0!+0

NEW TECHNIQUES IN HARD X-RAY ASTRONOMY

A thesis submitted by NEIL JOHN CURWEN SPOONER

for the degree of DOCTOR OF PHILOSOPHY

fromTHE UNIVERSITY OF LONDON

and theDIPLOMA OF MEMBERSHIP OF THE

IMPERIAL COLLEGE OF SCIENCE AND TECHNOLOGY

December 1985

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For Mum, Dad, and Jane.

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ABSTRACT

Instrumentation for hard X-ray astronomy is introduced with a historical review of the principal types of detector used: proportional counters; scintillation counters; and solid state detectors. The interplay between the nature of the interactions of X-rays, the development of instrumentation, and the acquiring of astronomical knowledge, is highlighted.A summary of the present trends in new techniques for hard X-ray astronomy is given, followed by a resume of the observational results achieved with past detectors.

A promising new technique is the "escape gated" detector. It is shown that such detectors can provide high sensitivity. Monte Carlo calculations, and measurements with a multiwire proportional counter and a detector made of a scintillator and a proportional counter, are presented. Emphasis is placed on measurements of background. This is found to be =15 times lower in an escape gated instrument than in the same instrument without escape gating.

An experiment to determine the significance of high energy loss events in the active shield of the Imperial College Germanium detector is then outlined. Data is presented of the frequency and size spectrum of such events measured during a balloon flight. The events can result in periods in which the shield is ineffective. New techniques to reduce such dead time would improve detector sensitivity.

Finally an assessment is made of the novel technique of using a large area, wide aperture (2tt), detector to search for and study periodic X-ray flux from X-ray sources. This was achieved by analysis of data from a balloon flight of the Imperial College 6.3 m2 Y-burst detector. Observational results are presented for PSR0531+21, PSR0833+45,PSR1510-59, GX1+4, 4U1626-67 and OAO1653-40.

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TABLE OF CONTENTS

ABSTRACT 3LIST OF TABLES AND FIGURES 8PREFACE 11

CHAPTER 1 INSTRUMENTAL TECHNIQUES IN HARD X-RAYASTRONOMY 12

1.1 INTRODUCTION 12.1.2 FACTORS THAT INFLUENCE THE DESIGN OF

HARD X-RAY DETECTORS 131.2.1 Hard X-rays from Space 131.2.2 Interactions of Hard X-rays with Matter 151.2.3 Background Environments 17

1.3 DESIGN CONSIDERATIONS FOR HARD X-RAYDETECTORS 19

1.4 THE DEVELOPMENT OF INSTRUMENTALTECHNIQUES FOR HARD X-RAY ASTRONOMY 21

1.4.1 Proportional Counters 221.4.2 Scintillation Counters 311.4.3 Solid State Detectors 41

1.5 FUTURE PROSPECTS FOR HARD X-RAYINSTRUMENTATION 49

1.5.1 Imaging 501.5.2 Spectroscopy 521.5.3 Time-Variability Studies 55

1.6 CONCLUSION 57

CHAPTER 2 OBSERVATIONAL RESULTS IN HARD X-RAYASTRONOMY 58

2.1 INTRODUCTION 582.2 GALACTIC HARD X-RAY SOURCES 58

2.2.1 Massive Neutron Star and Black HoleBinaries 58

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2.2.2 Low Mass Neutron Star Binaries (GalacticBulge Sources) 62

2.2.3 White Dwarf Binaries and Other UnusualStars 62

2.2.4 Supernova Remnants and Isolated X-rayPulsars 63

2.3 EXTRAGALACTIC HARD X-RAY SOURCES 662.3.1 Seyfert Galaxies 662.3.2 BL Lac Objects 672.3.3 Quasars 672.3.4 Clusters of Galaxies 67

2.4 ISOTROPIC X-RAY BACKGROUND 68

CHAPTER 3 BACKGROUND IN ESCAPE GATED DETECTORSFOR HARD X-RAY ASTRONOMY 69

3.1 INTRODUCTION 693.2 THE ESCAPE GATING TECHNIQUE 703.3 BACKGROUND IN PROPORTIONAL COUNTERS 72

3.3.1 Background Detected in the Laboratory 733.3.2 Background Detected at Balloon Altitudes 733.3.3 Background Detected in Spacecraft 743.3.4 Background in Actively Shielded MWPC's 75

3.4 THE MWPC 763.4.1 Efficiency and Energy Resolution of the

MWPC 783.4.2 Background in the MWPC 81

3.5 THE ESCAPE GATED HYBRID DETECTOR 863.5.1 Efficiency and Energy Resolution

of the Hybrid Detector 883.5.2 Background in the Hybrid Detector 90

3.6 SENSITIVITY OF THE ESCAPE GATEDDETECTORS 93

3.7 CONCLUSION 97

CHAPTER 4 HIGH ENERGY LOSS EVENTS IN THE ACTIVEANTICOINCIDENCE SHIELD OF THE IMPERIAL COLLEGE GERMANIUM DETECTOR 99

4.1 INTRODUCTION 994.2 Nal(Tl) ACTIVE ANTICOINCIDENCE

SHIELDING FOR GERMANIUM DETECTORS 1004.2.1 The Scintillation Efficiency of Nal(Tl)

to Different Radiations 1014.2.2 The Imperial College Germanium Detector

and Shield 1034.2.3 Nal(Tl) Shield Dead Time Due to Cosmic

Rays 1054.2.4 The Effect of Phosphorescent States in

Nal(Tl) 1094.3 THE HEAVY ION EXPERIMENT 111

4.3.1 Detection of Large ScintillationPulses Using a Photomultiplier Operated with Low Gain 111

4.3.2 Electronics 1134.3.3 Preflight Calibration 114

4.4 FLIGHT RESULTS 1194.4.1 Flight Details 1194.4.2 Results and Interpretation 121

4.5 CONCLUSION 125

CHAPTER 5 USE OF A WIDE APERTURE DETECTOR TOSEARCH FOR AND STUDY X-RAY PULSARS 127

5.1 INTRODUCTION 1275.2 X-RAY PULSAR ASTRONOMY - THE NEED FOR

A NEW OBSERVATIONAL TECHNIQUE 1285.3 THE IMPERIAL COLLEGE 6.3 m2 DETECTOR

AND DATA 1305.3.1 The 6.3 m2 Detector 1315.3.2 Flight and Data 133

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5.4 DATA ANALYSIS 1345.4.1 Pulsar Period Prediction 1355.4.2 The Search for Periodicity in Data 1395.4.3 Interpretation of Pulsed Count Rates 152

5.5 RESULTS AND INTERPRETATION 1535.5.1 Observations of isolated X-ray Pulsars

Associated with SNR's 1545.5.2 Observations of Accretion Powered

Binary X-ray Pulsars 1775.6 CONCLUSION 194

ACKNOWLEDGEMENTS 196REFERENCES 197

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LIST OF TABLES

1.1 Future satellite borne hard X-ray detectors. 501.2 Future balloon borne hard X-ray detectors. 50

2.1 Properties of accretion powered X-ray sources. 592.2 Properties of Crab-like SNR's and pulsars. 64

3.1 Efficiency and resolution of the MWPC and EGHD. 813.2 Relative sensitivities of the detectors. 943.3 Sensitivity of the EGMWPC. 96

4.1 Flux and energy loss of cosmic rays in the shield. 109

5.1 Observation and analysis parameters. 1585.2 3a upper limit pulsed flux from Vela. 1735.3 3a upper limit pulsed flux from PSR1510-59. 1765.4 Pulsed flux from GX1+4. 1865.5 3a upper limit pulsed flux from 4U1626-67. 190

LIST OF FIGURES

1.1 Counter used by Chodil et al. (1967) (48). 261.2 OSO-7 counter. Clark et al. (197 3) ( 5 ). 281.3 HEAO-1 MWPC. Rothschild et al. (1979)(L). 291.4 Detector used by Harri et al. (1979) ( 70). 351.5 OSO-3 detector. Peterson et al. (1970)(68). 371.6 OSO-7 scintillator. Peterson et al. (1973)(74). 371.7 Imperial College LAPAD module. Harper (1983) ( 1 51+) - 401.8 MIFRASOIC detector. Baker et al. (1984) ( 8 3). 411.9 1972-076B detector. Nakano et al. (1974)(87). 451.10 UCSD Ge(Li) detector. Peterson et al. (1972) C 69). 461.11 NASA-SMM spectrometer. Nakano et al. (1976)(85). 471.12 ZEBRA telescope. Dean et al. (1983) ( 1 12). 521.13 OSSE detector. Kurfess et al. (1983) ( 1 18). 54

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1.14 FIGARO telescope. Agnetta et al. (1983) ( 124). 56

3.1 Mass attenuation coefficients of Xe. 713.2 The MWPC. 773.3 Calculated efficiency of the MWPC. 803.4 Background in the top 4 layers of the MWPC. 823.5 Background in the top layer of the MWPC. 833.6 Detection of 60 ± AE keV in escape mode. 843.7 Background in layer 1 of the MWPC gated by layer 2. 863.8 The hybrid detector. 883.9 Calculated efficiency of the hybrid detector. 893.10 Background in the hybrid detector. 913.11 Background in the detectors with/without shielding. 92

4.1 dL/dE vs. E for Y-rays in Nal(Tl). 1024.2 L vs. E for Nal(Tl) for protons and a-particles. 1024.3 Imperial College balloon borne Ge detector. 1044.4 Cosmic ray spectra measured near the Earth. 1064.5 Gain vs. EHT of the photomultiplier. 1134.6 Block diagram of the flight electronics. 1154.7 Diagram of the staircase generator circuit. 1154.8 Block diagram of the calibration electronics. 1174.9 EHT reference voltage vs. discriminator level. 1194.10 Balloon altitude profile. 1204.11 Balloon trajectory profile. 1204.12 Raw count rate from the Heavy Ion detector. 1214.13 Energy loss spectrum from the Heavy Ion detector. 1224.14 Heavy Ion count rate vs. time during the flight. 124

5.1 Imperial College 6.3 m2 y-burst detector. 1315.2 Efficiency response of the 6.3 m2 detector. 1335.3 Angular response of the 6.3 m2 detector. 1335.4 Pulse profile for period p and duty cycle 3. 1415.5 xN vs. period for rectangular pulses. 1445.6 vs. period with/without spurious data removal. 1495.7 Pulsed counts forced into the data vs. x2™- 150

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156157159159162163165167168169175182183184184186

Crab pulse profiles. Schonfelder (1983)(203).Crab pulsed spectrum. Mahoney et al. (1984) ( 188). XN vs. period for the Crab.Peak of the distribution for the Crab.Crab pulse profiles.Crab pulse profile c.f. Wilson et al. (1983) ( 176).Sensitivity of the 6.3 m2 detector to pulsars.Pulse profiles of Vela. Turner et al. (1984)(20°). Pulsed spectrum of Vela.XN vs. period for pulses from Vela.XN vs. period for pulses from PSR1510-59.XN vs. period for pulses from GXl+4.XN peak c.f. the "background" xN distribution.GXl+4 pulse period vs. epoch.Pulse profile of GXl+4.Pulsed y-ray spectra of GXl+4.

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PREFACE

This volume concerns work undertaken by the author in a programme of development and operation of instrumentation for hard X-ray astronomy by the astrophysics group at Imperial College. A variety of topics are included because a series of balloon failures resulted in the premature end of several projects before reaching their full potential.

Work by the author with the Large Area Phoswich Array Detector (LAPAD) included: taking part in the ill-fated 1983 launch campaign in Brazil; improving phoswich performance; and some data analysis. Work with the Imperial College Germanium Detector included: partial detector fabrication, assembly and testing; supervising the systems integration during the 1983 launch campaign in Canada; Design and construction of a technological piggy experiment and analysis of the data, described in Ch. 4.

Work by the author with "escape gated" detectors, included: conceiving and performing a series of experiments and Monte Carlo simulations, discussed in Ch. 3, to investigate their background rejection capabilities; and writing the projects paper. Work performed with the MIFRASOIC balloon borne detector, (in collaboration with laboratories from Milan, Frascati and Southampton), included:- fabrication of detector parts; taking part in the ill-fated launch campaigns from Texas in 1984, and Sicily in 1985; and developing and operating a piggy escape gated experiment.

Work by the author with data from the Imperial College6.3 m2 Y-burst detector, discussed in Ch. 5, involved: creating a set of computer programs to investigate periodicity in the data in order to assess its sensitivity to X-ray pulsars; and producing values and upper limits for pulsed flux from several pulsars.

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CHAPTER 1

INSTRUMENTAL TECHNIQUES IN HARD X-RAY ASTRONOMY

1.1 INTRODUCTIONX-ray astronomy has undergone remarkable advances since

its beginnings 20 years ago. Surveys by the UHURU and ARIEL V satellites, followed by high resolution images obtained with the EINSTEIN observatory have resulted in a well developed science for energies below 10 keV. In comparison progress at hard X-ray energies (E > 20 keV) has been slow.The only complete survey was undertaken with limited sensitivity by the HEAO-1 satellite(1). The number of hard X-ray sources adequately observed is relatively small. This is perhaps a surprising situation considering the importance in astrophysics of both hard X-ray continua and line emissions from astronomical objects. Such emissions are produced by processess fundamental to the dynamics and evolution of many sources^2). In many cases soft X-ray data is insufficient to determine the production mechanisms fully.

The lack of hard X-ray data has largely resulted from the inadequacies of instrumentation, compounded perhaps by a historical neglect. The difficulties arise principally from three fundamental factors; the sharply falling photon spectra of most astronomical sources for energies above 20 keV, the way hard X-rays interact with matter in a detector, and the rapid increase with energy of background contarn/aation.

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These factors, compounded by the fact that observations cannot be made at ground level, have resulted in the evolution of instrumental techniques less successful than those employed below 20 keV.

There is perhaps a complex interplay between the nature and interactions of X-rays, the development of instrumentation, and the acquiring of astronomical knowledge. It is not clear whether instrumentation develops to meet astronomical needs or astronomy develops according to the instrumentation available, or both. This chapter introduces instrumental techniques in hard X-ray astronomy with particular reference to how the interplay between the instrumentation and astronomical results has effected the nature of the instrumentation. Emphasis is placed on detector design and factors important in later chapters.

1.2 FACTORS THAT INFLUENCE THE DESIGN OF HARD X-RAY DETECTORS Three fundamental factors have influenced the design of

hard X-ray detectors for astronomy; i) the sharply falling photon spectra of most cosmic X-ray sources for energies above 20 keV, ii) the interactions of hard X-rays with matter, and iii) the increase in background contamination with energy. Each of these factors is outlined below.

1.2.1 Hard X-rays from SpaceIt is apparent that the spectra of most cosmic X-ray

sources drop sharply with energy above about 20 keV. The three physical processes most likely to produce falling

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spectrra in an astxophysical context are summarized below.(All spectra are given in terms of photon number flux dN/dE), (see Tucker (1976) ( 3) and Rybicki et al. (1979) ( 4) ) .i) Blackbody emission; This usually arises from an optically thick hot plasma. The spectral form is given by Planck's law:

dN/dE = 1.02 x 10 38E2[EXP ( (E/kT)-l) ]" 1 ph cm" 2s" eV” 1 (1.1)

where ;T is the electron temperature.ii) Power Law emission: This can arise from either synchrotron radiation due to acceleration of relativistic electrons in a magnetic field, or from inverse Compton scatters of low energy photons on relativistic electrons.Both processes result in spectra of the form:

dN/dE = AE"a ph cm"2s"1keV"1 (1.2)

where a is the photon number spectral index.iii) Thermal Bremsstrahlung emission: This can result from the deceleration of thermal electrons in the Coulomb field of other charged particles. The spectral form is given by:

dN/dE = G(E,kT)EXP(-E/kT) ph cm"2s"1keV"1 (1.3)

where G(E,kT) is the Gaunt factor. For an optically thin plasma of temperature T, the spectrum is exponential.

The photon number flux from some sources can be fitted by one or more of the spectra above but modified above a cut-off energy Ec by the empirical function EXP[(EC - E)/EF] where Ep and Ec are in keV^5).

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As a consequence of the sharply falling spectra outlined above, there is a need for detectors of high efficiency.

1.2.2 Interactions of Hard X-rays with Matter ..Hard X-ray photons emanating from cosmic sources will

impinge on the earth and interact with any matter present.The interactions will be principally with individual atoms. The probability of a flux F (ph cm"2s" eV" :) remaining from an initial flux Fq after traversing a distance d (cm) in a material of density p (gm cm-3), is given by(6):

F = FoEXP(-u(E)pd) (1.4)

where u(E) (cm2 gm"1) is the mass attenuation coefficient of the material which depends on the atomic interaction cross section. For X-rays of energy <2moc2 the main interactions are? the photoelectric effect, which is predominant below about 200 keV (depending on Z), and Compton scattering, predominant above 200 keV, (pair production can occur above 2moc2 (7)).i) The Photoelectric effect; An X-ray interacting photoelectrically with an atom will cause a photoelectron to be ejected from one of the electron shells. If the X-ray is of energy greater than the K-edge energy of the atom,Ex-edge' tlien a photoelectron from the K-shell can be ejected. The electron vacancy left in the K-shell is subsequently filled by free electron capture or, more likely, by an electron from the next highest shell, in this

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case the L-shell. This will occur with the isotropic emission of a characteristic K-fluorescence photon of energy Er, which can leave the atom with probability wR (the fluorescence yield), or be reabsorbed by the L-shell with the emission of an Auger electron. Remaining electron vacancies are filled with the emission of L,M, or N fluorescence photons and Auger electrons. (Fig. 3.1 shows the photoelectric mass attenuation coefficients for Xe, a common detector medium. See also Sec. 3.2). The cross section for photoelectric interactions above an absorption edge is: a * Z1+"5/E3 (7).ii) Compton scattering: An X-ray interacting with an atom by Compton scattering will be deflected through an angle * by one of the atomic electrons. The detailed distribution of scatter angles is determined by the Klein-Nishina formula(6) but generally forward scattering predominates. The energy imparted to the electron can range from zero up to a large fraction of the initial photon energy. The photon energy after the interaction is given by(8):

E' = EEo[E0 + E (1 - cos*)]-1 (1.5)

where Eo is the electron rest energy. If we neglect the angular dependance of E' and consider E < Eo, then it can be deduced from Equ. 1.5 that the fractional energy loss of a photon AE/E * E. Thus the lower the energy of the photon the less easily it can lose energy in a Compton interaction (Fig. 3.1 shows the Compton mass attenuation coefficients

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for Xe). The cross section for Compton scattering increases linearly with atomic number.

