.B.Khristiansen. Cosmic Rays of Superhigh Energies, publ. Moscow State University, Moscow, 1974. 2. A.A.Bishara, L.I.Dorman. Proc. 13-th Intern.Cosmic Ray Conf., Denver, v.2, p.1218, 1973- 3. L.I.Dorman. Proc. VT-th Leningrad Intern.Seminar, Leningrad, 1974, p.169. 4. L.I.Dorman. Subcosmic Rays and Their Role in Space. Report at the present Conference. 5. J.H.Chan, P.B.Price. Astrophys.J.,190,Ho.l,part 2,p.139 (1974) 6. E.N.Parker, J.Geophys.Res..69.1755 (1964). 7. L.I.Dorman. Izv.Akad.Nauk SSSR, ser.flz.,55.No.11.1832.1969. 8. L.I.Dorman, A.V.Sergeev. Izv.Akad.Nauk SSSR, ser.fiz.,40, No. , 1976. — 411 THE NATUiUi OF TOE OBfacMVBD CObUIG RAY oi'.-.CtWH.. II. IliTERVAIS (1) AWD (2)(5 10"-W^ еVI L.I.Dorman IZUIRAM, Moscow Region, U&SH Abatract. Analysis of the data on the cosmic ray >un- sotropy and spectrum taking account ol' the features of cosmic raj propagation in the Galaxy permits the following conclusions to be drawn: (i) interval (1) according to the classification given in [l] up to the highest observable energies E„~10*" eV as a whole is most probably of £alac- tic origin; (ii) interval (2) is undoubtedly generated in our Galaxy and practically does not suffer solar modula tion; (iii) the mode of the dependence of the transport scattering path A in the interstellar space on E„ is esti mated for interval (1) and (2); this dependence facilitates the unterstending of the transition from interval (1) to (2); (iv) there exists a definite relationship between the spectral shape, anisotropy, chemical composition, and transport scattering path; (v) as the energy increases up to 3-10/5 eV, the relative portion of the daughter nuclei decreases by an order, and at 10y -10 eV" it will decrease by two orders as compared with the composition at 10** eV; (vi) the found dependence of the path on E„ may be inter preted in terms of the magnetic cloud model or the model of inhomogeneities with the fields of various complexities in case of a definite combination of the paraaetersA,/^, U. , and Jb which determine the spectra of the clouds and inhomo geneities. 1. Cosmic ray origin in interval (1) (З-Ю^-Ю20 eV). In various years, and up to recently, many researchers were of the opinion that the cosmic ray particles in the superhigh-energy range (interval (1) according to the classification of fy) were mainly of metagalactic origin I"2-87 (this hypothesis was criti cally analyzed in 1963 in [9]). The following arguments favouring the metagalactic origin of the interval (1) (or its highest-energy side) were considered: (i) the absence of the known sources of such high energies (up to ~102° eV) in the Galaxy, (ii) the serious difficulties associa ted with the retention of the particles of very high energies in the Galaxy. The recent discovery and the study of the pulsars, however, have made it possible to suggest highly probable mechanisms of particle acceleration in the Galaxy up to ~10So eV /lO-147(in particular, it was argued in [12/ that the pulsars were capable of accelerating also the very heavy -nuclei up to superhigh ener gies). It is not excluded either that the particles of such high energies are generated in powerful processes taking place in the vicinities of the galactic center [ 15/. A serious argument fa vouring the galactic origin of the superhigh-energy cosmic rays is the absence of the spectrum cut-off at the high-energy side up to ~ 10*° eV (such cut-off should necessarily take place in case of metagalactic origin due to interactions with the 27°K universal microwave radiation /167). Work fl?J presents a number 4i: of additional argument» against the hypothesis оГ the metagalae- tie origin of cosmic rays, while work ?18/ argues aeriously that tae «^—«» ray» up to the higheat obaarvabl* energies are ol ga- laatio origin» 2. The anisotropy in intervals (1) and (2). Tb« anisotropy •ad mode