The nature of the two processes discussed above to some extent determines the requirements of detection media, the degree of atmospheric absorption of hard X-rays, and governs the nature of background environments.

1.2.3 Background EnvironmentsOutlined below are the major background components

present in the laboratory, at balloon altitudes (>40 km) and in spacecraft.i) Background in the laboratory: The laboratory background flux is made up of secondary cosmic rays and radiation originating from the ground and surrounding materials (8 9 10 11 12) The photon flux for energies <2 MeV is due mainly to trace quantities of Th, U, Ra and K, and is typically about 1 cm“2s"1 above 20 keV. Y-rays from the decay chains of these elements will undergo multiple Compton scattering in the surrounding materials. This results in a build up of low energy photons (see Sec. 1.2.2) which is terminated below 30 keV by the photoelectric effect thus forming a spectral peak at 50-85 keV depending on the Z of the surrounding materials. The contribution from radioactive dust and gases such as 222Rn in the air is about 10% of this radiation. Some 80% of the secondary cosmic rays are M-mesons with a mid-latitude flux of 10"2 cm"2s"1str"1 and a peak at 0.5-1.0 GeV(13 llf) . The remaining flux is mainly electrons and photons and a much

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smaller contribution of protons and neutrons.ii) Background at balloon altitude: The predominant photon and particle fluxes at 3 mbar are; i) downward moving cosmic diffuse Y-rays, ii) atmospheric Y-rays resulting from cosmic rays, iii) primary and secondary cosmic ray ionizing radiation and iv) atmospheric neutrons. Above 20 keV these contributions provide a total flux >10 cm"2s“ 1 at latitude 42°. Above 100 keV some 70% of this flux is due to atmospheric photons resulting from cosmic ray interactions with air nuclei(12). A peak in this photon spectrum at about 60 keV occurs for reasons as given in i). The spectrum above 60 keV can be fitted by a power law with index 1.8-2.0(8 15 16). The total photon contribution at the peak is approximately 0.1 cm"2s"^eV"1, a factor of about 20 times the sea level rate(12 15).iii) Background in spacecraft: Outside the radiation belts spacecraft will encounter fluxes of; i) cosmic rays, ii) earth albedo electrons, neutrons and photons, iii) diffuse cosmic X-rays, and iv) precipitating trapped electrons and protons^6 17 18). For mid-latitudes at the top of the atmosphere, relativistic cosmic ray proton fluxes of about 1 cm"2s_1 provide the major contribution. There is a strong altitude and geomagnetic latitude dependance on this flux such that it is least for a low altitude (<500 km) equatorial orbit. Spacecraft materials can produce secondary radiation similar to that in the upper atmosphere, hence further radiation consisting of cosmic ray secondary

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photons and electrons, and induced radioactivity, will be present in the spacecraft(19).

1.3 DESIGN CONSIDERATIONS FOR HARD X-RAY DETECTORSImportant detectors that fulfil the requirements for

hard X-ray astronomy are? i) proportional counters containing a Noble gas of high atomic number such as Xe, ii) scintillation counters using inorganic crystal such as Nal, and iii) solid state detectors using semiconductor material such as Ge.

Such detectors must be placed on rocket, satellite or balloon borne platforms. The atmosphere above suitable balloon altitudes (-40 km) is relatively transparent to cosmic X-rays of energy >20 keV (the K-edge of Oxygen

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occurs at -23 k&f). Within the technical constraints imposed by these platforms, and the finances available, the underlying drive of detector development has always been towards improved sensitivity(2°). The sensitivity for background limited detection of a source is given by:

S = f-1 (E)B1/2(E) [AAE( j )]“ 1/2 (1.6)

where S (ph cm" 2s" eV" 1) is the minimum detectable flux (la level), f is the fraction of incident photons detected, A (cm2) is the detector area, AE (keV) is the energy interval and T (s) is the on source observation time, half of which is spent recording the background rate B (ct cm" 2s" eV"l) .

In general, for well designed detectors, increases in f are difficult to achieve. Therefore, improvements in

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sensitivity are more often concerned with increasing A and T, and decreasing B. However, on satellites and rockets weight restrictions limit detector area, and observations from rockets and balloons are of limited duration (typically <10 min and -4 h respectively). Therefore, it is the development of background rejection techniques, usually less constrained by platform requirements, that have often dominated the development of detectors. The restricted weight of satellite payloads has tended to necessitate detectors of small area, used mainly for surveys or studies of individual sources over long periods. The larger areas and shorter observation times available with balloon borne detectors means they are best suited to sensitive short term observations of individual sources. The severe restrictions on both detector area and observation time for rocket experiments has lead to their demise.

Most background rejection techniques can be categorised as either passive shielding or active shielding. Passive shielding entails surrounding the detector with a material able to absorb a major portion of the particle flux that is causing the background counts. For a background flux of X-rays a graded shield comprising several layers of different elements is appropriate. Each layer is chosen to be an efficient absorber of any fluorescence X-rays produced in its outermost neighbour by the absorption of an incident X-ray(8 21). Active shielding rejects background energy losses in the primary detector by detecting a simultaneous

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and associated event in a shield detector. Such simultaneous events are unlikely to result from source X-rays(12). The shield signal can be used to electronically veto the background event in the primary detector.

Satellite borne hard X-ray detectors, usually with small active areas, tend to require sophisticated active shielding over 4tt if sensitivity is to be optimized. It is generally difficult to extend this 4tt shielding to large area balloon borne detectors. Active shielding for these often covers only 2n. In both systems there is usually a trade off between the gains in background shielding and the increased background that larger shields tend to induce in the primary detector. Similarly, the gains in detection efficiency obtained by increasing detector volume may be balanced by the increased background that this causes(2 12).

1.4 THE DEVELOPMENT OF INSTRUMENTAL TECHNIQUES FOR HARD X-RAY ASTRONOMY ‘Discussed in this section are the important features of

proportional counters, scintillation counters and solid state detectors. This is achieved for each detector by; i) outlining the operation and characteristics of a basic form of the detector, and ii) tracing the major steps in its development. In this way some of the factors that have Affected the nature of the instrumentation, such as the influence of observational results, can be highlighted. Emphasis is placed on detector design rather than other aspects of instrumentation such as collimators, electronics

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or data handling.

1.4.1 Proportional CountersSince their extensive use on rocket borne platforms in

the 1960's, proportional counters have been the most commonly used detectors in X-ray astronomy below about 30 keV. It is comparatively recently that their use has been increasingly developed towards detection at harder X-ray energies. This has only been possible by incorporating many features of design that have evolved in their low energy cousins.

1.4.1.1 Operation and Characteristics of Proportional Counters

i) Operation of proportional counters; The basic proportional counter consists of a gas confined to a box containing an anode wire and a cathode plane. A potential of -1000 V is applied between these electrodes. A typical cylindrical geometry is shown in Fig. 3.8. An X-ray entering through a suitable low Z window material, usually Be or Al, may interact photoelectrically in the gas (see Sec. 1.2.2). The resulting photoelectron, and any Auger electrons, ionize the gas to produce a cloud of electron-ion pairs. If no fluorescence photons escape the counter (see Sec. 3.2) then the number of electron-ion pairs produced is, on average, proportional to the energy of the incident X-ray. The applied field will cause the electrons to drift towards the anode where they gain enough energy to produce further ionization and more electrons. An avalanche of

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electrons is produced with a multiplication factor of 103-105(22). There is little noise(7). Collection of the electrons at the anode produces an electrical pulse lasting a few us with an amplitude proportional to the number of ion pairs produced and hence proportional to the energy of the X-ray.ii) Efficiency and energy resolution; Counter fillings are usually noble gases plus a few % of a quench gas (e.g.CHif/CO2) added to stabilize counter action(6). The efficiency for detection of X-rays of energy E, assuming no attenuation by the entrance window, can be approximated by:

t(Babove K-edge> = t(1 “ EXP(-x(E)pd))(1 - wR) +(1 - EXP (-t (Ek) pd) ) (1 - EXP (-T (E) pd) )wR ] (1.7)?(Ebelow K-edge* = (1 “ EXP <-t (e ) pd) ) (1.8)

where ER is the K-photon energy of the Noble gas, t (E) is its photoelectric mass attenuation coefficient, and other symbols have their previous meanings. The second part of Equ. 1.7 accounts for the possibility that a K-photon, produced with probability wR, is reabsorbed in the gas.

The energy resolution of a counter is given by(23 24):

* 236[ (F*£)£ ]1/2 (% of FWHM) (1.9)

where F is the Fano factor, which accounts for the non-Poissonian nature of the secondary process(25), and f, the dominant factor, accounts for the variation in charge multiplication(26). e is the mean ionization energy. The geometry of the counter and the anode also effect the

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resolution(27 28) . For a typical cylindrical Xe counter AE/E * 10% FWHM at 60 keV.iii) Background; Of the various background fluxes discussed in Sec. 1.2.3, the dominant components detected in a typical single anode coaxial proportional counter with a 2ir window will be diffuse background Y-rays, and cosmic rays and their secondaries (Sec. 3.3 contains further details).

1.4.1.2 Development of proportional Countersi) Use of proportional counters in X-ray astronomy: It was Geiger-Muller tubes (GM-tubes) that were deemed to be sensitive enough to perform the initial observations of X-ray sources^29 30 31 32). However, proportional counters, unlike GM-tubes, can provide energy resolution and are capable of discriminating against background events of energy different from that of the X-rays of interest(33).It was largely the latter of these benefits which quickly led to the use of proportional counters since this allowed an improvement in detector sensitivity. This was needed for observations of weaker sources but also because of the rapid advances in collimator design that occurred in the drive for better angular resolution(34 35 36 37). Spectroscopy was perhaps a lesser consideration(32 38).

The early drive for improved sensitivity was soon slowed by the weight limits imposed by rocket borne platforms. This ensured that increases in detector area were often impractical, and that observation times were

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often limited to <10 min( 5 3 3 3l+ 35 39) . However, these problems did encourage advances in background rejection techniques(39 49).ii) Passive shielding and anticoincidence guards; Early shielding techniques for proportional counters were not well optimized due partly to a lack of understanding of the background components(41). Simple passive shielding, though reducing the soft Y-ray background, could not reduce background resulting from charged particles but in fact increased it, since secondary particles were produced^32 1+2 **3). Thus active anticoincidence shielding techniques were soon adopted. Initially these took the form of a second independent "guard" counter situated behind the main detector and operated in anticoincidence ( 3 2 3 3 **** **5).Later, guard counters were often added to the sides of the main detector^6). Such systems were inefficient at eliminating Compton electrons produced in the walls of the detector by Y-rays. Some attempts to eliminate such events were made using plastic scintillator anticoincidence (see Sec. 1.4.2.1) but plastic is less efficient at detecting lower energy events (33 **** l+6).

The problem of restricted space, and the inefficiency of using separate guard counters which could not provide complete coverage of the detector sides, encouraged the development of the first multielement systems 46) . In these detectors anticoincidence guard anodes ran parallel to the main detector anode on 3 sides, separated from it by a

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metal mesh or ground wires and sharing a common gas (4+7 **8) (see Fig. 1.1). The reduction in materials used in such detectors lowered the background resulting from cosmic ray secondaries (4 6 I+9). Detectors became long and narrow to minimize the proportion of charged particles incident at the unprotected ends.

COUNTHIOOYCO UIM ATO t. G tK )

AND FO IL ASSEMBLYc

Fig. 1.1 Cross sectional views of the counter used by Chodil et al. (1967) i1*8) .Despite the improvements discussed above, residual

background events occurred due to charged particles, and Compton scatter electrons resulting from high energy y-rays incident on the detector walls. An effort to alleviate these problems perhaps helped to stimulate the development of the pulse shape discrimination technique.iii) Pulse shape discrimination (PSD): The photoelectric absorption of soft X-rays in the gas of a conventional proportional counter will produce a localised set of electron-ion pairs (see Sec. 1.2.2). However, a high energy charged particle, an energetic electron for instance, will give rise to electron-ion pairs along an extended track. The length of the track will depend on the range of

27

the particle in the gas. In the former case the electrons produced will arrive at the multiplication region simultaneously and hence yield a charge pulse of short rise, time. Electrons from an extended track will usually have a spread of arrival times at the multiplication region and therefore produce a pulse of longer rise time. Pulse shape discrimination electronics can be used to detect this difference in rise times and enable rejection of the charged particle events^ 0 5 0 5 1 ).

PSD was initially limited to about 2-20 keV because the difference in path lengths between electrons and energetic particles decreases with increasing energy^32 33 40). However, it was soon adopted by many experimenters in the late 1960 's (t+l+) partly because of the ease of applying the technique. It requires no modification to counters themselves and little increase in valuable weight(6).iv) Multiwire proportional counters (MWPC): The possibility of extending the energy range of proportional counters above 30 keV was not initially a priority (see Sec. 1.4.2)., However, the development of satellite platforms provided increased scope for proportional counter development. One development that increased the energy range of proportional counters was the use of multilayer-multigas systems in which the efficiency of X-ray detection is optimized at different energies by providing a stack of proportional counters sensitive to different energy ranges^38 4 * 52 5 3 5l+) (e.g.

28

OSO-7, see Fig. 1.2). Reduction in background could be achieved in such "multipod" detectors by running the counters in anticoincidence with each other.

Fig. 1.2 Schematic diagram of the 1-60 keV X-ray detectors on OSO-7. Clark et al. (1973) ( 5 4).

The use of multipod systems helped stimulate an interest in fully multiwire proportional counters as a means of further reducing internal background due, for instance, to' double Compton and Compton/photoelectric events^33 55). This is achieved in a MWPC because the individual wires of layers of parallel wires in the same gas are operated in anticoincidence with each other and with the 3 sides of guard anodes (see Ch. 3). The ability of a MWPC to be a position sensitive detector (see Sec. 1.5.1) certainly also provided a stimulus for their development(5 . The first large MWPC used in this field was developed for HEAO-A^1) (see Fig. 1.3). Later, the

29

development of the coaxial anode provided 5 sided end guards and thus a truly "wall-less" counter ( 6 1+9 57) .

The increased use of proportional counters on balloons appears to have arisen not because of the superior energy resolution compared with conventional scintillators (see Sec. 1.4.2.1), but as a result of the continued drive to increase sensitivities. The large areas and excellent background rejection provided by MWPClSat soft X-ray energiesf1) could be extended to harder energies by using Xe at high pressure to increase efficiency(24 25). The lower cost/weight ratio of MWPC^Salso made them ideal for the newly developing long duration balloon flights(58) .

Fig. 1.3 Cross section of HEAO-1 high energy MWPC. Rothschild et al. (1979) (*).

30

Further developments in MWPC's concentrated on increasing sensitivity by increasing active area, and by raising gas pressure and recovering K-escape photons to provide greater efficiency. Higher gas pressure also improves background rejection by pulse shape discrimination since the track lengths of photoelectrons are reduced. A typical contemporary MWPC for hard X-ray astronomy is POKER(2l+ 59) . This consists of 4 MWPC of *2700 cm2 each with 5 layers of 12 anodes and has a background rate of -3 x 10~1+ ct cm”2s“ 1keV“1 (20-180 keV).

It was the discovery of cyclotron lines in the spectra of some sources (see Ch. 2) that highlighted the

good energy resolution obtainable with large MWPC’s. This stimulated their development towards use for high sensitivity spectroscopy (2 2Lf) . One advance has been the development of "escape gating" (see Ch. 3).v) Gas scintillation proportional counters (GSPC): Interesting results have been obtained with GSPC's(60 61 62). In these counters, detection is achieved in a low field region of the gas. The collected charge is then accelerated through a high field region resulting in light scintillation that can be detected by a photomultiplier. Energy resolutions -3% FWHM at 60 keV can be obtained. GSPC's have flown on the EXOSAT and TENMA satellites and are likely to play an important role in the future as the drive for improved energy resolution continues.

31

1.4.2 Scintillation CountersScintillation counters were used in the first balloon

borne hard X-ray observations of the mid 1960's. Since then they have not only been established as the most commonly used hard X-ray detector for astronomy, but also as a popular form of active anticoincidence shielding for many forms of hard X-ray detector including proportional counters(53).

1.4.2.1 Operation and Characteristics of Scintillation Countersi) Operation of inorganic crystal scintillators; The most successful scintillation detectors use optically transparent alkali-halide crystals, usually Nal or Csl, doped with an activating impurity such as T1 or Na. Energetic electrons, arising in the material from the interaction of an incident X-ray, excite molecular, binding states, lattice states or atomic states. De-excitation of these occurs with the release of optical photons which can escape the medium. The light can be collected and amplified by a photomultiplier to produce an electrical pulse, lasting typically 0.25-1.0 us, of height proportional to the energy loss in the crystal.ii) Efficiency and energy resolution of inorganic scintillation counters: The quantum efficiency of such detectors for X-rays below about 200 keV is -1 since both Nal and Csl have high atomic numbers (max. 53 and 55 respectively). The energy resolution is poor in comparison with other detector types since e, the energy required to produce ionization, is higher. Resolutions of -30% FWHM at

32

60 keV are typical. In practice the resolution is determined by the variance of the photon transfer efficiency t « Gqd, where G is the fraction of scintillation photons which reach the photocathode, q is the quantum efficiency of the photocathode of the photomultiplier, and d is the efficiency of collection of the electrons at the 1st dynode. G is affected by the quality of the crystal and the light collecting system. The uniformity of the photocathode efficiency is also important (61+ 55) .iii) Operation of organic scintillators: The most important organic scintillators used in X-ray astronomy are the plastic scintillators. These consist of a scintillation compound bonded into an optically transparent material such as polystyrene or polyvinyl toluene (6I+) . Energy dissipated in the material by an interacting X-ray or particle, may excite the organic molecules. Subsequent de-excitation occurs with the emission of U.V. or optical photons ( 3“+) . These photons can be detected and amplified by a photomultiplier in a way analagous to that used in inorganic scintillation counters.iv) Efficiency and energy resolution of organic scintillation counters: The quantum efficiency of organic scintillators is, for most energies, much lower than for the inorganic types since the Z (-6) of the material is much lower. The scintillator light output is also typically only -20% of that for Nal. A high Compton scattering fraction means any energy deposited is difficult to measure. Thus

33

plastics cannot easily be used as the primary detection medium in a hard X-ray telescope, except for certain specialised applications (see Ch. 5). However, its easy workability and the efficiency with which it can detect charged particles, makes plastic scintillator ideal for anticoincidence shielding.v) Background in inorganic scintillation counters; At balloon altitudes the important background contributions result from: i) Atmospheric and diffuse background X-rays. These can themselves deposit energy in the detector, or may interact in surrounding materials to produce secondaries such as Compton photons and electrons which subsequently deposit energy in the detector. ii) Cosmic rays and secondaries, and the radioactive nuclides which they generate in the detector. iii) Atmospheric neutrons^15 66). In Nal these can be captured by 127i or 23Na nuclei to produce unstable isotopes that can subsequently decay with the emission of ionizing radiation(57).

Similar background contributions will arise in satellite borne scintillation detectors (depending on the satellite orbit). This is because a spacecraft can cause cosmic ray secondaries similar to those arising from the atmosphere above balloon altitudes(68).

1.4.2.2 Development of Scintillation Countersi) Use of scintillation counters in X-ray astronomy: As aconsequence of the restrictions imposed on proportional

34

counters by rocket borne platforms (Sec. 1.4.1.2), a need arose to open up new techniques to find and locate new sources(32 33). it was largely this which encouraged the use of the first balloon borne hard X-ray scintillation detectors since detectors on balloon platforms can have longer flights and be more massive. The possibility of improving spectroscopic studies with the extended energy range of scintillators, or a desire to make use of a form of X-ray detector already developed in other fields, were probably lesser incentives.

Since any detector used above balloon altitudes is limited to detecting X-rays of energy >20 keV, a high Z detection medium is required. Scintillation counters provide such a system. They have an additional advantage over proportional counters that minimum ionizing particles can generally only deposit energy greater than typical upper energy thresholds for X-ray astronomy (-MeV's).ii) Passive shielding and plastic anticoincidence; Although balloon platforms have lower weight restrictions than rocket platforms, the massive nature and expense of scintillators do still impose restrictions on detector area. Thus improvements in background rejection, as a means of improving sensitivity, quickly dominated detector design^69).

Early balloon borne scintillation counters generally used a single Nal crystal coupled to a photomultiplier and surrounded by unsophisticated passive shielding and/or plastic anticoincidence shielding^69 70 71 72 73). This

extended forward to form a cylindrical collimator (see Fig. 1.4) providing typically a 2-4 fold reduction in background from diffuse Y-rays and charged particles.

35

Fig. 1.4 Schematic of the X-ray detector used by Harri et al. (1969) ( 70>.Early scintillation counters of this sort, used with

modulation collimators for instance, succeeded in discovering a few sources, including GX1+4, and providing some positional and angular information on more luminous sources(32 33). However, the low S:N ratio often achieved suggested that hard X-ray observations should best be used to provide additional high energy information on the emissions of known sources^ 3 2 3 3 3t+ 59 72 74) . This stimulated the recognition that the mechanisms of X-ray sources could be determined by performing spectroscopy over a large energy range(26 33).The increased importance placed on spectroscopy was also

36

part of the necessary shift from survey observations to observations of specific sources^33), imposed largely by the limitations of the platforms available.iii) Active anticoincidence shields: Plastic veto shielding is inefficient at dealing with Compton scatter photons from the shield or crystal, or interactions in the crystal that result in deposition of only part of the original energy of a source photon. Massive anticoincidence shields of inorganic crystal (introduced by Frost et al. (1966)(75)) evolved largely through a need to reduce such background.The trade off between the relative benefits of large detector area or good background rejection with these shields, has since resulted in the development of two broad forms of scintillation detector. Those with small active area and -4tt shielding, generally used on rockets and later on satellites, and large area detectors with =211 shielding, generally used on balloons(6 8).

Early scintillation counters with thick active anticoincidence shields mimicked the cylindrical form of previous detectors by replacing their plastic shielding with an inorganic crystal, viewed by a photomultiplier below the main detector(68 75 76). Fig. 1.5 shows an example. The drive for increased detector area meant simple cylindrical active collimators became unacceptably heavy if adequate angular resolution was to be achieved. Angular resolution was desirable largely to maximize the S:N ratio rather than for angular information on sources^ 6 1+ 3 6 8) . it was perhaps

37

a logical step to expand the cylindrical collimator by developing the drilled honeycomb active collimator(20 69). These provide 4tt anticoincidence. Something not possible with passive collimators. Several balloon experiments adopted such active honeycomb collimators ( 69 71+) .

No T CRYSTAL ASSY SO M L BERYLLIUM ENTRANCE WINDOW

• PLASTIC PM POTTINf ROOT

rVOLTASE DIY10CR NETWORKS

LEAOS

DC-t-OOSTOPTICAL COUPLINS SREASC

C l | CRYSTAL

SC SP-IOITOPTICAL COUPLINS FLUID

0 I t* 1 *SCALE/IN

S

Fig. 1.5 Schematic of the OSO-3 satellite borne detector of Peterson et al. (1970) (68).It was partly the greater background problems

encountered in early satellite detectors, such as OSO-3 (7l+ 77) and OSO-5(78), that further inspired a particular drive for 47T anticoincidence in satellite experiments. Fig. 1.6 shows one example.

Fig. 1.6 Schematic of the OSO-7 detector with active honeycomb collimator. Peterson et al. (1973) (71+).

MBTllMi

38

iv) Phoswiches - the LAPAD detector: An important source of background in the detectors described above are components arising from Compton scattering of Y-rays in the photomultiplier and material of the primary detector, within the shield. The development of the phoswich detector enabled elimination of this material thus reducing the background and allowing an increase in area to weight ratio at reduced cost, an important consideration for detectors on both satellite and long duration balloon flights^12 20 69) .

In a phoswich detector two different scintillators, a "primary" and a "veto" crystal, with different decay time constants, are placed in optical contact with each other and a photomultiplier. Rise time discrimination of the photomultiplier signals can be used to determine whether an event deposited energy in one crystal or the other, or both. Only events depositing energy wholly in the primary crystal are accepted. Photons which Compton scatter and deposit only part of their energy in the primary are vetoed, as are all the events normally rejected by an anticoincidence shield. Thus 2ir shielding is provided.

Though phoswiches have been designed with organic/ inorganic scintillator combinations(43 79), it is the wholly inorganic Csl(Tl)/Nal(Tl) combination, first suggested by Cline et al. (1961) ( 80), that are the optimum for most purposes. In either form the veto crystal is usually situated nearest the photomultiplier. However, in detectors that use an active collimator to provide 4tt shielding, the

39

active collimator itself can be used as the veto crystal. These systems provide improved light collection from the collimator(20). This was important since as detector areas increased so did the problems of light collection in active collimators. Problems associated with radioactivity induced in active collimators by cosmic rays, spallation products, and neutrons, were also revealed(12 20 32). This helped to stimulate the use of passive collimators.

Though a passive collimator induces extra Y-ray background, the phoswich system can compensate for this. Also, less detector area is obscured than with an active collimator so sensitivity can be increased. Furthermore the system is less complex, lighter and cheaper, factors particularly important for balloon borne detectors. Several recent such systems have been flown(21 81), including the Imperial College LAPAD phoswiches (see Fig. 1.7). The systems usually consist of several separate phoswiches in a modular design since single large crystals are more difficult to fabricate^7 76).

Despite the poorer energy resolution of CsI(Na) and Csl(Tl) compared with Nal(Tl), many early phoswiches used these crystals as the primary detector due to background or cost considerations(21). It was perhaps only after an observation with a phoswich revealed discrete hard X-ray line emission from some sources (see Ch. 2), that more attention began to be paid to energy resolution. Nal(Tl) then became the natural chfioce of primary. The practical

40

limit to the energy resolution of scintillation counters is due primarily to photocathode non-uniformity(12 65 70) (see Sec. 1.4.2.1). The author has improved the resolution of various phoswiches by diffusing the scintillator light evenly over a photomultiplier tube by using light pipes coated with highly reflective PTFE paint(82). 20-25% FWHMat 60 keV has been achieved.

Fig. 1.7 Diagram of an Imperial College LAPAD module, v) The MIFRASOIC detector: A novel scintillation detector developed and flown by Southampton University and 2 Italian groups in collaboration with I.C., is the MIFRASOIC detector(16 83). This has achieved a resolution of -18% at 60 keV. The detector incorporates the technique, adopted by

41

several earlier experimenters, of using two photomultipliers to view a single Nal(Tl) crystal slab(73 78). All this is enclosed in a light tight box coated with PTFE paint (see Fig. 1.8). The photomultipliers can be run in anticoincidence to reduce noise. The box diffuses the scintillation pulses over the photocathodes and results in the good energy resolution characteristics. An optically separate veto crystal below the Nal(Tl) provides active anticoincidence against Compton events.

(a) (d)

Fig. 1.8 Cross section of a MIFRASOIC detector unit: (a)plastic anticoincidence; (b) Copper collimator; (c) graded shield; (d) photomultiplier; (e) primary Nal crystal; (f) Nal anticoincidence shield. Baker et al. (1984) (83).

1.4.3 Solid State DetectorsSolid state detectors have not been as successful as

proportional counters or scintillation counters, largely due

42

to problems of cost and availibility. This may change as the latter detectors reach a high level of maturity, and improvements in solid state detectors continue (8I+) .

1.4.3.1 Operation and Characteristics of Solid State Detectorsi) Operation of solid state detectors: X-rays incident on a solid state detector, usually of Ge or Si semiconductor, will produce energetic electrons which quickly convert their energy into electron-hole pairs. An applied electric field of a few keV cm” 1 allows collection of the charge carriers at electrodes. The total charge collected is proportional to the energy deposited in the crystal, therefore amplification of the signal enables the energy to be measured.ii) Efficiency and energy resolution: The quantum efficiency C of a solid state detector depends on the Z of the semiconductor (see Sec. 1.2.2), and on the depth of the depletion region(85 86). However, below -200 keV 5 is nearly 1, though the entrance window and surface dead layer will cause a cut-off at lower energies.

The energy resolution of solid state detectors is given by

FWHM « [(detector noise)2 + (electronic noise)2] 1/2 (1.10)

The intrinsic detector noise is 2.36(eFE)1/2 keV, (symbols have their previous meanings). For Ge, e = 2.9 eV (about 100 times smaller than for scintillators), and F - 0.13 depending on the variations in charge collection. Thus solid state detectors can achieve high energy resolutions,

43

typically about 1.0% at 60 keV (30-40 times better than scintillators(85 87)). However, since there is no intrinsic amplification in solid state detectors, the resolution also depends heavily on noise in the amplifying system(88).iii) Background: An unshielded Ge detector will be particularly sensitive to the various background components discussed in Sec. 1.2.3, however the important contributing background components are similar to those discussed for inorganic scintillators(67 89 90) (see Sec. 1.4.2.1). Fast neutrons and protons can additionally cause damage resulting irt the degradation of energy resolution(89).

1.4.3.2 Development of Solid State Detectorsi) Use of solid state detectors in hard X-ray astronomy: In the 1960's the unequalled energy resolution obtainable with solid state detectors inspired their rapid development for spectroscopy in medical and nuclear physics(9 86 92). This was despite the small size of the available semiconductor crystals , which limited detector efficiency. In X-ray astronomy at the time, any spectroscopy interests lay largely in source continuum spectra which requires only moderate energy resolution. It was this factor, asmuch as the small size and high cost of crystals, that restricted the early development of solid state detectors for hard X-ray astronomy.ii) Si(Li) and Ge(Li) detectors; The few balloon borne solidstate detectors flown before 1970 used coaxial Ge, doped

44

with Li to compensate for impurities ( 69 9 3 91+) . Unfortunately Li diffuses easily in Ge(Li) therefore these detectors require continuous cooling, even in storage. The problem is less severe in Si(Li) detectors which consequently received earlier development(95 96) . However, Si(Li) detectors have lower efficiency for hard X-rays because they have lower Z (14 c.f. 32 for Ge). Si(Li) detectors have thus been used mainly for detection of X-rays <50 kev(32 85 93) (e.g.HEAO-b (97)). Both Ge(Li) and Si(Li) detectors also require cooling during operation if background noise arising from the thermal excitation of electron-hole pairs is to be minimized. Good energy resolution has only been possible with the advent of low noise amplifiers. Of particular importance has been the development of the Field Effect Transistor(9 3).

Some early improvements in the sensitivity of solid state detectors occurred through improved crystal perfection and better crystal purity. However, in the 1970's the development of satellite platforms allowed more sensitive observations with these detectors since longer observation times became available. More importantly, theoretical low energy Y-ray astrophysics was far outstripping the exploratory observations. Specifically there developed a need to find monoenergetic line features and sharp continuum breaks predicted to occur in some source spectra(6 85 87 98 "). Such was one objective of the first solid state satellite borne detector, designated 1972-076B, of

45

Nakano et al. (1974)(87) (see Fig. 1.9). This was a Ge(Li) detector sensitive to X-rays of 40 keV-2.8 MeV.

Fig. 1.9 Schematic cross sectional view of the 1972-076B detector. Nakano et al. (1974) ( 87).

Although lines had already been observed in solar flares(85 87), no lines had been observed from cosmic sources(98 10°). Therefore there was no observational precedent to perform spectroscopy. It is possible that the objective mentioned above arose partly to justify a desire to exploit a new detector system which was developed rapidly for other fields of physics^85).

The 1972-076B detector incorporated 4.4 cm of plastic anticoincidence shielding^85 87). The plastic was ineffective at suppressing lines due to (n,Y) interactions

46

in the Ge, neutron induced events, and Compton scatter events. This stimulated the development of the more effective thick (>6 cm) inorganic crystal anticoincidence systems used on many later detectors for both shielding and collimation. Fig. 1.10 shows an example. Both Csl (Na) and CsI(Tl) shields have been used (e.g. for the HEAO-C Ge(Li) detector(98)). However, with the larger shields that come with improved sensitivity, consideration of shield dead times becomes important because extra internal shield events arise from induced radioactivity(85). Some experimenters

c

have therefore used Nal (Tl) shields which, apart from lower cost, have faster pulses(" 101). In polycrystalline form Na(Tl) shields also have improved strength.

Fig. 1.10 Ge(Li) detector used by UCDS. Peterson et al. (1972)(69).iii) HpGe detectors: The cryogenics systems needed to provide permanent deep cooling in Ge(Li) detectors involve considerable cost and continuous development problems. They provide the ultimate limit to the lifetime of such detectors

47

and add weight that could be used for sheilding. The development of the high purity intrinsic coaxial Ge detector (HpGe) certainly alleviated these problems and helped the development of multiple detector arrays of Ge. HpGe requires cooling only during operation and does not involve Li drifting in fabrication. Fig. 1.11 shows the HpGe multi-crystal detector of the NASA-SMM program(85).

Fig. 1.11 Schematic cross section of the Y-ray detector for the NASA-SMM program. Nakano et al. (1976) ( 85).

135*KTHERMAL SHIELD

improvments in the efficiency of Ge detectorswere achieved by widening their intrinsic layers. This

48

tended to increase surface dead regions and hence constrain their use to energies >50-80 keV(6 90), particularly in p-type HpGe. It was really only after the realisation that p-type is particularly susceptible to proton and neutron damage^67 89 102), and Compton scatter background from the dead layer, 'that the more expensive n-type become more acceptable, n-type can also provide a lower low energy threshold. This advantage became particularly important with the discovery of discrete emission lines below 100 keV.

The increased interest in spectroscopy below 100 keV encouraged the development of planar n-type HpGe detectors sensitive to X-rays of energy down to about 10 kev(90 101). Optimisation of the sensitivity of Ge detectors at these energies allows the use of passive collimators rather than active ones because background induced by the collimator is less. Examples of detectors with all these features include the LEGS detector of Paciesas et al. (1983) ( 90), that of Helfand et al. (1981) ( 1 0 1 103), and the I.C. Ge detector (see Ch. 4).iv) CdTe and Hgl2 detectors: Although the development of HpGe has allowed some simplification of Ge detector systems, their sensitivity is still limited in part by restrictions of weight, size, and detector lifetime, imposed by the need for Cryogenic systems. Sensitivity is also ultimately limited by volume dependant background rates (8I+ 1Q1+). The development of room temferature, high Z, low volume semiconductors like CdTe (Z = 48, 52) and Hgl2 (Z = 80, 53)

aims to solve these problems and increase efficiency. However, though Hgl2 detectors have been flown on balloon platforms(105), both CdTe and Hgl2 suffer from problems concerning charge collection, carrier trapping, crystal uniformity, and low hole mobility. Their areas are so far also restricted to -1 cm2(86 106). Energy resolutions of -2 keV at 60 keV have been obtained for Hgl2 ^101t .

1.5 FUTURE PROSPECTS FOR HARD X-RAY INSTRUMENTATIONX-ray astronomy has now gone through its initial stage

of discovery. In the soft X-ray waveband, surveys and imaging of the sky by EINSTEIN and other satellites has been particularly successful. This has certainly encouraged the development of future soft X-ray imaging telescopes such as ROSAT, AXAF and LAMAR^103). Though the current soft X-ray telescopes EXOSAT(107) and TENMA(108), are now largely devoted to spectroscopy, there is a continuing drive to carry out further survey work in the soft X-ray waveband(1 09). For instance, ROSAT will perform an all sky survey below 2 keV. Thus with soft X-ray experiments performing sensitive survey and imaging work, it seems apparent that new hard X-ray instruments should continue to be optimized to provide solutions to specific astronomical problems which can only be addressed at these energies. Tables 1.1 and 1.2 list the important proposed future satellite and balloon borne hard X-ray instruments.

Three areas of development appear to be particularly

50

favoured by hard X-ray astronomers; imaging (using position sensitive detectors), spectroscopy, and time-variability studies(103 110). The somewhat conflicting demands these techniques place on the detector used in an instrument often means the roles played by the major forms of detector (see Sec. 1.4), are somewhat blurred.

Experiment Geometric area (cm2)

OSSE (GRO) 2685GAMMA 1 1400SIGMA 825SAX 1250SALUT 7 800XTE 12000

Table 1.1 Proposed future satellite hard X-ray detectors (10 keV-10 MeV).

Experiment Geometric area (cm2)

LDBF(MSFC) 8100JUPITER 3 340LAPEX 3400MIFRASOIC 5000POKER 11000ZEBRA 5000GRIP 1000DGT 715

Table 1.2 Proposed future balloon hard X-ray detectors (10 keV-lOMeV).

1.5.1 ImagingIn recent years there has been a considerable drive

towards the development of instrumentation capable of producing hard X-ray/low energy Y-ray images of the sky( 103 ill 1 1 2 1 1 3 ) Most involve a coded aperture mask (CAM) placed above a position sensitive detector plane (PSDP). There is certainly some need for such instruments, for instance, to identify Y-ray sources discovered by COS-B, to identify X-ray emitters with sources at other wavelengths, to measure the angular structure of X-ray sources, and to resolve complex source fields.