or propagation in the Galaxy ol the superhigh-energy oosmle rajs аха or apeeial interest in connection with tba exa mined problem. The published data of the measurements of df lai— gaat-aiaa KAS with four BAS arrays at Sydney, Volcano Ranch, Ha- veraa Park, and Yakutsk have bean uaod in /"197 t0 study the dis tribution of the arrival of the j= 2-10" eV particles to the Earth. The aearch for sidereal anieotropy on the basis of the above data has givan a Talue of ~ 60% for the amplitudes of the first and seeond harmonics. The possibility was analyzed in [19) that the obtained results were due to the particle arrival from the galactic clusters (superclusters). The estimates of f19/show, however, that, if the superclusters contain from 3-10J to 104 ga laxies of the type of our Galaxy, the flux of the superhigh-ener- gy particles expected from such superclusters proves to be at least 400 times as small as the flux observed on the basis of EAS measurements. The data of measurements in the lower-energy range also presented in /197 show that the amplitude of the sidereal anieotropy (see Pig.l) in the 10 -3-10i5 eV energy range (inter- Fig.l The amplitude of the jnn. • 1 icT* sidereal cosmic ray anisotro- ш* ' -*"" **» py A sect as a function of,,, energy Е к in the 10" -10 eV range (the points with the vertical bars denoting the measurement errors)[19]. The solid line shows the de pendence on E* according to i0" 1Q" (5) and the dependence E*a'7 J? eV /DCBfc.) taking account of (5) val (1) according to the classification of /17) varies very little with energy and remains constant (~0.1%) with the peak near 19 of sidereal time, which corresponds to an inconsiderable flux of cosmic rays along the force lines of the galactic spiral field*. The time of the maximum is shifted with increasing the particle energies to 13 of sidereal time, which corresponds to the ap pearance of the* drift flux of cosmic rays with energies s 3-10/J eV from the Galaxy across the magnetic force lines. In this case, the anieotropy amplitude increases rather rapidly with energy and reaches several tens of percent at lO'-* -10*° eV. The discussed re sults on the anieotropy may be approximated by the expression •The data or various observations displayed in Pig.l show that the beat agreement with the experimental data can be obtained on the assumption that the amplitude of the sidereal anisotropy Asia in the 1Q.-3-10 ^eV range is not constant but increases approxi mately as Aj,tf°cB,£ . This corresponds to an increase of ASiU by approximately an order (from ~ 0.02% to ~- 0.2%) as energy increa ses from 10" eV to 3'10i5eV. Such regularity in the variations of Asid agrees with the mode of variations of the content of the daughter nuclei depending on the particle energies in the 109-10и eV (see below. Section 3). 413 (КьУ}-1015)0*2, if 10U»V ig. - З-Ю1'-*»/ ( Г *aid V я- 0.2* ^ * CD (B^/3-1015)0,6, if З-М^В^Ю^еУ U i5 which gives ABld^0.03% *t K=3-10 eV; Aeid*0.2Ji at Ek«=3*10 e\ 1 and Aeid«i 35% at Е^г-Ю * eV. the transport scattering Path, it can be easily shown (see in f b П5]) that any diffusion model of oosaio ray propagation in the Galaxy involves a certain relationship between the observed cos mic ray spectrum D(E4) and the total spectrum of generation in all sources P(E4), the mean penetrable amount of the interstellar matter d(E«Xdetermining the relative content of the damgbmer nuc lei of the type of Li, Be, В and some secondary isotopes which are explicitly absent from the sources), the anisotropy amplitude Ajiaf(E^), and the transport scattering path in the Galaxy: 3XE«)°*F(£H)Al4E*)' C2) where Ла(Ек)±в the transport path averaged over the entire region of particle propagation; Ac-ei*.