51

However, other non-astronomical factors have also influenced the trend. For instance, there is a historical precedent that improved angular resolution produces scientific results. This was true in soft X-ray astronomy with the invention of the modulation collimator, and more recently with the success of imaging telescopes( 11 **). A further advantage of imaging is that it can provide an improvement in sensitivity since background can be measured simultaneously with the source. Perhaps as significant is the general development of PSDP's especially in the field of medicine, with the development of Y-cameras (to be used on the SIGMA telescope( 113)) , and in nuclear physics, with the development of Xe proportional counter PSDP's (L15).

An important example of an imaging hard X-ray telescope based around a coded aperture mask is ZEBRA(112 116 117)(see Fig. 1.12 and Table 1.2). The main PSDP for ZEBRA comprises of 2 arrays of Nal(Tl) bars covering the energy range 100 keV-10 MeV and sensitive to =3% Crab at 1 MeV with 1 ms time resolution. Each bar is viewed by 2 photomultipliers, one at each end* The bar width defines the position of an event in one axis, whilst the relative magnitudes of the signals in the 2 photomultipliers is used to provide position information for the other axis.Operation of the bars in anticoincidence with each other provides an effective charged particle and Compton shield.

Mounted above the main detector is a spectroscopic position sensitive proportional counter (SPC) developed from

52

POKER (see Sec. 1.4.1.2). This provides sensitivity for X-rays below =150 kev(117). Position information in the SPC is provided in one direction by knowledge of the cell in which the X-ray was detected, and in the direction along the anodes by the process of current division. A point source produces a unique shadow-gram of the CAM on both detector planes. This can be deconvolved to provide angular information on the source to a few arcmin.

r . TUNGSTEN RA/CCM / ' MASK M ----- /

I NE102 KA5TC — I A C COUNTERS

z!

I P H P 2 N-lN

N.lr r PM TUBES VEWING

ENOS OFHTBRD SCAMCR CRYSTAL

,— POSTON SENSTTTVE / XENON COUNTER

ARRAY OFANTOONCCENCE UMTS C

Fig. 1.12 Schematic view of the ZEBRA y-ray telescope.Dean et al. (1983) (112).

New forms of PSDP are being developed for the future. These include Barium Fluoride and BgO arrays, new Ge arrays and improved GSPC PSDP's^111 112 115). Charge coupled devices (CCD's) are also a possibility(103).

1.5.2 SpectroscopyImaging of the X-ray sky can certainly isolate sources

and help in the understanding of their morphologies. However, it is clear that many fundamental properties of sources can best be investigated using spectroscopy( 10 1). Also, there have been many recent spectroscopic discoveries

53

and predictions in the hard X-ray/low energy Y-ray wavebands, particularly concerning line emissions. Spectral lines have now been observed from about 50 sources (see Ch.2). It is apparent that the study of spectral lines can provide important model independent information on many source parameters(101).

Despite the importance of spectroscopy there are comparatively few future instruments capable of performing sensitive spectroscopic observations, particularly of line features. This is largely due to technological reasons. Ge detectors are the most suitable for the application (see Sec. 1.4.3), but they are expensive if high sensitivity is required. Perhaps for this reason only one major Ge instrument is planned for the future (JUPITER 3), though a few contemporary balloon borne Ge detectors may be reflown. It is likely that semiconductor detectors will be important in the future only if internal background in the crystal and shielding can be reduced(1 0 1). Ch. 4 discusses one possible background effect in the shield of a planner Ge detector system that may be causing problems. A further development might be the use of Hgl2 detectors with Bragg crystal concentrators.

Probably the most important planned instrument to perform sensitive hard X-ray/low energy Y-ray spectroscopy, is the Orientation Scintillation Spectrometer Experiment (OSSE) for the Y-Ray Observatory(118) (see Fig. 1.13).This has four Nal(Tl)-Csl(Tl) phoswich detectors,

54

covering the energy range 0.05-10 MeV, each with a Csl (Tl) anular shield and a Tungsten alloy collimator. OSSE has a continuum sensitivity of 3 x 10"5 ph cm"2s-1, and is -10 times more sensitive to Y-ray lines than previous detectors. The energy resolution of 8% (0.66 MeV) is of course much less than that of solid state detectors. The time resolution is 4 s (spectra) 2 ms (pulsars). The angular resolution is -10'.

Fig. 1.13 Schematic view of an OSSE detector module.Kurfess et al. (1983) ( 1 18).SPC's, such as incorporated on ZEBRA (see Sec. 1.5.1),

are likely to play an important role in spectroscopy below 150 kev(119). In ZEBRA it is primarily the imaging capabilities of the SPC that are paramount. However, this detector incorporates the new technique of "escape gating" which provides improved energy resolution. Ch. 3 discusses this technique and shows that it can also provide excellent background rejection and therefore high line sensitivity.

55

The development of escape gating in GSPC's will also be important. The use of "Penning" gasses in SPC's may improve their energy resolution in the future.

1.5.3 Time-Variability Studies ^The need to study the time-variability of X-ray sources

has been recognized since the early days of X-ray astronomy ( 3 3 7J+) . However, only after the discovery of pulsations from X-ray sources by Kurfess et al.(1971)(12°), and the subsequent observations with SAS-3 and ANS(109), was the importance of such studies highlighted. More recently there has been further interest, inspired largely by new observational results (see Ch. 2).

The design of an instrument capable of performing variability studies is generally not limited by technology. However, the desire for such studies has, in the past, had some influence on instrumentation. It has certainly encouraged the use of balloon borne detectors since these are well suited to fast timing measurements(2 S8). The use of modular scintillation counters with fast electronics(6) has also been encouraged, since a smaller crystal has faster time constants than a large one(121).

The increasing drive for time-variability studies has perhaps exerted an even greater influence on the design of future instruments. The soft X-ray satellites XTE and ASTRO-C(18 3) and the hard X-ray satellite SAX(122) will all emphasize time-variability studies as will the LDBF phoswich

56

detector(123). An extreme example of an instrument designed specifically to perform time-variability studies is FIGARO(12I+) (see Fig. 1.14).

!■ OCTETOW ASS

HAST® OOHC 0.5 an TNeh

1 1 1r n Nal c ICTBCTWO Saw TNekJ

l ~ L ***f,„m.J__1

I:::::!::::

WcmLATtRMJHCLOMB * N«| BLOCKS lernTTiH

I

Fig. 1.14 Schematic representation of the FIGARO telescope.Agnetta et al. (1983) ( 124).

FIGARO is designed specifically to study sources of low energy Y-rays with a periodic temporal signature. The main detector consists of 12 Nal(Tl) crystals shielded on the bottom by 5 blocks of plastic scintillator and on the sides by 14 Nal(Tl) crystals covering the energy range 0.15-6 MeV. The field of view is -tt str. There is no need for fine angular resolution for observations of pulsed X-ray flux since sources can be identified by their periodicity.

57

FIGARO has timing resolution of 25 us and is sensitive to the Crab pulsar in 3 min (5a).

The increased interest in the study of the time-variable characteristics of pulsars has certainly encouraged the development of FIGARO. Ch. 5 contains a discussion of the feasibility of using an instrument with a 2tt aperture to search for and study X-ray pulsars.

1.6 CONCLUSIONIt is apparent from the previous discussions that hard

X-ray instrumentation does not develop solely in response to astronomical needs. Also, however, the development of instrumentation is not solely determined by the technological advances that become available. The instrumentation is in fact determined by both astronomical needs and technological advances.

It seems likely that future instrumentation will be based on large area scintillation detectors, large semiconductor arrays, and spectroscopic proportional counters. Possibly all these will have imaging capabilities by using coded aperture masks or modulation collimators. Future developments such as these should not only account for the astronomical needs at the time, but also of technological advances. Only then can the maximum science be obtained in the shortest time. X-ray astronomy can then continue its rapid advance.

58

CHAPTER 2

OBSERVATIONAL RESULTS IN HARD X-RAY ASTRONOMY

2.1 INTRODUCTIONAlthough most observations of X-ray sources have been

undertaken using soft X-ray telescopes, some important data has been obtained using the hard X-ray detectors discussed in Ch. 1. Presented here is a summary of the hard X-ray characteristics of the major classes of X-ray source, revealed in part by these detectors. Emphasis is placed on topics important in later chapters.

2.2 GALACTIC HARD X-RAY SOURCESGalactic X-ray sources are concentrated towards the

galactic plane and centre. Many are associated with evolved binary systems containing a compact object and can be classified by the nature of this object and its companion star. They often show periodic and/or aperiodic X-ray emission. Those displaying a regular period are usually termed "X-ray pulsars". Most of these are accretion powered X-ray binary pulsars, but some are isolated pulsars.

2.2.1 Massive Neutron Star and Black Hole Binariesi) X-ray emission from massive binaries: X-ray emission can occur in binary systems as a result of matter accreting from a supergiant primary onto a compact companion. Such binaries are concentrated towards the galactic plane and

59

account for -30% of galactic X-ray sources. They usually evolve from systems that originate as two main sequence stars. The more massive of these evolves into the compact object, usually a neutron star but occasionally a black hole(125) (CYGX-1 is a candidate black hole binary). The other becomes the supergiant. Most of the class exhibit pulsed X-ray emission at two characteristic periods, one arises from the orbital motion of the system, the other from the rotation of the compact object. They can often be classified as either "standard" massive binaries(103) (e.g.CENX-3), or Be/X-ray binaries (e.g. A0535+26). (Table 2.1gives some important properties of the best studied accretion powered X-ray pulsars(5) (see also Ch. 5)).

Source Pulseperiod(s)

Orbitalperiod(d)

Log[L (erg s x) J (0.5-60 keV)

SMCX-1 0.71 3.9 38.7HERX-1 1.24 1.7 37.44U0115+63 3.61 24 37.0CENX-3 4.84 2.1 37.9-38.24U1626-67 7.68 0.03 37.1LMCX-4 13.5 1.4 38.8OAO1653-40 38.2 ■? 35.4-36.8GX1+4 115 ? 38.04U1258-61 272 >20 35.84U0900-40 283 9.0 35.9-36.44U1145-61 292 >20 35.04U1538-52 529 3.7 36.64U1223-62 695 41 36.44U0352+30 835 >40 33.6

Table 2.1 Phase averaged properties of accretion poweredX-ray(1983)

sources (5) ) . (adapted from White et al.

The more common "standard" binary system has an 0 type

60

companion star which fills, or nearly fills, its Roche Lobe. The X-ray emission is permanent. In a Be/X-ray binary the companion, is a B-emission star which lies deep within its Roche Lobe. The X-ray emission from these varies from complete absence to occasional large outbursts lasting a few weeks. If the primary star has filled its Roche Lobe then matter will stream through the inner Lagrangian point to form an accretion disc which spirals onto the compact object. These systems usually have orbital periods £4 d, pulse periods -S14 s, and exhibit high spin up. If the primary is enclosed in its Roche Lobe then matter aa^etes onto the

Acompact object by a stellar wind. Systems of this sort have orbital periods 9 d, and erratic pulse periods >102 s.

The X-ray emission in both forms of binary system arises from heating of the accreting matter, powered by the gravitational potential energy released as the material falls. The intense magnetic field of the compact object constrains the accreting plasma to impact on small areas near to the magnetic poles. This creates a "hot spot".X-ray emission from this is believed to be either of fan beam configuration, in which the emission is normal to the magnetic axis, or of a pencil beam configuration, in which two beams of photons follow magnetic field lines from the magnetic poles of the star.

Many X-ray binaries are observed to have two peaks of emission per pulse period. This can result, for instance, if the orbit inclination is =90° and the magnetic axis is

61

perpendicular to the rotational axis. The phase averaged spectra can usually be represented by a flat power law steepening above -20 keV (see Sec. 1.2.1) but line emission has been observed in some sources and is predicted to occur in Black Holes( 1 0 1 118 127).ii) Problems of massive X-ray binaries: Many of the properties and theories of accretion powered binary X-ray pulsars are uncertain. A key problem is deciding whether the compact object is accretion disc fed or stellar wind fed, and how the transferring matter is decelerated. This can be crucial in determining whether the emission is of pencil type or fan beam type(103) . The role of the accreting matter in determining the pulse profile and pulsed spectrum is also unresolved. The characteristics and variety of the pulse profiles, their spectra and time-variability and how these factors are related to the period and period derivative is not yet sufficiently understood to constrain present models^5 l28).

The nature of the spectral lines and spectral breaks observed above 20 keV is also not fully understood. However, they are related to the source temperature, density, bulk motion, isotopic abundancies. Of particular significance are cyclotron lines. The intensity and time-variability of thesere l a t e d to the magnetic field of the neutron star and \s> therefore fundamental to the mechanism of such stars(5).

62

2.2.2 Low Mass Neutron Star Binaries (Galactic Bulge Sources)About 1/4 of galactic X-ray sources show a different

spatial distribution to the massive binaries. They occur wider in galactic latitude and concentrated towards the galactic bulge. They are almost certainly neutron stars in low mass close binaries (e.g. SCOX-1). Many are found in globular clusters where they are probably formed by capture processes. Those in the galactic bulge may have escaped from globular clusters or evolved from cataclysmic variables in which the white dwarf component has collapsed (see Sec. 2.2.3). Their characteristics differ from those of the massive binaries in several respects. TheX-ray photon spectra are usually softer and show no regularperiodicity, and the X-ray to optical luminosity is a factorof -105 lower. Many emit one of two types of X-ray burst;type I last -20 s, occur at intervals of hours to days, and have blackbody spectra. Type II are less common than type I, occur on timescales of 10-100 times shorter, and have spectra resembling that of the source in its quiescent state.

2.2.3 White Dwarf Binaries and Other Unusual StarsLow mass binary systems in which one member is a white

dwarf are believed to be the explanation of some cataclysmic variables and some novae. X-rays are produced by matter accreting from a K,M or N type primary directly onto the degenerate dwarf. The X-ray luminosities are a factor of

63

103 less than for the massive neutron stars and black hole binaries. Two basic classes can be distinguished: i) Those with characteristics like AM HER. These have high magnetic fields, hard thermal bremsstrahlung X-ray emission, and pulsed X-ray emission synchronized with the binary period, ii) Those like DQ HER which have weaker magnetic fields and no observed hard X-ray emission. Some white dwarf binaries, and some cataclysmic variables, do not fit either category.

Other unusual X-ray sources include: transients, which may be massive binary systems or recurrent novae or flare stars and which become luminous for only a few weeks; nondegenerate stars, such as Algol; and non-accreting white dwarf systems, like Sirius B. Spectral lines have now been observed from some of these sources.

2.2.4 Supernova Remnants and Isolated X-ray Pulsarsi) X-ray emission from SNR's and isolated pulsars: X-ray emission can arise from supernova remnants (SNR's) as a result of the heating of ejected material and interstellar gas by the expanding shock wave. Such systems account for a few % of galactic X-ray sources. Though they generally have thermal emission spectra there is a group of SNR's, often termed "Crab-like”, which do not. The best known example of these is the Crab nebula itself, though Vela and MSH15-42 are others (see Table 2.2 and Ch. 5). These systems all contain an isolated radio/y-ray pulsar, probably a rotating neutron star. The stellar remnant from the supernova. The

64

kinetic energy from the acc^/^tion of relativistic electrons by the magnetic field of the pulsar is injected into the SNR(129 13°). This results in a "filled centre" SNR with extended y-ray emission, a smooth synchrotron type spectrum and a flat radio spectrum. In the case of the Crab radiation is emitted over 20 decades of energy and is pulsed at a period of =33 ms with 2 peaks per period. Variable line emission has been observed from the Crab pulsar(10°).

Source SNR log IL„ (erg s”1) J (0.2-4 keV)

Pulsar loq[Lx (erg s“ M ]A (0.2-4 keV)

Pulsar per iod

(s)Crab 37.4 36.0 0.033G29.7-0.3 36.5G21.5-0.9 35.4MSH15-42 35.4 34.6 0.1503C58 35.0G74.9+1.2 34.9Vela X 33.4 32.5 0.089CTB80 32.9

Table 2.2 Properties of Crab-like SNR's and associated pulsars (adapted from Becker (1983) ( 129) and Seward (1983) ( 1 3°)).

ii) Problems of SNR's and isolated pulsars: Many important questions concerning isolated X-ray pulsars and SNR's remain unanswered(103 13°). For instance, many models of these sources are based on the Crab system, but surveys at radio and X-ray energies suggest that objects like the Crab are quite rare and therefore atypical(131). Prior to EINSTEIN the only other known example was the Vela system. However, current birthrate figures for SNR's indicate that there ought to be about 25 "Crab-like" systems in the galaxy^129

65

131 132). Apart from this problem there is a strong need to elucidate the radiation mechanisms of the isolated X-ray pulsars. This could, for instance, help explain the possible turnover observed in the hard X-ray spectrum of Vela and PSR1822-09 ( 1 2t+ 133).

Precise mechanisms to explain the emissions from the Crab itself are unclear. The beam geometry of the pulsar is complex and not compatible with two pencil beams from the poles (133). Different emission regions or mechanisms probably account for the observed differences between the two pulse peaks and between the same peak at different energies. This is also likely for Vela(133 134). The broad single X-ray pulse observed from PSR1510-59 in MSH15-42 is difficult to reconcile with the pulse profiles of the Crab or Vela. More importantly the high dp/dt of PSR1510-59 indicates a characteristic age of -1600 y whereas the SNR appears to have an age of =104 y. Some models can explain this discrepancy(1 35), but Bergh et al. (1984) C136), has suggested that the pulsar is only projected onto the SNR.iii) Recent results in X-ray pulsar astronomy: There have recently been some important observational and theoretical results concerning X-ray pulsars. For example: a 50 ms Crab-like pulsar, designated 0540-693, has recently been discover in the LMC by Seward et al. (1984) ( 1 3 1) ; a 67 mspulsar has been found by reanalysis of HEAO data(137); recent analysis of EINSTEIN data has revealed 7 SNR's with compact objects of which one is thought to be pulsing at 3.5 s(138);

66

from recent observations evidence is accumulating for the existence of fast X-ray oscillations with periods of about 10 ms associated with variable X-ray sources( 137), examples include MSB1728-34 and several others(137); new theoretical predictions suggest that ms rotational periods in low mass X-ray sources should exist(139).

2.3 EXTRAGALACTIC HARD X-RAY SOURCESExtragalactic hard X-ray sources are distributed

uniformly in the sky. Their X-ray emissions do not appear to be of stellar origin. The most intense sources are associated with various forms of active galactic nuclei (AGN). X-rays from these probably arise from inverse Compton scattering of low energy photons off relativistic electrons. Synchrotron processes can account for the photons whilst a central black hole may provide the electrons. The power of the X-ray emission often dominates the rest of the spectrum. Time-variability is observed in some AGN's and spectral lines are predicted(101 118 127). Differences in the various types of AGN may arise from different orientations of otherwise similar systems. Normal spiral galaxies also emit X-rays, however the bulk of this emission is probably due to accreting neutron stars in binary systems.