(E*) is ***в local transport path in the region of the measurements of the anisotropy. It follows from the comparison between (1) and (4) ГЕ22, if lO^V^E^sj-lO^eV Аа1ос<Ек>°<^ - - - «> 6 15 20 /E£ , if 5 10 eV«Eki€l0 eV It should be expected that, though Ac. and Aotce. may be qualitati vely different, the mode of their dependence on SK is most proba bly the same. Then, it follows from (2) that, if the total spec- trum of cosmic ray generation ?(E«)c< EiЕ*'' , the observeobse d spectrum D(E«)ocE^ in inter^ajL (12 (3-10"6V-3-10'*"eV) and р(Ех>»сЕк** in interval (2) (5-10-/fljV). Thus, the above presented data on the anisotropy are in agreement with the assumption of the unified It should be noted that, if AC(E„) la described by the de pendence of the type (5), the following Important oomolualoa may be drawn from ff(ЗУ: the penetrable Mount of interstellar medium d at Ew~3-10 eV should be an order smaller than that at 1«~ ~10меУ. Such mode of variatioms in oi dei>endine on 8» is con firmed by the- data of the direct measurements of the cnsmical composition of the ooeaio rays in the energy rang** 3 10"ef (where it was found that ««<£«•*' ). Of oourae, it would be extremely important to verify if such trend also takes plus at higher energies (the available data on the chemical eompo* itlom in the high- and superhigh-enorgy ranges ere not reliable yet). Thus, if (5) is valid, this means that the relative portion of the daughter nuclei and secondary isotopes should rapidly de crease with increasing energy and already between intervals (1) and (2) the observed cosmic ray composition should be dose to the composition of the accelerated particles in the sources. 4. Interpretation of the found dependence ЛсСЯ ). Ям da- pendenoa or д» on JS * (determining the characteristic feature* of the apectrum, anisotropy, and chemical composition of eosmls rays in Intervals (1) and (2) found in Section 3 ••7 be understood on the basis of the results of the calculations presented in/17. Analysis of Fig.4 from fl"] shows that, as a first approxlaatlom, the expression (5) may be Interpreted in terms of the magnetic cloud model in case of the appropriate selection of the parnma- ters Лi/Я.* oL , and fi . It follows, however, in thi4 s case (in cluding the necessity that the dependence Лс°с £к ' should be maintained for at least 4 orders) 3that: (ijAJhelO"' is expli citly invalid, (ii) if Л,/Л,£ 10: , thenot^i but, sinmA*_i * £ 10 kp sl t ~ 3-10"cm and na^2-10"'gsл , already at R* 10 300 na A2*2-10 the dependence Л&«< Е£ should turn out to ba toe dependence Л^ Е5 and DCf*;°cfv4 Since this is explicitly at variance with the experimental datA, it follows that, if the .magnetic cloud modek is realised, the paxtiole» of the hignest- -energiea should be heavy nuclei (it will be noted that the above estimates are very rough since the differential approach Is Inap plicable when the mmcimm size of the clouds and me transport path are of the order of the else of the system and mors strict examination on the basis„of the kinetic theory is necessary)J , tbm transition from A to Л© «е £K°'* at Kx~3-10 *sT eerrem ponds to the variations of the spectrum exponent dence Л&°£ £«' tuma out to be Лв~ const (шев Pig.4 in ft]), while in tens of the model of field inhomogeneities of various complexities a rapid increase in Л6 is expected when the par ticle energy is below soae critical energy determined by the equality or the Larmor radiur of particles and the smallest scale of inhomogeneities. From this viewpoint, it is of great importance to appropri ately analyze the experimental data in the low-energy range, i.e. in intervale (3)-(5) /2iy. References 1. L.I.Dorman. The Mature of the Observed Cosmic Ray Spectrum. I. Classification and Propagation Features of Cosmic Par ticles. Report at the present Conference. 2. G.Cocconi. Astrphys.J.Suppl.Ber.,4, 417 (I960). 3. M. Johnson, Observatory, 90, No. 97$"» 31 (1970). 4. C.E.Fichtel. Phys.Rev.Left. ,11, No.4, 172 (1963). 