2.3.1 Seyfert GalaxiesSeyfert galaxies appear to contain a compact, highly

luminous, nucleus from which are observed intense, but variable, optical emission lines. Two types of Seyfert

67

galaxy can be distinguished: Seyfert I galaxies (e.g.NGC4151) emit both broad and narrow optical lines; and Seyfert II galaxies emit only narrow lines. The hard X-ray emission observed from Seyferts seems to be steady and non-thermal.The photon spectra can adequately be fitted by power laws.

2.3.2 BL Lac ObjectsBL Lac Objects occur at the cores of some giant

elliptical galaxies (e.g. MKN501). They display non-thermalcontinuum radiation with only faint optical line emission.Their hard X-ray emissions have steep power law photon spectra .of variable intensity and spectral index on timescales of days. The spectra of some BL Lac objects flattens below 20 keV.

2.3.3 QuasarsActive galaxies known as Quasars are of two general

types; those which emit strongly in the radio and those which emit only quiet radio signals. The former type may be related to BL Lac objects whilst the latter are often indistinguishable from Seyfert I galaxies. Though more than 70 galaxies have been detected at soft X-ray energies only a few have been “observed at hard X-ray energies (e.g. 3C273) .These have hard X-ray photon spectra that can be fitted by power laws. They show large variations in intensity on timescales of months.

2.3.4 Clusters of galaxiesSome clusters of galaxies show extended X-ray emission.

68

This emanates from hot gas contained in the individual galaxies of the cluster, or in the cluster as a whole, or in both (e.g. Virgo). The hard X-ray photon spectra are generally thermal. However, for some clusters a hard X-ray tail is evident, probably arising from active galaxies within the cluster.

2.4 ISOTROPIC X-RAY BACKGROUNDAn isotropic X-ray background has been observed from

soft X-ray energies up to about 100 MeV. There is evidence that =35% of this background may actually arise from distant active galaxies which have not been resolved by detectors. However, part of it may also be due to thermal bremsstrahlung emission from a proposed intergalactic medium.

69

CHAPTER 3

BACKGROUND IN ESCAPE GATED DETECTORS FOR HARD X-RAY ASTRONOMY

3.1 INTRODUCTIONIn X-ray astronomy measurements above 20 keV there is

an increasing interest in sensitive spectroscopy (see Sec.1.5). For example, pulse shape spectroscopy of pulsars isrequired to study the cyclotron line features in sourceslike the Crab nebula and HERX-l(100 lt+0). Such investigationsrequire instrumentation with good energy resolution and high

<

sensitivity. Solid state detectors can provide good energy resolution, but reducing the high background rates in these detectors requires massive anticoincidence shields and large areas of costly semiconductor crystal^85) (see Sec. 1.4.3 and Ch. 4). An alternative instrument which has been shown to achieve good energy resolution (typically 5% at 60 keV), is the escape gated multiwire proportional counter (EGMWPC)(21* 11+1 11+2). Large areas of this form of detector can be readily constructed.

The sensitivity for background limited detection of a source is proportional to the square root of the background rate (see Equ. 1.6). Therefore efficient background rejection is required, particularly in the high background environments in spacecraft or at balloon altitudes. The escape gated detector is inherently capable of efficient rejection of charged particle and Compton events and it is

70

this which determines the good sensitivity of the escape gating technique, which can, of course, also be applied in position sensitive detectors.

Presented in this chapter are measurements made in the laboratory background environment, supported by Monte Carlo calculations, and an assessment of the potential of escape gated detectors for X-ray astronomy. Emphasis is placed on consideration of the spectral shape and origin of the residual background in escape gated detectors. Results are used to make a prediction of the sensitivity that could be achieved by such detectors(143).

3.2 THE ESCAPE GATING TECHNIQUEThe technique of fluorescence escape gating in gas

scintillation proportional counters (GSPC)(ll+l+ 11+5 11+6 lt+7) and multiwire proportional counters (MWPC)(2l+ 11+1 ll+2) for X-ray astronomy applications and X-ray fluorescence analysis(148), has been used for some years. The process of escape gating relies on the high probability that an incident X-ray of energy Ex, greater than the K-edge energy Er of the counter gas, can interact in the gas to produce a K-photoelectron of energy Ep = Ex - ER. This electron is often accompanied by emission of a characteristic fluorescence K-photon (see Sec. 1.2.2). Fig. 3.1 shows the photoelectric mass attenuation coefficient for Xe.

Above the K-edge the coefficient increases since there is then the additional possibility of absorption with the

71

emission of a K-photon. The energy of the K-photon is necessarily just below that of the K-edge energy for the gas so its interaction cross section is small. Thus there is a high probability that the photon will escape from the detection medium. In a conventional proportional counter this results in an "escape peak" in addition to the primary peak hence complicating the spectral response. However, in an escape gated detector an attempt is made at detecting the escaping K-photon separately. A valid count is then deemed to occur only when two pulses are detected simultaneously, and one corresponds to the K-photon energy.

Fig. 3.1 Mass attenuation coefficients for Xe: (a)photoelectric; (b) Compton.

72

The detection technique has two advantages over conventional detection. Firstly, the energy resolution is improved, especially near the K-edge. This is because the energy is obtained from Ex = Ep + ER and since ER is known, only the statistical uncertainty in Ep determines the energy resolution. The second advantage is the reduction in background events. This results from the strict condition that a valid event only occurs if there is a coincidence between pulses from two detection regions with one pulse corresponding to the energy of the K-photon. Thus most charged particles and Compton scattered photon-electron events will be rejected. These events are particularly prominent in space and high altitude environments.

3.3 BACKGROUND IN PROPORTIONAL COUNTERSTo qualitatively relate background measurements made

with proportional counters in the laboratory to those likely to arise at balloon altitudes or in space, an understanding is needed of the background effects in proportional counters in each of these environments. The incident background components in each environment were discussed in Sec.1.2.3. Outlined below, for each environment, are; i) the background components that will be detected by a proportional counter with a 2tt aperture, and ii) the effectiveness of both conventional and escape gated background rejection techniques for reducing the dominant background components. Attention is then given to the

73

effectivness of active anticoincidence shielding.

3.3.1 Background Detected in the LaboratoryThe laboratory background detected in an unshielded gas

counter, without end guards or rise time discrimination (see Sec. 1.4.1.2), results mainly from local Compton scattered photons and y-mesons. The photons contribute only about 50% of the rate since the efficiency for photon detection in the low density medium is small and rapidly drops off with energy, (see Sec. 1.4.1.1). Also, as a consequence of the low density of the gas, y-mesons, likely to be minimum ionising, can deposit energy in the range of the detector, (around 10 keV in a 1 atm. Xenon counter, 1 cm deep). Thus the background in a detector will peak at 10-20 keV depending on its typical dimensions and the gas pressure. A long high energy tail will, result from the combined effects of y-mesons off the minimum ionizing plateau or passing through the gas at small angles to the anode, and high energy photons detected with diminishing efficiency. Radioactivity from instrument materials and contaminants in the gas, particularly Tritium, can provide further counts (**6 1 **9) .

3.3.2 Background Detected at Balloon AltitudesFor a proportional counter at balloon altitudes (>40 km)

the background detection below 100 keV is dominated by atmospheric photons through the aperture^6 15 2l+) . Above about 60 keV photoelectric absorption of photons directly

74

penetrating the detector wall becomes important( 15). Lower energy counts can result from Compton scattering of atmospheric y-rays in the gas or wall, or from extended photoelectron tracks ( 6 8 15 2 4 150 151). Cosmic rays and accompanying secondaries will deposit energy along extended tracks in the gas or interact in surrounding materials to produce further Y-rays and additional Compton events (21+ ****) . They will also produce small numbers of radioactive nuclei, spallation products and neutrons(12 17 19).

3.3.3 Background Detected in SpacecraftFor a proportional counter in a spacecraft the

principal background sources result from energy deposited by minimum ionizing cosmic rays and their secondaries, Y-rays and electrons. For low altitude orbits cosmic rays account for approximately 20% of the flux above about 1 keV, whilst their secondaries in the upper atmosphere contribute 20-30% by way of earth albedo and inner radiation belt particles(17). A Shuttle orbit (typically 58° inclination) has cosmic ray contributions of about 50%. Y-rays will interact in the detector wall or gas to produce a spectrum of Compton electrons. These may be minimum ionising and so deposit energy in the range of the detector along extended tracks, or be of low energy thus depositing similar energies but along short tracks (electrons with energy between about 0.1 and 100 MeV deposit energy above the detector threshold). At low energies some 20% of the X-ray

75

background is from Compton electrons, the rest largely due to cosmic rayst1*1*). In large spacecraft such as Shuttle or Space Station, the increased flux of induced Y-rays will result in more Compton electrons.

3.3.4 Background in Actively Shielded MWPC'sThe summary above suggests that many of the background

events in proportional counters arise from charged particles passing through the detector. Active shielding such as inter-layer anticoincidence and rise time discrimination is good at reducing these events (see Sec. 1.4.1). However, in an escape gated system the stringent requisite for a valid event means that rejection of charged particles is much more certain. For instance, the escape gated mode will reduce the occurrence of counts resulting from charged particles, which deposit energy above the lower threshold in only one cell. A conventional counter will count this type of event because energy must be deposited in at least two cells or along an extended track for an event to be rejected.

Other residual background in conventional counters arises from the results of Compton interactions. Compton events in the gas in which either the scattered photon or electron is lost or leaves a short track, will not be vetoed by a conventional system^152). Compton events in the walls can emit into the gas, photons, short track electrons, or electrons parallel to the anodes. These are also unlikely to be rejected. An escape gated system, however, is more

76

likely to reject such events because of the requirement of coincidence pulses from two cells. This condition will also help to reduce counts due to fluorescence or bremsstrahlung photons, radioactive decay photons, Auger electrons and electron-positron pairs, from the walls or gas.

The occurrence of charged particle and Compton type events is greater at balloon altitudes and in space than in the laboratory. Thus any advantage that an escape gated mode has in background rejection as measured at ground level, is likely to be increased in these environments.

3.4 THE MWPCTests were performed with the MWPC shown in Fig.

3.2(ll+9). The counter is mounted in a cylindrical vacuum box with a carbon fibre entrance window and is filled with 95% Xe and 5% CO2 at a pressure of 0.97 atm. The electrode design provides for 6 sided anticoincidence (see Sec.1.4.1.2) and 6 detection layers, each 2 cm deep. Each layer has 20 anode wires separated from neighbouring layers by cathode planes. The outer wires in each layer are used to form 2 side guards 2 cm deep. The cathode signal technique is used to form 2 end veto sections also 2 cm deep. A 7th layer forms a rear guard cell whilst a front guard cell is formed by a layer of cathode wires 2 mm from the front window. Each layer has its own amplifier, pulse height selector, and rise time discriminator (RTD) ( 1 5 3). Each guard has its own amplifier and discriminator with

thresholds set just above noise. The layer can be operated in two halves since alternate anode wires are connected.

77

GUARD

ANODE CATHODE

Fig. 3.2 The MWPC: (a) plan view (with window removed); (b)cross section.

Experiments undertaken with the counter operated conventionally (normal mode) used a computer controlled logic system to process the signals. This selects for pulse height analysis, any events from a chosen layer, that are of correct pulse height and are in anticoincidence with signals from all guards. Spectra or count rates from any layer can

78

then be summed. The low gas pressure of this counter results in long ionisation tracks for X-rays >30 keV, which makes RTD very inefficient at high energies. Hence, for maximum efficiency in normal mode, measurements were performed without RTD. Anode voltages were fixed at =1650 volts to give best energy resolution at 60 keV.

Experiments in escape gated mode were performed using pulse height analysis of signals from each of the 6 layers in turn, in coincidence with events corresponding to the K-photon for Xe (29.8 keV) from each other layer in turn.The rates were then summed. Events with energies within a window of ±AE keV about 29.8 keV were assumed to be K-photons. AE was adjusted to be as narrow as possible whilst consistent with achieving good efficiency (typically AE = 5 to 10 keV). RTD was performed on the K-photons but not on the signal events. No end guard anticoincidence was used.

3.4.1 Efficiency and Energy Resolution of the MWPCThe detection efficiency for one layer of a MWPC for

X-rays of energy E incident normally, when operated in normal mode, is given by Equ. 1.7 and 1.8. In escape mode the efficiency for one layer is given by:

e(E) = x(E)Ta (E)Rb(E)pKwKTc (K)Rd (K) (3.1)

where t (E) is the transmission of the window material, whichfor events of interest here is near 1. T (E) = EXP(-y„x )a E arepresents the probability that the X-ray reaches a layer

79

distance x into the counter where y_, in this case, is the mass absorption coefficient of the gas (i.e. the mass attenuation coefficient * density (see Sec. 1.2.2)). R^E) = 1 - EXP(—^exfc>) represents the probability that the photon is absorbed in the layer of depth x . pK is the probability that the interaction results in ejection of a photoelectron and wK is the fluorescence yield of the K-shell. Tc(K) = EXP(-ukxc) represents the transparency of the layer to the K-photon where is the mass absorption coefficient of the gas for K-photons, and xc is the path length of the X-ray in the layer of origin. R^(K) = 1 - EXP(-uKx ) is the probability that the K-photon is absorbed elsewhere in the gas, where x^ is the path length therein.

It is not possible to solve Equ. 3.1 analytically because parameters xc and x^ follow a statistical distribution. The efficiency can, however, be obtained using Monte Carlo computer programs to simulate the detection of X-rays in the MWPC. Such programs were developed and Fig. 3.3 shows a plot of efficiencies calculated with them for both escape mode (EGMWPC) and normal mode (MWPC) operation of 4 layers of the counter.Also shown are measured efficiencies taken with X-ray sources evenly illuminating the whole detector window. Comparison is made in escape mode with different escape gating energy window widths AE. It can be seen that operation with a window of ±5 kev is only slightly less efficient than ae = ±10 keV.

80

df>

>4uzCdMOMCuEuCi3

Fig.of the MWPC: (a) normal mode? (b) escape mode AE =±10 KeV; (c) escape mode AE = ±5 keV.

Table 3.1 also gives results for efficienciescalculated for both escape mode and normal mode and measured values for the energy resolution of the detector in the two modes of operation.

81

E n e r g y(keV)

M o d e of O p e r a t i o n N u m b e r of M W P C

l a y e r s

M e a s u r e dE f f i c i e n c y

(%)

M o n t e - c a r l oE f f i c i e n c y

(%)

E n e r g yR e s o l u t i o n

(%)

4 4 . 2 E G M W P C 4 3 . 2 ± 0.5 3.84 4 . 2 E G M W P C 6 - 7.0 -

4 4 . 2 M W P C 4 1 2 . 0± 2.0 1 1 . 2 1 2 . 2± 0 . 24 4 . 2 M W P C 6 - 14. 5 -

4 4 . 2 E G H D (latm) - 3 . 0 ± 0.4 5.0 4. 5 ± 0 .54 4 . 2 E G H D (2atm) - - - -

4 4 . 2 N M H y b r i d (latm) - 4 . 5 ± 0.4 - -4 4 . 2 NM H y b r i d (2atm) - - - -5 9 . 5 E G M W P C 4 3 . 2 ± 0.4 3.9 -

5 9 . 5 E G M W P C 6 5.6 ± 0.6 6.2 5.1 ± 0 . 25 9 . 5 M W P C 4 6 . 9 ± 0.5 6.2 8 . 2 ± 0 . 25 9 . 5 M W P C 6 1 0 . 0± 1.5 8.0 -

5 9 . 5 E G H D (latm) - 2 . 2 ± 0.4 3.0 4. 9 ± 0. 15 9 . 5 E G H D (2atm) - 4 . 2 ± 0.5 - 5.0 ± 0. 55 9 . 5 N M H y b r i d (latm) - 2.8 ± 0.4 - 8 . 2 ± 0 . 259. 5 N M H y b r i d (2atm) - 6. 3 ± 0.6 - 1 0 . 0± 0. 58 0 . 0 E G H D (latm) - 0 . 6 5 ± 0 . 0 5 - 5. 9 ± 0.48 0 . 0 E G H D (2atm) - - - -8 0 . 0 N M H y b r i d (latm) - 0 . 9 5 ± 0 . 0 5 - 8 . 8 ± 0.480. 0 NM H y b r i d (2atm) - - - -

1 2 2 . 0 E G H D (latm) - 0 . 0 8 ± 0 .04 - 6 . 5 ± 0. 51 2 2 . 0 E G H D (2atm) - 0 . 3 5 ± 0 . 0 5 - 9 . 0 ± 0. 51 2 2 . 0 N M H y b r i d (latm) - 0 . 1 4 ± 0 . 0 4 - 7.1 ± 0.51 2 2 . 0 N M H y b r i d (2atm) — 0 . 6 0 ±0. 0 5

"1 0 . 0± 2.0

Table 3.1 Efficiency and energy resolution of the multiwire proportional counter operated in normal mode (MWPC) and escape mode (EGMWPC), and of the hybrid detector in normal mode (NM Hybrid) and escape mode (EGHD).

3.4.2 Background in the MWPCBackground spectra were taken in the laboratory without

a collimator or additional passive shielding, though the steel vacuum box has a thickness equivalent to a few mm Lead. Fig. 3.4 compares the normal mode and escape mode background rates obtained using the techniques described in Sec. 3.4.1. The general form of the normal mode background is explained well by the predictions of Sec. 3.3. The peak at about 30 keV results from the detection of the Xe K-photons. The escape gated background spectrum shows a high energy tail-off similar to that of the normal mode spectrum, however the count rate is lower by a factor of

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about 15.

MWPC: (a) in normal mode; (b) in escape mode.The effectiveness of the various background rejection

techniques is illustrated in Fig. 3.5. This shows background spectra for the top layer of the detector. The reduction in normal mode background when anticoincidence endguards and interlayer anticoincidence is switched on, is consistent with rejection of y-mesons and a slight reduction in efficiency (7% at 60 keV). The remaining count rate results from photons entering through the window.

The escape gated background with RTD shown in Fig. 3.5 is less than the normal mode background with anticoincidence. The reduction can be explained by a lower efficiency for the detection of aperture photons (at 60 keV

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this efficiency is down by a factor of 3), and more efficient rejection of y-meson and Compton type background. From Fig. 3.5 it can be seen that with this system the effect of RTD on the escape gating events is important. It is thought likely that the high escape gated count rate without RTD results from an inefficiency of the the single channel discriminator at stopping events with long

_________■ ■ i — i— ■— I i i i . —0 10 20 30 40 50 60 70 80 90 100 110 120ENERGY (keV)

Fig. 3.5 Background in layer 1 (top layer) of the MWPC:(a) operated alone; (b) in anticoincidence with layers 2,3 and 4; (c) in anticoincidence withlayers 2,3,4 and endguards; (d) escape gated by layers 2,3 and 4 with no RTD on these layers? (e) escape gated by layers 2,3 and 4 with RTD.

Fig. 3.4 shows a peak in the escape gated background at about 60 keV which is present in all escape gated background spectra. It is explained as follows; consider the simple escape gating system consisting of two neighbouring layers of Xe gas, layer A and layer B (see Fig. 3.6).

60 ± AE keV PHOTONS

Fig. 3.6 Detection of 60 ± AE keV photons in two elements,A and B, of an escape gated detector. The two processes leading to detection are shown.

Analogue signals from layer A are accepted only if they are in coincidence with events of about 30 ± AE keV detected in layer B (the k-photon energy is about 30 keV). The width, ±AE, set for the window on the single channel analyse (SCA) is dependant on the energy resolution of the layer. Choice of a suitable width is critical to the escape gating sensitivity since if it is too narrow efficiency is lost, while if too wide the background increases. A photon of energy 60 ± AE keV can be detected by this system in two ways. It can interact in layer A producing a 30 ± AE keV Xe K-photon which is detected in layer B. By adding 30 keV to the energy of the event in the first layer we obtain

the 60 keV result in the usual way. In the second situation, the incident X-ray passes through layer A to interact in layer B to produce a 30 ± AE keV photoelectron and a 30 keV K-photon. The photoelectron has energy within the range of the gating window and its detection is indistinguishable from a 30 keV K-photon. If the K-photon is now detected in layer A, then the result of adding 30 keV to the energy deposited in layer A is again 60 keV. The second process is not negligible and the resulting count at 60 keV can occur for any incident photon in the range 60 ±AE keV interacting by the second method. Events outside this range can only be detected by the first method. Thus there is a peak in the escape gated background at 60 keV and, of course, also in the detection efficiency for signal photons around this energy. The peak can never be avoided, even if the two layers were to contain detection media with different K-escape energies there would still be a peak at the sum of the two energies. While the sensitivity may not be greatly affected by the peak it could cause problems with spectrum unfolding. Genuine spectral features around this energy could become confused if the spectral response is not accurately known.

The background rate of layer 4 escape gated by all other layers increases by about 22% at 60 keV when the layer is operated as two halves each escape gating each other.This is associated with a 24% increase in efficiency.

86

Fig. 3.7 displays the effect on the escape gated background in layer 1 when it is operated in anticoincidence with the 6 side guards. The background is reduced by a factor about 2 at low energies without loss of efficiency. This may be due to reduction of Compton events from the walls depositing energy in 2 layers.

lOIo

I><D

IWegI6ocaE-i<BSzZOCJ ENERGY (keV)

Fig. 3.7 Background in layer 1 (top layer) of the MWPC escape gated by layer 2: (a) withoutanticoincidence with endguards; (b) in anticoincidence with endguards.

3.5 THE ESCAPE GATED HYBRID DETECTORAlthough the relative transparency of Xe to its

K-photon enables this photon to escape from the MWPC layer in which it was produced, it also makes difficult the subsequent detection of the photon in other layers. High efficiency can therefore only be achieved with a deep counter. An attempt to overcome this problem is provided by the escape gated hybrid detector (EGHD). This system has a

87

low density Xe proportional counter above a scintillation detector. The scintillator is used to detect the K-photons from the proportional counter. The signals from the scintillator are then used to escape gate pulses from the proportional counter. The scintillator will have near 100% efficiency for detecting the 30 keV photons entering it. A detector of this sort also has the advantage that high energy photons, for which the Xe is fairly transparent, can pass through to the scintillation counter. Thus the energy range is extended, though with lower energy resolution for the scintillation events.

To assess the performance of the hybrid system, experiments were made with cylindrical single anode Xe proportional counters escape gated by a LAPAD phoswich scintillator (see Sec. 1.4.2 and Fig. 1.7). Fig. 3.8 shows the geometry used. This was technically easy to achieve but by no means represents an ideal configuration. Monte Carlo calculations were used to extrapolate the measurements to other possible geometries. RTD was used on the phoswich to determine in which scintillator an event occurs and hence reject non-X-ray events ( 15**) . A single channel analyser was used to select K-photon events. The proportional counter had no RTD or endguards. Experiments were performed on the escape gated system to assess the effect of passive shielding and different gas pressures.

88

1.9mm LEAD

SHIELD S

n11II ALUMINIUM HOUSING/

ALUMINIUMWINDOW\

PHOTOMULTIPLIER

T “ Csl :|14mm | •"Nal 2mm

BERILIUMWINDOW XPROP-COUNTER

Xe 24mm

c

65mm

COLLIMATOR

Tzr

1.5° FWHM

PHOSWICH

Fig. 3.8 The escape gated hybrid detector (EGHDl

3.5.1 Efficiency and Energy Resolution of the Hybrid Detector

The detection efficiency for X-rays of energy E incident normal to an EGHD is given by:

ME) = Rb (E)pKwKTc (K)sK (3.2)

The transmission of the window materials is taken to be unity. Rb (E) = 1 “ EXP(-uExb) represents the probability that the X-ray is absorbed in the proportional counter gas of depth xb, iiE is the mass absorption coefficient. pR and wR have the same meaning as in Equ. 3.1. TC(K) =EXP(-yKxc) represents the transparency of the gas to the K-photon where x is the path length for K-photons and \iv isC i\

89

the mass absorption coefficient of the gas for K-photons. sR is a geometric factor to account for the proportion of K-photons leaving the counter which then pass into the scintillator. The efficiency of the scintillator is taken to be 1. Table 3.1 displays efficiencies of the EGHD, using 1 and 2 atm. of Xe, measured with collimated sources, and results of Monte Carlo efficiency calculations. Table 3.1 also shows results for the energy resolutions of the EGHD at both pressures. These compare favourably with those for the EGMWPC.

Fig. 3.9 shows a plot of the calculated efficiency for a 1 atm. hybrid system compared to values measured with the geometry of Fig. 3.8.

<#P

O2WM<JMCmCnca

Fig. ENERGY (keV)3.9 Calculated and measured efficiencies of the hybrid

detector.The calculated efficiencies of the escape gated hybrid

detector are in broad agreement with the measured values and are similar to those for the escape gated MWPC (EGMWPC).

90

The geometry used for the Monte Carlo calculations of the hybrid detector involved a layer of Nal and a layer of Xe of the same area as the MWPC layers. This geometry is more efficient in escape mode than that of Fig, 3.8. This explains the slightly higher calculated efficiencies in Fig. 3.9 compared to the measured values.

3.5.2 Background in the Hybrid DetectorLaboratory background spectra were taken with the EGHD

in both normal mode (the proportional counter operated alone, (NM Hybrid)) and escape mode (the proportional counter escape gated by the phoswich, (EGHD)). Fig. 3.10 compares the background spectra for 1 and 2 atm. gas pressures taken in these modes using the geometry shown in Fig. 3.8 but with the passive shielding removed.

The change in spectral shape and the overall increase in background rate observed in normal mode when the pressure is increased from 1 to 2 atm., can be interpreted as arising from 2 factors. Firstly, an increased efficiency for the detection of background photons and secondly the increased energy deposited in the detector by y-mesons. The 30 keV K-fluorescence peak is still present. The increase in escape gated background with pressure is due mainly to a higher photon detection efficiency. The peak at 60 keV is due to the same effect as described in Sec. 3.4.2. The peak is in the same position as that of the EGMWPC system since the K-photon for Iodine has nearly the same energy as that

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for Xe.

Fig. 3.10 Background in the hybrid detector: (a) in normalmode with 2 atm. Xe; (b) in escape mode with 2 atm. Xe; (c) in normal mode with 1 atm. Xe; (d) in escape mode with 1 atm. Xe.

The effect of passive shielding on the background rate in the hybrid in both normal mode and escape mode is illustrated in Fig. 3.11. It can be seen that increasing the thickness of passive shielding, especially benefits the escape gating technique. This agrees with the hypothesis that most of the escape gated background counts result from background Y-rays after most of the y-meson and Compton event background has been eliminated. In the normal mode proportional counter only about 50% of the background is due to photons the rest resulting from y-mesons. The effect of passive shielding is to reduce the background that is due to

92

photons and hence reduce the escape gated background proportionally more than the normal mode background. Any increase in the flux of Compton electrons caused by the shielding is also efficiently eliminated by the escape gated system. Thus passive shielding plays an important role in escape gated detectors.

passive shielding: (a) MWPC in escape mode; (b)hybrid detector in normal mode with no shielding;(c) hybrid detector in normal mode with 1.9 mm of Lead shielding; (d) hybrid detector in escape mode with no shielding; (e) hybrid detector in escape mode with 1.9 mm of Lead shielding.

The effect on the background rates of the presence of a nearby 60Co source was also studied. High energy photons from this source will undergo multiple Compton scattering in the surroundings and increase the flux of low energy photons to a level more like that encountered at balloon altitude (8 **5) *

93

Since the escape gated system has a smaller M-meson and Compton background contribution than the normal mode system, the 60Co tends to increase the background in escape mode proportionally more than in normal mode. The overall change in rejection efficiency is also influenced by the better rejection of Compton electrons that is inherent in the escape gated technique. There is still efficient rejection of Compton events in the EGHD though there are more of them in this detector than in the EGMWPC, due to the high density of scintillator present.

Comparison with the EGMWPC shows that the EGHD background spectrum is similar (see Fig. 3.11).Differences at lower energies can be atributed to differences in efficiency for the rejection of u-mesons.- A phoswich gating layer is likely to be less susceptible than Xe to low energy deposits resulting from y-mesons. However, the lack of RTD and endguards in the cylindrical proportional counter of the EGHD, in fact makes this system relatively more susceptible than the EGMWPC. The low energy background of the MWPC normal mode system is greater than that of the hybrid system in normal mode because of the larger dimensions of the gas layers(17).

3.6 SENSITIVITY OF ESCAPE GATED DETECTORSUsing Equ. 3.1 and the results given in Sec. 3.4 and

3.5 it is possible to estimate the relative sensitivities of the escape gated systems in the laboratory background, and

94

to predict the sensitivity of a hypothetical full size balloon borne EGMWPC instrument for hard X-ray astronomy. Table 3.2 compares the sensitivities in the laboratory, at approximately 44 and 60 keV, of the EGMWPC and EGHD.

Energy(keV)

Mode of Operation Relative Sensitivity (B17 2E“ 1)

44.2 EGMWPC 2.2 x 10“344.2 WMPC 2.9 x 10"344.2 EGHD (latm) 2.6 x 10“359.5 EGMWPC 3.7 x 10“ 359.5 MWPC 4.7 x IQ”359.5 EGHD (latm) m1Or—1X00•m

Table 3.2 Relative sensitivities of the various detectors.

The efficiency data used to calculate the results in Table 3.2 is from Table 3.1 whilst background data used for the EGMWPC is from Fig. 3.4. Background data for the EGHD was obtained using the geometry of Fig. 3.8 but with the collimator removed in order to make a fairer sensitivity comparison with the EGMWPC. The lead shielding provides efficient attenuation for X-rays <200 keV, similar to that of the vacuum box of the MWPC. Table 3.2 suggests that the sensitivities of the laboratory systems are broadly similar. However, none of these systems represents an optimum design. The results given in Sec. 3.4 and 3.5 indicate that in all cases improvements in the background rejection capabilities can be made. This is especially true for the EGMWPC. The effect on the background and efficiency of operating a mid-layer in two halves and using endguard anticoincidence,

95

indicate a potential increase in sensitivity below about 60 kev, of perhaps 50% (see Sec. 3.4.2). However, this increase in sensitivity is small compared to that achievable by optimizing the gas pressure.

It can be seen from Equ. 3.1 that the EGMWPC efficiency is a complex function of gas pressure, incident photon energy and cell dimensions. Optimizing these parameters, and especially that of gas pressure, is very important. Monte Carlo calculations suggest that for optimum efficiency at 60 keV for this EGMWPC, a gas pressure of about 6 atm. is needed. This pressure results in an increase in efficiency by a factor of 5 at this energy (a similar optimization is possible for the hybrid system). Results given in Sec. 3.5.2 suggest that this increase in gas pressure will also raise the background rate due to the increased efficiency of detection of background p h o t o n s 6). However, this can be alleviated by passive shielding. The effect of increasing the gas pressure will have an additional advantage of shortening the ionization tracks resulting from X-ray events. This will increase the efficiency of rejection by RTD of background events leaving long ionization tracks since these will have rise times more easily distinguishable from those of X-ray events (see Sec. 3.4). This will be especially true for high energy Compton scattered electrons(1 52).

Using the Monte Carlo efficiency predictions, the laboratory background data, and the background rates

96

measured at altitude for normal mode operation of the balloon borne MWPC POKEr (21+) (see Sec. 1.4.1.2), the 1 sigma sensitivity of a hypothetical 1 m2 balloon borne EGMWPC detector can be estimated. Table 3.3 shows the values obtained. The results assume cell dimensions as for the MWPC shown in Fig. 3.2, a 1.5° passive collimator, and a 10^ s observation at 2 mbar. The gas pressure is taken to be 6 atm. Factors to account for the reduction in background due to operation of endguards, of layer subdivision, and the effect of the collimator were included.

Energy(keV)

EGMWPC minimum defcetectable flux (ph cm“2s“ 1keV”1)

Flux from Crab (ph cm”2s” eV” 1)

Flux from HERX-1 (ph cm” 2s” eV" 1)

44.259.5

3.0 x 10“6 3.7 x 10“6

4 x 10“31 x 10"31 x 10”43 x i c r 5

Table 3.3 Estimated minimum detectable flux for a hypothetical1 m2 EGMWPC compared with the total flux from the Crab and HERX-1.

The detector sensitivities shown in Table 3.3 compare very favourably with those of the best state of the art proportional counter and scintillator instruments( 1 55), although improvements could still be made. The cell sizes for example were not optimised for this application.Although the prediction is based on laboratory background measurements, the conclusions of Sec. 3.3 and 3.4 would suggest that the escape gated background at altitude is unlikely to be much greater than in the laboratory, due to

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the high efficiency of charged particle and Compton background rejection. Experiments at balloon altitudes (using a modified ARIEL VI MWPC) are planned to verify this assumption. An optimized hybrid system is found to provide no real benefits over the EGMWPC, especially when one considers the complexity and cost of fabricating a full size system.

3.7 CONCLUSIONThe escape gated multiwire proportional counter

(EGMWPC) and the escape gated hybrid detector (EGHD) not only have good energy resolutions, typically 5% at 60 keV, but are also shown to have excellent background rejection efficiency, and therefore potentially high sensitivity, though this is limited to energies above the K-edge of the gas used. A hypothetical 1 m2 EGMWPC is predicted to have a sensitivity of about 3 * 10“6 ph cm_2s”1keV” 1 at 44 keV for a 10^ s observation. The technique can therefore provide a significant advance in the field of hard X-ray spectroscopy from balloons or in space. However, the special ability of the technique for rejecting Compton type events as well as charged particles, makes the escape gated detector particularly suitable for use in high background environments. With the increasing use in the future of large space craft such as Shuttle and Space Station, instruments with this type of background rejection capability will become more necessary. The escape gated

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detector seems to provide a solution at relatively low cost, and has the additional advantage that an EGMWPC system could be adapted for imaging applications.

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CHAPTER 4

HIGH ENERGY LOSS EVENTS IN THE ACTIVE ANTICOINCIDENCE SHIELD OF THE IMPERIAL COLLEGE GERMANIUM DETECTOR

4.1 INTRODUCTIONDiscussed in Ch. 3 was a background rejection

technique which will help to improve the sensitivity of proportional counters to spectral lines. This is needed to satisfy an increasing desire for improvements in the sensitivity of spectroscopic studies of hard X-ray sources (see Sec. 1.5). Though the energy resolution of proportional counters is adequate for some of the objectives of spectroscopy, many such studies would benefit from the unparalleled energy resolution that can be provided by solid state detectors, especially Germanium detectors (see Sec. 1.4.3.2). Thus there is a drive to improve the line sensitivity of these detectors as well as proportional counters.

As part of the drive there is a continuing need for a more complete understanding of the detection characteristics of the active anticoincidence shielding used with Germanium detectors. Such improved knowledge can help in work to improve the efficiency of operation of the shields.Efficient shielding is important if sensitive Germanium detector systems are to be developed at reasonable cost.This is because semiconductor crystals are particularly susceptible to certain background components, and because increasing detector sensitivity by increasing Germanium

crystal area is costly since large areas of crystal are expensive to produce and require expensive cryogenics (see Sec. 1.4.3).

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One problem with active anticoincidence shields has been causing particular concern. It has been observed by several workers using large (>15 kg) Nal(Tl) shields(66 1 56 1 57 ) The effect results from events almost certainly due to the passage of heavy cosmic ray nuclei. These nuclei can lose large amounts of energy in a shield, causing light output greater than the normal detection level by orders of magnitude. Some events appear to have a noisy tail with a decay time constant much longer than is normal for Nal.This can result in extra periods when the shield is producing veto signals. The electronic design largely determines the degree of this additional shield "dead time".

The purpose of the experiment descibed in this chapter was to assess the importance of the high energy loss events using the Imperial College Germanium detector, by determining their frequency and size spectrum. This knowledge could be useful for optimizing the electronics design to deal with them in the best way. A flight of the Germanium detector, with its 15 kg Nal(Tl) shield, on a high altitude balloon in Canada in Aug. 1983, provided a good opportunity for investigating these high energy loss events.

4.2 Nal(Tl) ACTIVE ANTICOINCIDENCE SHIELDING FOR GERMANIUM DETECTORSOutlined in Sec. 1.4.2.1 were the characteristics of

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inorganic scintillation counters important for their use as hard X-ray detectors. Outlined in this section are the important characteristics of the Nal(Tl) scintillation counters used for shielding Germanium detectors in general, and for shielding the Imperial College Germanium detector in particular.