5. H.L&ster. Phye.Rev., 135. No75&, 1274 (196*). 6. U.Oda. Inst.Nucl.StudyTUniv. Tokyo, 1961, No.l, p.10. 7. V.S.Berezinsky, S.I.Grigorieva, G.T.Zatzepin. Izv.Alcad.Nauk SSSH, ser.fiz.,38. No.9» 1791. 197*. 8. A.H.Hillas. Proc. 14-th Intern.Cosmic Ray Conf., ttttnchen, 1975, v.2, p.717; S.A.Colgate. Proc. 14-th Intern.Cosmic Ray Conf., Ifthchen, 1975, T.2, p.723. 9. V.L.Ginzburg, S.I.Syrovatsky. Cosmic Ray Origin, publ.Acad. Sci. USSR, Moscow, 1963. 10. V.L.Ginzburg. Comments Astrophys. and Space Phys.,1, No.6, 207 (1969). 11. J.E.Gunn, J.P.Ostriker. Phys.Rev.bett., 22, No.14, 72S (1969) 12. G.Silvestro. Letter Nuovo Cimento, 2, No7?6, 771 (1969). 13* S.A.Colgate. In: "Origin of Cosmic Rays" (eds J.L.Osborne and A.W.Wolfendale), Dadrecnt-Holland, 1975, p.425. 14. S.A.Colgate. In: "Origin of Cosmic Rays" (eds J.L.Osborne and A.W.Wolfendale), Dadreeht-Holland, 1975, P.447. 15« L.I.Dorman. Kosmich.Isaled., 7, Sb.3, 402 (1969). 16. O.F.Prilutsky, I.L.Rosental. izv.Akad.Nauk SSSH,ser.fiz.,33, No.11, 1776, 1969; K.Greisen. Pnys.Rev.Latt.,16,7*8,1966; G.T.Zataepin, Y.A.Knamln. Pis'ma ZaBTF, 4, 7вТ^966; A.M.Hillas. Can.J.Phya., 46, 9623, 1968." 17. V.L.Ginzburg. Astrophys. апЗГЗрасе Sci., 1, No.l, 125 (1968). 18. S.I.Syrovatskii. Comments on Astrophys. and Space Phys., 3, 155 (1975). 19. A.M.HUlaa, M.Ouldridge. Nature, 253, No.5493, 609 (1975). 20. R.D.Davies et al. Observatory,J3_7xl2 (1974). 21. L.I.Dorman. The Nature of the Observed Cosmic Ray Spectrum. III. Intervale (3X5) (10 -3 Ю eV). Report at the present Conferenca. 416 THE MATURE OF THE OBSERVED COSMIC RAY SPECTRUM. III. HTBRTAIB (3M5)(10*-3-10U eV) L.I.box EMMA», DHOW Region, USSR Abatragt. The nature of cosmic rays in intervals (3)- (5) •Mimmig to the classification given in /1/ haabeen analysed on the basis of the data on the chemical composi tion and the energy spectrum modulation in interplanetary •pace Including the features of particle propagation. The arguments are presented that spectral interval (3)* though significantly distorted by solar modulation, is undoubted ly of galactic origin. The upper boundary of solar modula tion is interpreted. The mode of the energy dependence of the trensport path in the interstellar and interplanetary apsee baa bean determined. It has been shown that not the magnetic clouds but the 1nhomogeneities with the field structure of various complexities should be of major signi ficance to the low-energy particle scattering. In this case the »1п*^и» size of the imhomogeneities varies from ~ 3 x x 10 " cm during solar »-»п1мп» to — 6 x 10'* cm during so lar —-rfi»- It is expected that the particles penetrating to th* interplanetary space from the Galaxy due to the rapid increase of Л with decreasing E. will contribute to inter vals (4) and (5). K 1. Chemical composition in the 10°-3-10 eV/nuoleon range and the expected dependence of A» and A*i on By. According to numerous experimental data on the chemical composition of cosmic rays, the relative content of the daughter nuclei in the 109-10" ev/nueleon range decreases with increasing energy as EV . There fore, the penetrable amount of interstellar matter dV will also similarly decrease with increasing E*. Since, according to (3) ча v from /2/, d.&°c A~(? , then AecC E* in the above mentioned energy rang*. At the same time, it follows from (4) In f6] that the stellar anleotropy amplitude A ц is (other conditions being equal)се Л5. whence A«oc B* , i.e. as the particle energy de creases, AM should also decrease and reach ~ 0.01%, which agrees with the oaaaureaent data/3,4,/* 2. Chemical composition in the 3'10'-10 _ey/nucleon energy ran«*~and the nature of the scattering elements in the Galaxy. Since, as the particle energy decreases further down to the lower •In the above mentioned energy range the ground-based measurements of the stellar anisotxopy are difficult due to the additional scattering of particles and distortion of their trajectories in interplanetary space which give rise to the problem of correct interpretation of the observation results. Therefore, the avai lable data on the stellar anisotropy in the low-energy range should be treated as rough estimates. '/ />iv />/'.'