4.2.1 The Scintillation Efficiency of Nal(Tl) to Different Radiations

An ionizing particle incident in a Nal(Tl) crystal will give rise to scintillation photons. These can be detected and amplified by a photomultiplier to produce an electrical signal of amplitude proportional to the number of photons produced (see Sec. 1.4.2.1). The efficiency with which the scintillation light is produced, dL/dE, is determined by the nature of the incident particle and the specific fluorescence dL/dr. L is the scintillation light output, E is the energy loss of the particle or photon and r is the distance it has, traversed in the scintillator. For Y-rays dL/dE is nearly constant with energy above about 150 keV. Below this energy the light output is only approximately proportional to the amount of energy released in the crystal (see Fig. 4.1). Below about 100 keV the dL/dE vs.E characteristic for Y-rays is closely related to that for electrons because at these energies the scintillation pulses due to Y-rays can be atributed to photoelectrons in the crystal arising from photoelectric interactions( 1 58).

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Fig. 4.1 dL/dE vs. E for Y-rays in a 1 cm diameter x 1 cm high Nal(Tl) crystal. Kaiser et al. (1962) ( 158).

The scintillation efficiency for protons is approximately constant with energy but drops slightly at energies below a few MeV. The scintillation efficiency is greater than for Y-rays by a factor of about 40% at 5 MeV(61+). For alpha particles of energy below about 20 MeV dL/dE increases with increasing energy (see Fig. 4.2). The response for heavier particles (atomic mass 10-20) is similar but is linear above about 6 MeV nucleon” 1. In both Csl (Tl) and Nal(Tl) the scintillation pulse heights for a given particle energy decrease with increasing particle mass.

concentrations to protons and alpha particles. Birks et al. (1964) (6l+)w

Fig.

103

The time structure for the main component of the scintillation emission intensity in inorganic crystals, is given by:

nt (ph s"1) = N [ EXP (—t/T) - EXP(-t/tR)] (4.1)

For Nal(Tl) t = 230 ns at room temperature. tr - 60 ns for Y-rays but only a few ns for alpha par tides (6 4 157).

4.2.2 The Imperial College Germanium Detector and ShieldThe astrophysics group at Imperial College has

developed a prototype fine energy resolution Germanium detector system for balloon borne astronomical observations at hard X-ray energies in the range 20-500 keV. It is based on a PGT model IGP 1510 (159) detector consisting of a 46 cm2 array of 3 equal area 1 cm thick planar intrinsic Ge crystals cooled by a common dewar. Fig. 4.3 shows the system as configured for flight. The energy resolution obtainable is 0.82 keV at 122 keV.

The anticoincidence shield for the instrument consists of a 4 cm thick 15 kg Nal(Tl) crystal in 4 segments viewed by 10 2.5 cm photomultiplier tubes. This reduces background due to charged particles and atmospheric and cosmic X-rays arriving outside the acceptance angle by vetoing events detected in the shield in coincidence with events in any of the 3 Ge crystals. The acceptance angle is defined by 3 Tantalum honeycomb collimators of 5° FWHM, one above each Ge crystal. The overall 3a sensitivity at 50 keV for a 3 h

104

balloon flight with a 2 keV wide energy channel is estimated to be -2.1 x 10-1+ cm-2s_LkeV” L (assuming 50% atmospheric attenuation).

Fig. 4.3 The Imperial College balloon borne Ge spectrometer

105

4.2.3 Nal(Tl) Shield Dead Time due to Cosmic RaysPeriods during which an active Nal(Tl) anticoincidence

shield is ineffective at vetoing background events, i.e. periods when a veto signal is being produced (the shield "dead time"), can occur as a result of high background rates caused by the interactions of cosmic rays and their secondaries in the shield. Such shield dead time can significantly reduce the useful detection time of a solid state detector during an observation. The dead times can occur because large energy deposits cause enough light output to saturate the photomultipliers or electronics.During the dead times valid X-ray events in the Ge crystal are either not counted or are counted without discrimination from background, depending on the electronic configuration(85).

The degree of shield dead time is dependant on the numbers and energy distribution of cosmic rays since this determines the size and duration of the fluorescence pulses produced in the shield (see Sec. 4.2.1). This then effects the choice of anticoincidence gate time. For large shields the dead times can be >10%(9°) and even approach 100%. The composition and fluxes of cosmic rays at balloon altitude has been well studied, thus, neglecting saturation effects, some prediction can be made of the resulting possible dead times in a shield.

The cosmic ray flux at the top of the atmosphere is composed mainly of energetic protons, together with Helium, Carbon, Oxygen and Iron nuclei, though all nuclei are

106

present(15 0 161) (see Sec. 1.2.3 and Fig. 4.4) .

Fig

JHPSEdZid

UMEdZMX\XDJCbEdJOHHOS<cCU

4.4 Cosmic ray spectra for various nuclear species as measured near the Earth. Different curves for the same species represent measurements at various levels of solar activity. Meyer et al. (1974) ( 16°)

particles with energy below the local cut-off rigidity R which varies with geomagnetic latitude, will not penetrate the Earth's magnetic field and so not contribute to the flux, for a balloon flight from Gimli, Canada, (Lat. 50° N), R - 1 GeV nucleon-1. Above this rigidity the cosmic ray spectra can be approximated by a power law:

dN/dE = kE“ 7 (4.2)

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Integrating the differential fluxes of Fig. 4.4 for allenergies >1 GeV nucleon-1 using Equ. 4.2 and using data fromOrmes et al. (1970) ( 1J+) and Sood (1983)( 161), yields theapproximate integral fluxes shown in Table 4.1. Protons ofenergy above several MeV will be minimum ionizing, thus,providing they do not lose energy comparable to theirinitial Kinetic energy, they will deposit (dE/dr) - 2PMeV(gem-2)- 1 in the Nal(Tl). A minimum ionizing particle of atomic number Z will deposit energy given by (dE/dr)z - Z2(dE/dr) ( 1 57). Table 4.1 shows estimates of the energiesdeposited in the IC Ge detector shield by the minimum ionizing cosmic rays.

To a first approximation we can assume that the pulse of scintillation photons generated in the Nal(Tl) by charged particles can be represented by a single exponential decay with initial height proportional to the energy deposited (see Sec. 4.2.1):

nfc(ph s-1) = N EXP (-t/x) x =* 250 ns (4.3)

An amplified signal from a photomultiplier responding to this light pulse will have similar characteristics.

An anticoincidence shield system is designed to produce a veto pulse in response to any energy deposited in the shield that is greater than some threshold energy.Thresholds are set at energies typically >50 keV. This prevents too many unnecessary veto signals being produced by low energy events, such as X-rays, which may anyway be

108

prevented from reaching the Ge detector by the passive shielding effect of the anticoincidence system. An upper limit to the time during which the shield may be producing a veto signal (the shield dead time) is given by the time that the photomultiplier signal is higher than a signal corresponding to the threshold energy. During this time it is possible that the photomultiplier will saturate due to space charge effects caused by the high currents and possible depletion of charge on the dynode coupling capacitors(121 156), (this process can sometimes be beneficial since it may prevent electronics saturation).The time is given by:

T = x ln(Ed/K) (4.4)

where is the deposited energy and K is the threshold energy Table 4.1 shows estimates of the maximum shield dead

time per particle calculated with Equ. 4.4 with K = 50 keV.It can be seen that they are generally of a similar duration as typical veto pulse times of several us. Thus they pose no real difficulties for electronics design. However, this assumes that there are no long term recovery problems after photomultiplier saturation. More important than this is the result of long lived phosphorescent states in NaI(Tl)(15 121 156 157 162) (see Sec. 4.2.4).

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Particle Integral flux in shield (>1 GeV)

(s“ MEnergy deposited in shield per particle (GeV)

Shield dead time per

particle (us)Protons 1 x i o 5 3 . 0 x i c r 2 3 . 3

Helium 5 x 102 1.2 x 10” 1 3 . 7

Li Be B 10 4.8 x 10” 1 4.0C N 0 P 40 1.9 4.4Iron 1.2 20 4.5

Table 4.1 flux and energy Nal(Tl) shield.

loss by cosmic rays in the

4.2.4 The Effect of Phosphorescent States in Nal (Tl)In addition to the fluorescent light component in

Nal (Tl) , with its 250 ns decay constant, it is known that further decay components can be excited by X-rays or charged particles. Various time constants are possible depending on temperature, doping and differential energy loss in the scintillator such that(61*):

nt(ph s"1) = LOHj/Ti) EXP(-t/t.) (4.5)

where N^ photons are emitted in a component of decay time x . States with x up to several hours have been observed(15 64 121 156). However, the major contribution to these slower components is due to a phosphorescent state with decay time -1.5 us. This emission results from activation of deep metastable levels or delayed diffusion of electrons and holes from lattice traps. Under y-ray excitation about 10% of the light emitted by Nal (Tl) is in the long lived phosphorescent stated121 156 162). The scintillation efficiency for the phosphorescent state is

110

linear with differential energy loss below -1.5 MeV, although it may increase for higher dE/dr(157 152).

Several groups using inorganic Halide crystals in anticoincidence shields or as Y-ray burst detectors, in space or at balloon altitudes, have observed background counts or bursts of events which can be attributed to non-Poissonian fluctuations in the tail of long lived phosphorescent light pulses^66 156 157). These pulses are probably originally stimulated by large energy deposits due to cosmic rays. The bursts of events are also observed at altitudes below 60 mbar so are unlikely to be due to y-ray bursts. Rates of -0.005 m-2s-1 for E > 10 GeV nucleon-1 have been measured in large detectors(121). The events produce a series of unwanted veto pulses in an anticoincidence shield lasting -0.1-1.5 s (121 157)^ ana can thus continue long after the passage of the cosmic ray. This results in additional dead time in the shield.

The frequency vs. time distribution of the events, and thus the time for them to die away, depends strongly on the number of photoelectrons needed to trigger the discriminator, and on the discriminator resolution( 157). A higher discriminator threshold means they die away quicker. The distribution is also determined by the integration time of the pre-amp(163), the time constants of the dynode chain, and the decay time of the phosphorescent state. Ordinary RMS fluctuations in the number of photoelectrons from the photocathode responding to very large energy deposits could

Ill

also be large enough to simulate X-rays of energy above the shield threshold( 1 56).

Calculations of the energy deposited in a crystal by cosmic rays, assuming they deposit energy by ionization only and that the light yield is comparable to that following Y-ray excitation, indicate that the flux of Iron group nuclei (Z = 26) could account for the events. The flux of Z > 30 cosmic rays is to small. Additional energy losses due to nuclear interactions suggest that lower Z particles could also contribute.

4.3 THE HEAVY ION EXPERIMENTA flight of the Imperial College Germanium detector

from Gimli, Canada, in 1983, provided a chance for investigating the flux and energy loss of heavy nuclei, such as Iron, passing through the detector shield and likely to cause the sort of problems discussed in Sec. 4.2.Described here is the development and operation of the "Heavy Ion" experiment designed to achieve this.

4.3.1 Detection of Large Scintillation Pulses Using a Photomultiplier Operated with Low Gain

The cosmic ray events of interest in the IC Germanium detector anticoincidence shield, are likely to deposit energies up to -100 GeV (see Sec. 4.2 and Table 4.1). A system used to investigate such shield events should be capable of detecting scintillation pulses of this energy, without saturating. It must therefore have low gain. A

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wide dynamic range of perhaps 0.1-100 GeV is also desirable so that energy losses from a large range of nuclei can be detected. A photo-sensitive diode could be used as the photon detector, however, such devices have very different spectral response characteristics from the shield photomultipliers, hence calibration is difficult. An alternative technique, which was adopted for this experiment, involves the use of a low gain photomultiplier run at low voltage. Varying the EHT supply to the photomultiplier enables a large dynamic range to be obtained.

Fig. 4.5 shows the gain vs. EHT characteristic of the photomultiplier adopted for the heavy ion detector. Data for this graph was obtained using 137Cs and •133Ba X-ray sources irradiating a small Nal crystal placed in light contact with the tube. The gradient of the graph is nearly constant with a mid range value of d(log gain)/d(log EHT) =6.3 ± 0.4. Also shown in Fig. 4.5 is the manufacturers plot for a typical photomultiplier of the type used. The gradient of this graph is 7.5 ± 0.5, implying a larger change in gain with EHT. This is to be expected since the tube chosen for the experiments had an unusually low gain.

Fig. 4.5 suggests that if heavy ions are to be detected, without dropping the EHT below the recommended minimum operating voltage of 400 volts, then the gain of the photomultiplier must be reduced by a further factor of about 2 orders of magnitude. This was achieved by partially masking the tube from the scintillator. The number of

fluorescent photons reaching the photomultiplier after an interaction in the shield is then reduced, and hence the gain of the system is lowered.

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Fig. 4.5 Gain vs. EHT characteristic of the photomultiplier adopted for the Heavy Ion experiment: (a) measuerdwith X-ray sources and; (b) measured with the LED. (c) The characteristic as measured by the manufacturer for a typical tube.

4.3.2 ElectronicsFig. 4.6 shows a block diagram of the electronics

developed to run the masked photomultiplier during the balloon flight. An EHT unit supplies the high voltage to the photomultiplier which views the Nal(Tl) shield through a hole in a metal mask. The photomultiplier is surrounded by a Mu-metal anti-magnetic shield. The voltage generated by the EHT unit is determined by a low voltage reference signal supplied by a staircase generator. This ramping reference voltage results in a ramping EHT and hence a changing

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photomultiplier gain. A dynamic range in gain of about 1000 is possible. Analogue signals from the photomultiplier pass to a discriminator with a fixed, but adjustable, threshold voltage. The EHT reference voltage, and the discriminator count rates, are incorporated into the telemitery data format for transmission to ground and subsequent analysis. Frame periods of 1 ms were used.

Fig. 4.7 shows the staircase generator circuit. It provides an EHT reference voltage which repeatedly rises from 2 to 7 volts in 1024 s, in 256 4 s steps. A cycle time of 1024 s was chosen because a time much longer than this might result in significant changes in the environment around the telescope during the cycle. A shorter time might result in too few counts per cycle. Equal voltage rises per step were chosen since this produces an exponential increase in the energy channel width when the differential spectrum is formed. This partially compensates for the lower count rate expected for higher energy events. The EHT unit(151+) was adjusted to produce output voltages of 400-1200 volts, depending on the reference voltage supplied.

4.3.3 Preflight CalibrationAn increase in EHT supplied to the photomultiplier

results in larger photomultiplier pulses for the same energy loss in the scintillator. This represents a decrease in the energy equivalent of the threshold voltage. At a given EHT the discriminator count rate corresponds to all events with

115

Fig 4.6 Block diagram of the flight electronics.

Fig. 4.7 Diagram of the staircase generator circuit

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energy loss greater than this threshold energy. Thus the difference between count rates recorded for successive EHT values gives the differential count rate for events with energies in a channel of width determined by the difference in energy of the energy thresholds corresponding to the EHT values. To calibrate the system it is necessary to find the relationship between the EHT reference voltage and the effective discriminator threshold energy level.

A calibration procedure was developed to obtain a plot of EHT unit reference voltage, vs. equivelent energy level of the discriminator. Fig. 4.8 shows a block diagram of the system used. A gr-aph of relative photomultiplier gain vs. EHT was first obtained over the whole EHT range by using a green light emitting diode to simulate scintillator pulses. This was placed on a shield window. The amplitude of the diode could be continuously varied thus effectively providing a variable event energy. The energy could be calibrated by using a pulse height analyser (PHA) to compare the pulse heights due to X-rays from radioactive sources with pulse heights due to the diode emission. The dynamic range obtainable was much larger than could be achieved with the available X-ray sources alone.

The gain vs. EHT characteristic taken with the diode is shown in Fig. 4.5. The slope of this was found to be the same as that obtained with a 137Cs source. A tailing off of the characteristic at high energies is due to EHT unit saturation, whilst non-linearity at low energies is an

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effect of the photomultiplier when run at low voltage.'

Fig. 4.8 Block diagram of electronics used to calibrate the Heavy Ion detector.

The discriminator was adjusted so that when the photomultiplier views the scintillator without a mask, and is supplied with its maximum possible EHT of =1200 volts (corresponding to maximum reference voltage), then the threshold voltage is set at an equivalent energy of 1 MeV. This was achieved by adjusting the discriminator threshold voltage until the discriminator count rate was the same as that measured above 1 MeV with a PHA. With this calibration, and the graph of Fig. 4.5, it is possible to find the effective discrimination energy at any EHT reference voltage. The addition of a mask to reduce the photomultiplier gain could then enable the range of effective discrimination level to be adjusted to about

118

100-100,000 MeV.A mask consisting of a thin A1 disc painted matt black

on .both sides with a -2 mm diameter hole at the centre, was found to give the most reproducible reduction in gain. This mask fitted neatly into the photomultiplier housing between the tube and the shield window. It resulted in a gain drop by a factor of 110, measured using X-ray sources and a PHA. Various other masks were produced with gain reduction factors of 1 to 1000 (±5%).

The final calibration graph is shown in Fig. 4.9. It must be emphasized that implicit in this calibration are several significant assumptions. It is assumed that the form of the gain vs. EHT characteristic of Fig. 4.5 is independent of the size of the scintillation pulse or diode light pulse used to measure it. Experiments using the largest diode pulses possible indicate that linearity is maintained up to X-ray equivalent energies of -4 x 105 keV. However, the assumption was also made that the diode pulses can simulate scintillation pulses due to X-rays over the whole energy range. (Fig. 4.5 suggest this is true over at least part of the range). It is also assumed that the scintillation efficiency for X-rays is linear up to energies of 100,000 MeV and is similar to that for other radiations. This may be true only over some of the range (see Sec.4.2.1) .

Thus the discriminator threshold energy as measured can at best only represent an estimate of the energy equivalent

of a scintillation pulse if the pulse had been produced by an X-ray, and if the X-ray scintillation efficiency of Nal (Tl) is constant up to very high energies. Also, it should be noted that the size of the photomultiplier signals was found to vary by a factor of about 2 depend -)y}cj on theposition of the light pulse in the shield.

vref (v)Fig. 4.9 Graph of EHT reference voltage vs. the energyequivalent of the discriminator threshold voltage.

4.4 FLIGHT RESULTS

4.4.1 Flight DetailsThe 1983 flight of the Imperial College Germanium

detector was made in collaboration with the University of Calgary, Canada, who provided the steerable balloon platform. A successful launch was achieved from Gimli,

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Manitoba, (lat. S O ^ T ^ S 11 N) at 13h24 UT (08h24 local time) on 24t 1 Aug. 1983, using a 22 Mcf Winzen balloon.The balloon reached altitude after 3.5 h and driftedwestwards for a further 10 h before a slow leak necessitatedan early termination.

Fig. 4.10 and 4.11 show the altitude and trajectory of the balloon. The spectrometer made observations of the Crab, NGC4151, and HERX-1. Unfortunately the data has been difficult to analyse due to microphonics interference fromthe Germanium and periods of noise spikes from the shield,probably due to coronal discharge. The slow descent of thedetector during the flight resulted in a graduallydecreasing sensitivity.