/ - ' 'rK л , . V /7 L/4 . , о:. ; L' „-.'.] .. : 1 l.,t.-- , Ш-Х/ D L' '; ' uab also Seal or." L4- -. of •Jbci - . JU s чел t:ie iiiuoii.;' uec _.i o. ..J — . •r'jv :.:'• .. , . 1Ц . AcooiMin^ t; '..•.' aaal;'3is /~p~c / Cnrrie Г caT Oi; :;гГ :2з: r<:'— /ear unner-- loOni utabi^rylHUILtG 01' lii . -.;os:lil t i-.v У ",U ..„ • ' .'u'.u /lie r:.r- , u ue ilpubl ' •' ... j -1'. -• La. :.;^ u-./maar-y 01 CL.iitiic -o,V :::w c ail". Ou in ••'• xie «ary spa oe _. ipo-p^O JtfV ;il -Aiii "o -' ilO O ^ •') bliii i, th is !J •..) ^.. la Ir y ua.ri.es iiiar- .-•aiy tm-ou^hout the i.''.-у :•••:;;:• s j _L -j.! • •'."^.les, the moauia :;.1оь . •_; «OL-u,: . i la li. . J< .*it :i In o aia 3 lXLg tills par- ;;le ripidii/ ft r..-; Я "~ ... .' 0 WO:", DU :Г< I-.• • 11 •''-'•J > W: ifrtas at lower i,e:'^ias t,ny .-.'.pvcUr-uui uiod ; ta ti 0П. io.. >'.<; o' ;V"t'4' i / :H /,'.»^- V '-"Г" The •ipppr n-.-r^y boundary oi l;.e :JO j.ai" iu-:xlu ia. io ;i ;.K :iU tj ;.o uh '.f lac t t r:a t ;,e magnii tic ini.oibOfpjneit.y spest-t-uM ia iuir.zroiwi: r.ary space is .united at the siae of Я by оеги;п A.. , tne largest scale oi ,he Lahomogeneities. According to f a], tne е le-mtua os of the sec- ."ir structure are mo;.:t probably su.cn largest scale of inoonioge- ''..lie cosmic ray modular ion in interplanetary space is very sig nificant in the studied energy rung'.: (the particle flux is modu lated by a factor of from ~ 2~y to ueveral thousands), la aecor- iance with the nuuerous theoretical ana e-.tperimtntal data, how- .ver, the amplituae of the spectral modulation ia determined ou- i.y by the rigidity й aad velocity v of tie particles. Therefore, :,ue modulation amjslitude for ail partxclef; with tne saiae ratios . •'3 and the same eaergy/nucleon (i.,o., the same R and IT ) is also :.he same, i.e. their relative chemical ooaiposition iiiSide the ii;terpiane tai-y space will be the same as that in the interstel lar space. •il.N ueiLiea. Tuerel'ore, at. a •>- *00 ica/atc velocity ol' the aol*r «lad, the complete revolution ol' а force line ul Ши regular component oi the interplanetary magnetic lie Id (in the form ol um Arctwao- des spiral) is in each 6 a.u., i.e. for four sectors .<>, -> l.b n.u. whence tbe upper boundary .-;*m,»~- 270 GeV at II - • lo'* ga. -hen moving away from the Sun, Л, is almost invariant, but K decrea ses (approximately inversely to the dietnuc«» r to the Sun) and, therefore, Еим, (г) ~ 270(г/гл) ' , i.e. the siae ol' the effective modulation region decreases with increasing the particle energy. According to Fig.3 »t j=l and Fig.4 from [I], at ь1*.'л «««*• ш transport path in interplanetary spaceЛ=с f' and, since the mo dulation depth "СЛ1 at the ultrarelativistic energies, it will be <*K~K which explains the results of/5-ti7. At K«< Е*«•».», the de pendence of Л on EK is determined [1] by the parameters A t/Az , d., and fi> ; and (as it can be seen from Fig.5 &t j=l sjid Fig.4 irom /l7) the appropriate selection of these parameters may en sure that Л°С£/'' in a broad energy range from E«*«* to ~ J GeV. 4. Solar modulation of protons, electrons, and nuclei with various 2; the mode of the dependence of л on particle rigidf^ ty R. Vast experimental material has been accumulated at present ior the spectrum modulation and the radial gradient of protons, electrons, and various nuclei during a period of about one solar activity cycle (and during almost four 11-year сус1вз for moderate energies from ground-based observations). Analysis of these data on the basis of the modern theory of modulation including the dif fusion, convection, and adiabatic cooling maKes it possible to obtain fairly complete information on the mode of the dependence of A on R in the interplanetary space for a broad energy range. The results of the theoretical calculations (both analytical and numerical, see /"97) £°r the spectrum modulation may be presented, as a first approximation, in the form a. a- - *-r -uP.otseitr=i-L>poe. ; •L'e.oSsetv"-Уеое '•"ifSsnv1*, for the protons, electrons, and nuclei Z respectively. The para meter OL in (1) is determined by the solar wind velocity, the size and geometry of the modulation region, while the particle velocities for relativistic electrons are interrelated as 2 /г &--{&myfr*, М±.=[&(Ап,,ф) У К-' (2) Using the results of determination of the interstellar elec tron spectrum D го and the data on the nonthermal galactic radio emission, the interstellar spectra of protons and nuclei Z may be determined from (1) without knowing the modulation parameters: where Vt/Vj, and Ik/14 are determined by the relations (2). The parameter CC//L(R) may be determined from the comparison between the observed spectra on the basis of (1). Such analysis was car ried out in many works (see, for exa mple,/"10-16/). it was found in [Vo]% in particular, that the coordination of numerous experi- 419 rncuta can bo achieved on th* assumption that Л should markedly increase with decreasing KK (see Fig.l) la th* low-*n*r«y пас* (:•:«;: 100 l*sV for electron* and E„* 10 MaV Tor protons) in ao*or- da*:ce with Fig. 5 from fl]. Thin means that the aula rola in th* low-energy particle scat tering in th* interplaneta ry apace should b* played not by th* magnatio olemds but by tbe inhomog*n*ltl*a with th* field structure of various complexities. Tbe maximum sis* of th* imhoao- g*n*iti*s may b* **tim*t*d on the basis of comparison г»ч"--* y——-jГ between Figs 1 and 3 from IM*¥ >C'.':.i:.i, ..imai, :...'. ТГГ'. iiiio . i' .'.'c :•, v.. , ;.. , , 1 •. . .'. '^^Mfi,".' i- 11. • :ч>с. . — ' . :. . . - n . " \ • ,t'. . , - :,v- :•, v.. , ;•.,<', 1 - -. .'.. . v wL.i:i.'., . . \. ^'^L-U'I и . •;' . /. , , _ -• » •* *• . . o :•::-:::, ' . Г . '.'Ln):i:j..ci. -i: ,. .!:• "• ..л..о :-••., •;•.. o -> C O" , j. '' o # r>r:-.. , 1^_, :o.u, i )••.:', i •. • . -ji: xi. - .-'•.>--roti. , l , :;i_'.... i•• •, i •', 422 атгаямлпс ш STOCHASTIC ACCELKRATIOJI OP COSMIC RAYS BY RADIATIOH Y.M. Charualn Institute of Space Research, U.S.S.R. Aoademy of Soienoea, Moaoow, U.S.S.R. Yu.P. Ochelkov Institute of Applied Geophysic Moacow, U.S.S.R. Tbaoratieal В Experimental • BothQ Among the cosmic ray sourceв there are pulsars, super- потае, nuclei of galaxies, quasars which are known to be characterised by the relatiYistic brightness temperatures and high radiation densities. Therefore the process of cos mic rays acceleration due to this radiation in this sourceв is тегу important. In this paper the systematic and stocha stic acceleration of relativistic electrons due to induced eomptoa scattering is calculated. The heating efficiency was found to be strongly depending on spectral distribution of radiation. In the case of broad frequency band, where radia tion has the relativistic brightness temperatures the rate of electron energy increase can be higher than spontaneous energy losses. The problem of acceleration of the cosmic rays and form ing of their energetic spectrum is one of the basic problems of cosmic rays origin. Fermi mechanism is the most appropriate mechanism of cosmic rays acceleration. The problem of this mechanism is the necessity of injection to the acceleration region of t'he particles with the energies K>£ min, i.e. the preliminary particle acceleration is required. From the energy 413 consideration* it should be apparent that auperaovae, pulaare and galaxy nuclaoua may ba the injeotora. However, the charact er of partiola acoalaration ia not clear eo far. On the other hand it haa been known that thsae aouroaa aa well aa the galaxy nuclei and quaaare are the power ra \ пло . : . • .-• v:::t ; е :i г..;, -.0. У/ . .// • t- jeat тзг1г>я '...f !i n-:io-. •••<.--:• ''••? : ••ти1 • ' >: ;' ••:;• .-. DhifE with fi'ccu °r cy w .м:. j- , .if'1, i -'i'.jL ->ri-;:vy piipe :• J). At.' ьт'' I'.-o.,, L» >'.;..'v ->.'ii*j :i • . iiiint'i»,,, r гтда l i c bud td. • IM 't:'': .'. i.itiJf;.i- >',' f '• ::•':' roi'. •_' if.- vy ti'.rr-::. ол t.n.9 inte.npii;y ^> ( v ). Let 11a coneifler Hi*;'. '\ :> t tv fir'icii-n j Jl' lh..' ! ;. '.it- wx\ \ . and BtochaBtic poceleration of electron I, rad LU t i;.i: i \v; t r: t i spectre occuri'ig, under the aetrophys 1 ca i < or;a i tiony. [ll til»,- next formulae p - the nun at;on energy ;• :ini. :•• w i '.'• 0I"_ •hlr r temperature I ( v> ) on frequency \) f o, _ - Thompson croos uect.i щ <• • I I' -.. •' ) '• 1- \ . * V <Л' t •••" * 1D Billion years) similar to tha result of tha Shlcago c^oup (Garclc-Munor at.al. 1973). Above Interpretation is done with great hesitation and humility due to the poor statistics. 5. Rgfergqwa. Garcia-ttunoz,M., Mason,G.M. and Simpson, J.A., 1975 (confe
1 GeV/nuc due to uncertainties In the Be/B Ratio aa a Function of Energy fragmentation parameters. Theme (maxi mum) arror Ьага ага •hone la Figure 4. Energy (GeV/nuc) Be/• Batlo The difference between complete ,ese decay and survival la t>2.8 times th« 0.182-0.249 0.240*0.0П» estimated maximum uncertainty In either 0.2*9-0.305 0.26ttO.017* valua—thua It ahould ba poaalbla to de 0.350-0.450 27310.021 termine In an abaolute aanaa even froa 0.450-0.600 28И0.017 a a In |la measurement whether leBe haa 0.600-0.800 27210.013 mainly decayed or survived (e.g., O'Dell 0.800-1.10 27310.014 U at., 1975). 1.10 -1.50 28310.015 1.50 -2.00 28110.017 The run of our data polnta below t* 2.00 -3.00 291Ю.015 GeV/nuc and the point fro* Garcia Munoc 3.00 -4.00 31510.020 U at. (1975a, b) at -HOC MaV/nuc follow > 4.00 30710.020 cioaely the expected Be/B ratio If nearly >10.5 39 10.075 complete decay of 10Be haa occurred. >11.3 GV 33610.030+ Thla agreement between prediction and > 4.80 37510.038+ experiment would aeem to aubatantiate 1 ft From dE/dx x Е analyala. the accuracy of the theoretical calcu- "•"From flight at Parana, Argentina, latlone and argue that indeed we are on Pc - 11.30 GV. the curve for *°Be decay and net the one A/Z assumed to be 2.00 for Be, for survival or aome intermediate eltui 2.14 for B. tion. A further test of this assumption would be made if the observed ratios increase at energlea >4 GeV/nuc because of the time dilation effect on the 10Be half life. (At these energies the fragmentation cross sections are essentially energy independent.) As pointed out earlier, the estimated age from direct l0Be measurements ia £l0 x 106 years. If this long age is correct it is fortuitous In the interpretation of the Be/B ratio, since almost all 10Be with energlea <4 GeV/nuc will therefore have decayed and the observations would be expected to follow the simple en ergy dependence calculated for 10Be decay as is observed. To illustrate the effects of the cosmic-ray lifetime on 10Be decay, we show In Figure 4 the expected variation of the Be/B ratio for various average cosmic-ray lifetimes In units of 10s years. As the lifetime becomes longer, the energy at which the ratio changes from the decay to survival curve moves to higher and higher energies—and becomes more difficult to measure I Our data at the highest energies is not precise enough to define a mean ingful age; however, the much more accurate data up to "v.4 GeV/nuc may be used to set a reasonable lover limit to the cosmic-ray age. If we accept the agreement between the measured ratio and predicted Ba/B ratio for no 10Ba de cay below