£ 160cno 120 o o- 80ta Qg 4 0M EH

12 16 20 0 4UT (24/8/83)

Fig. 4.10 Balloon altitude profile.52

cn<DT35 1

Eh<

50LONG. (deg)

Fig. 4.11 Balloon trajectory profile.260 262 264

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4.4.2 Results and InterpretationFig. 4.12 shows the raw count rate from the heavy ion

detector as a function of EHT reference voltage, integrated for the period 16hl5-lh35 UT. Since the expected count rate was very low, only single heavy ion counts occurring in any of the -1 ms frame periods, were accepted as valid. A count of greater than 1 was likely to be a bit error. The expected increase in count rate, due to the decrease in energy threshold as the EHT increases, is clearly seen. Unfortunately an electronics error in the interface with theCanadian telem^^ery system reduced the ramping voltagesupplied to the EHT unit by about 1 volt and hence increasedthe threshold energy values of the discriminator.

10°

~ 10“ 1 Icn4Ju

g 1 0 “ 2 <Oh

H2DOCJ

10“ 3

10“**1 10 100 1000

ENERGY THRESHOLD (GeV)Fig. 4.12 The raw count rate from the Heavy Ion detector.

t

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Fig. 4.13 shows the differential energy loss spectrum for the whole period of 16hl5-01h35 UT, using the calibration of Fig 4.9. A best fit power law to the data gives a spectral index of =3.0. The calibration indicates a range of energy losses from 1.2 GeV to 2.5 * 103 GeV. However, as indicated in Sec. 4.3.3, it must be emphasized that this energy scale represents only a crude estimate of the energy deposited in the shield and assumes the events behave like X-rays.

Fig. 4.13 Differential energy loss spectrum from the Heavy Ion experiment for the period 16hl5-01h35 UT.

The total count rate over the whole energy range is0.23 ct s"1. Comparison with Table 4.1 indicates that this isat least 2 orders of magnitude less than would be expectedif the energy calibration is correct. An upward shift in

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the energy scale by a factor of only 2 could account for this discrepancy by insuring that only nuclei with atomic weights equivalent to Iron or above can deposit enough energy. The discussion in Sec. 4.2.1 indicates that a factor like this could easily be accounted for by the apparent decrease in Nal scintillation efficiency with increasing heavy ion atomic number.

The count rate observed can be accounted for by the detection mainly of nuclei of atomic number >20. For such particles the nuclear interaction cross section becomes significant( 1 65 1 57). This partly accounts for the dropping scintillation efficiency with increasing atomic number. The nuclear interaction mean free path, X, for cosmic rays of atomic number Acr, interacting in a medium of atomic number At, is given by:

X « (At1/3 + Acr1/3)-2 (4.6)

Foe Al, X = 10.5 gra cm-2. The value is smaller for elements of higher atomic number, thus one might expect materials in the surroundings of the detector to play a part in determining the count rate.

Fig. 4.14 shows the heavy ion count rate above the lowest discriminator threshold against time during the flight. The count rate appears to vary during the flight and by the end of the flight it has dreflped by a factor of about 2. The drop in altitude of the detector from a residual atmosphere of 1.5 gm cm”2 to 4.5 gm cm”2 (see Fig. 4.10)

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could only contribute a maximum of 10% to the general drop in count rate. This is more than offset by a predicted increase in count rate of =15% due to a drop in cut-off rigidity resulting from the northerly component of the balloons trajectory (see Fig. 4.11). However, the changes do show some correlation to detector pointing, and therefore to changes in the orientation of the detector with respect to the other matter in the payload. For Iron nuclei in the air of the upper stratosphere A = 12 gm cm"2(165).

_0.4—HIcnj_> 0 . 3 o“ 0.2< ongo.ioCJ

05 10 15

TIME (45 min intervals)Fig. 4.14 Total Heavy Ion count rate vs. time during the

flight.A sample of 45 minutes of data was examined to

determine whether any heavy ion counts detected with themarsked photomultiplier, were associated withnon-statistical fluctuations in the normal shield count ratefrom the normal photomultipliers. Shield counts in 1 mstime bins close to each heavy ion event were verified toobey Poissonian statistics. However, of the 131 heavy ion

events studied, all with energy above about 1.2 GeV (according to the calibration of Fig. 4.9), 12% were associated with shield counts per bin that were >3a above the mean, occurring within -5 ms of the heavy ion event. 4 of the heavy ion events were associated with shield counts per bin which were >10a above the mean. 3 of these heavy ion events had an equivalent energy >3 GeV. Of all heavy ion events with energy >3 GeV, 21% were associated with shield counts per bin >3a above the mean.

Thus coincidences between heavy ion events and bursts of shield counts do seem to have occurred and appear more likely to do so if the heavy ion event is of high energy loss. However, it is important to realize that the form and frequency of shield bursts will be determined mainly by the electronics design (see Sec. 4.2.4). This makes interpretation of the events difficult.

4.5 CONCLUSIONThe results outlined in Sec. 4.4 suggest that the

events of Fig. 4.13 are almost certainly due to high energy nuclei with A > 20. They occur at a rate of =0.23 ct s” 1 and deposit energy greater than about 1 GeV. Although the count rate is low, the event energies are quite high, thus they could be producing significant long term (>1 ms) effects in the shield. Certainly some spikes in the shield are correlated to heavy ion events.

A future flight of the Germanium detector may need to

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include some system or electronics to ensure that the high energy loss events do not cause dead time in the system.For instance, an anticoincidence shield of plastic scintillator, which is less susceptible to long lived fluorescent states than Nal, could be arranged with a high threshold energy and used to veto the long lived Nal shield events( 1 56 1 57)% Detailed electronic analysis of the rise time or spectra of the Nal events could be used to gate them out, perhaps by suppression of shield photmultiplier gain by dropping the EHT, or by changes to the electronics time constants.

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CHAPTER 5

USE OF A WIDE APERTURE DETECTOR TO SEARCH FOR AND STUDYX-RAY PULSARS

5.1 INTRODUCTIONThe developments in escape gated detectors outlined in

Ch. 3 were inspired mainly by technological ideas rather than any specific astronomical need. The discussion in Ch. 4 concerning Germanium detectors was inspired largely by a general astronomical need to improve the sensitivity of these detectors. Discussed in this chapter is an assessment of a proposed instrumental technique which has been inspired more directly by a specific astronomical need. The technique involves the use of a large area, wide aperture (2tt) , detector to search for and study periodic X-ray flux from X-ray pulsars. Such an instrument could help solve some important questions in X-ray pulsar astronomy.

Assessment of the technique was performed by using data from a flight of the Imperial College 6.3 m2 y-burst detector. Although this detector was not designed to search for or study X-ray pulsars, its large area, high sensitivity, fine timing resolution and wide aperture, fulfil most of the requirements for a detector proposed for this application. It was considered that a search in the data from this detector for periodicity from known X-ray pulsars would provide a sufficient test of the proposed instrumental technique.

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5.2 X-RAY PULSAR ASTRONOMY - THE NEED FOR A NEW OBSERVATIONAL TECHNIQUESince the discovery of X-ray pulsars in the early

1970's, their study has become an important branch of X-ray astronomy. However, although a considerable amount of observational and theoretical knowledge has been accumulated, many important questions remain unanswered (see Sec. 2.2). For example, why are there so few Crab-like pulsars? Progress in instrumentation for hard X-ray/low energy Y-ray astronomy is helping to confront some of the problems, especially those concerning the X-ray binary pulsars( 118). For instance, the OSSE experiment proposed for GRO (see Sec. 1.5) will have 1 of 4 modes of operation dedicated to X-ray pulsar work. The proposed balloon experiments FIGARO(1 ), and ZEBRA( 11 2), will have large areas and also be capable of pulsar work. However, the specific problems concerning isolated X-ray pulsars might best be tackled with instrumentation specifically designed to search for X-ray pulses from new examples of these sources. The pulse characteristics of those discovered could then be compared with the models (121f 131).

Some recent observational and theoretical results have highlighted the possible success that searches for X-ray pulsars can have (see Sec. 2.2.4). The discovery of new pulsars of all forms would be beneficial, for instance in determining the distribution of pulsar periods(125 166 167). Many theories predict that fast pulsars and pulsars in

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nebulae should be easier to find at X-ray energies than in other wavebands( 1 32 1 33 169 17°).

Though recent X-ray detectors have had some limited success at finding new pulsars, and a few of the proposed new detectors have some potential for furthering the work, most are not well suited to a general pulsar search. For instance, many of the detectors (see Table 1.1 and 1.2) have

/ rfields of view ti_ypically£ 2 0° and so require pointing at/specific targets. This has necessitated making assumptions

about the location of pulsars. Past targets have included; i) suspected ‘Crab-like SNR's^1 3 1 17°), ii) radio pulsars predicted to emit X-rays ( 1 33 131+ 171), iii) sources known to show intensity variations( 1 2 5 172), and iv) sources known to have hard X-ray spectra( 1 66). Also serendipitous detections of X-ray pulsars have occurred( 173).

An improved pulsar search technique might be to use a large area detector with an -2 tt aperture, to search for any periodic X-ray flux down to a few ms. This could not only establish the existence of new X-ray pulsars more easily but also provide much information about their pulsed flux. The precise periodic nature of the sources and their relative rarity suggests that identification of different pulsar periods in data would not be a problem(116). The study of known pulsars could be achieved since they could be identified by their periodicity(124). Though the detector could provide no angular information some might be obtainable by identification with the periods of known radio pulsars.

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A wide aperture detector would be incapable of measuring the non-pulsed components of pulsar emission.However, this would not be a severe disadvantage since for many such sources the pulsed flux arises from a different mechanism than the non-pulsed flux and can be treated separately( 1 74 175). Measurement of the profile and spectrum of pulsed X-ray flux can therefore help to answer some of the astronomical problems mentioned in Ch. 2(176).For instance, general observations of pulse shape and pulsed spectra together with measurements of correlations with intensity variations could be used to infer emission mechanisms^5 1 33 177). Indeed, accurate measurements of the p and dp/dt of isolated pulsars is very important since the rotational energy loss is generally accepted to be a possible origin of the energy output. Measurements of dp/dt alone can help to find the orbital period in a binary system^ 1 3 2 178).

Thus it appears that a sensitive, wide aperture, detector could provide a useful contribution to X-ray pulsar astronomy.

5.3 THE IMPERIAL COLLEGE 6.3 M2 DETECTOR AND DATAThe feasibility of using a large area open aperture

detector to investigate periodic X-ray flux from X-ray sources was assessed using data from the 1978 balloon flight of the Imperial College 6.3 m2 y-burst detector.This detector was built as part of the IC y-burst astronomy programme (Mills (1978) ( 179), Beurle (1983) ( 16i+)) . Outlined

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below are the essential characteristics of the detector and the nature of the data obtained from the 1978 flight.

5.3.1 The 6.3 m 2 DetectorFig. 5.1 shows a diagram of the IC 6.3 m 2 y-ray

instrument. Detection of y-rays was achieved with 18 slabs of NE 102A plastic scintillator. Each slab is viewed by 2 photomultipliers. The slabs are arranged in trays and mounted on an Aluminium raft together with Lithium batteries, EHT units, and electronics. Some trays were inclined at 30° to the horizontal to provide some directional sensitivity to y-bursts. Polystyrene sheet, 20 cm thick, provides a thermal jacket.

WijaMilf1 # 1 3

0TRAY F 0

O0

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o TRAYE O 0

TRAY A0 c

TRAY A o TRAY C O

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0 0u u o o o

o TRAYE O

—

TRAY B 0 o TRAY BO o TRAY C o

mu V— ■ ' 0TRAY 0O

0o

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— 10 ' -----------------------------------------------------

O P H O T O MU L T I P L I E R Sf U S f e j / f i e t i P O iA T T f

Fig. 5.1 Schematic diagram of the 6.3 m2 y-burst detector.

132

The energy range of the detector was chosen to be 50- 2000 keV, consistent with the requirements of y-burst astronomy. Both upper and lower energy levels were determined by discriminators. The upper level, set to reject cosmic ray events. Unfortunately all events detected within the energy range are summed into one energy bin, thus providing no energy resolution. Fig. 5.2 shows the efficiency response of the detector, and Fig. 5.3 shows the angular response. These were both calculated using Monte Carlo programs, and verified experimentally with 57Co and 60Co sources. (Beurle (1983) ( 164)).

Provision was made to measure the altitude of the detector with a barometer. However, this had ill-defined temperature characteristics and was thus subject to error. Provision was also made to measure the orientation, range, and direction of the payload. A stable quartz crystal frequency standard kept UT to a precision of <10- 6 s.

Valid events from each individual slab, and each tray, and for the whole detector, were accumulated into 1 , 1/16, and 1/256 s data time bins respectively. These count rates were arranged into a telemetry frame, along with house keeping data, of 4096 bits, for transmission to ground once per second.

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Fig. 5.2 Efficiency response of the 6.3 mz detector.

Fig. 5.3 Variation of efficiency with angle.

5.3.2 Flight and DataThe 1978 launch of the detector took place at 20h00 UT

on lSth Nov. (0530 CST 16th Nov.) from Alice Springs, Australia, (lat. -23°42', long. 133°52')* A float altitude of 3 mbar residual pressure was achieved at 23h39 UT.Residual pressures of <6 mbar were maintained until cutdown

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at 08h00 UT 17fch Nov. Unfortunately telemetry became noisy in the later part of the flight and was lost completely at 02h30 UT until cutdown, except for 1 h of data recieved at a down range station.

The detector functioned quite satisfactorily during the flight. Original analysis of the data for the purposes of identifying Y-bursts was successful in revealing 2 candidate Y-bursts (Beurle (1983) ( 161+) ) . Most slabs produced reasonable Poisson noise distributions throughout most of the flight, however a few slabs produced almost random numbers of counts per time bin for some of the flight. These spurious counting rates resulted from coronal discharge in EHT potting, or instabilities in EHT and discriminator units. They took several forms; intermitant spikes of increased count rate lasting a few ms found in most slabs at various times, sizable drops in count rate lasting several minutes, and sudden fluctuations in count rate. Other variations in count rate were associated with diurnal changes in altitude and a few geophysical events.

5.4 DATA ANALYSISIt would be possible to conduct a search in the data

from the 6.3 m 2 detector for periodicity due to any X-ray source with a period >4 ms (The best time resolution of the detector). However, for the purpose of assessing the sensitivity of the detector to X-ray pulsars, a search for a selection of known pulsar periods was considered sufficient.

135

This would make more economic use of valuable computing power and research time but could still provide some useful observational results. The data analysis procedure adopted to achieve this can be divided into 3 steps: i) To know what period to search for in the data a prediction is needed of the period of each chosen pulsar at the position and epoch of the observation. ii) Errors in the predicted period mean a limited search in the data for indications of the actual pulsar period is still required. In this way either values for the detector pulsed counts, with an assigned significance, or an upper limit for it, can be found. In the case of a positive detection a pulse phase histogram (light curve) can then be deduced. iii) The count rates obtained should be converted into values for the pulsed flux from the source for comparison with previous observations or theories, and to determine the sensitivity of the detector.

5.4.1 Pulsar Period PredictionA prediction of the expected period of a pulsar

during a new observation can be calculated using previous measurements of the pulsar period p and period derivative dp/dt. These parameters are usually quoted in solar ecliptic coordinates at some epoch not necessarily contemporary with the new observation. To find the period of the pulsar at the time and location of the new observation two adjustments to the data are needed: i) A correction for the pulsar period drift between the old and

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new observations. ii) A correction to find the period of the pulsar at the location of the detector.

5.4.1.1 Correction for Drift of PeriodMost X-ray pulsars are undergoing some degree of

spin up or spin down (see Sec. 2.2), therefore the period of a pulsar will change between the epoch of the observation used for the prediction, Jzt and the epoch of the new observation, Ji. This period drift can be corrected for using the equation:

Pjl = Pj2 + (Jl - j2> if (5.1)

where pj i (s”1) and pj2 (s”1) are the periods for Julian day Jl and J 2/ and dp/dt (s d-1) is the period derivative.Since the observation will occur over a finite length of time, account must be taken of the pulsar period drift during the observation. We can consider the situation as a drift in pulsar phase with the actual period approximately constant. If the period drifts from p to p + Ap (s) in time t (s) then the time derivative of the fractional change in period is given by (l/p)dp/dt = Ap/(pt). The phase drift in one period per period is 6pi = p2Ap/(pt). The drift after n periods is given by:

6pn = 6pn_! + np2( || ) = ip2( |E ) = 0.5n2p2( || ) (5.2)

We can consider <5pn to be significant if after a time T = np 6pn > O.lp, in which case:

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T > [ 0.2p2t/Ap] 1 / 2 (5.3)

For the Crab pulsar, which has a high dp/dt (-4 * 10” 13

s s” 1), T > 6.5 h. Thus, for the short observations available from balloons, the effect is usually small.However, the extrapolation of Equ. 5.1 is often subject to several other effects that are more difficult to assess: i) p and dp/dt may be ill-determined. ii) There may be period derivatives of higher order than dp/dt. iii) Jumps or "glitches" in the pulsar period may have occurred in the time between the two observations. iv) The pulsar period may be subject to irregular fluctuations or effects due to incorrect evaluation of its position.

5.4.1.2 Corrections for Detector Locationi) Correction for the Doppler shift due to the Earth's orbit: The Earth's orbit around the sun produces a Doppler shift in the period of a pulsar as observed from the earth.To correct for this it is usual to first find the position of the pulsar in ecliptic coordinates (3, A) by transforming from coordinates in right ascension and declination (a,6) using the equations( 1 80):

sin3 = sin<$cose - cososinasine (5.4)sinA = (cos(Ssinacose + sinSsine) /cos3 (5.5)

where e is the Earth's tilt angle. These ecliptic coordinates can then be converted into solar ecliptic coordinates (3 ,A') using the relation:

138

X' = X - (DN - 80)360/365 (5.6)

where DN is the day number. The component of the Earth's velocity in the direction of the pulsar is then given by v = -Vcos8sinX' (V - 2.9 * 101* m s-1). The Doppler correction, given in terms of the pulsar frequency, is thus:

(Af/f)orb = v/c = -(V/c)cosesinX' (5.7)

The period, p2r as observed from the earth is then given in terms of the period, px, in solar ecliptic coordinates, by:

p2 = Pi ( 1 + (v/c) )“ 1 (5.8)

For most observations from balloons the change in orbital Doppler correction over the duration of the observation can be neglected.ii) Correction for the Doppler shift due to the Earth's spin: For observations from a balloon borne detector for which the latitude and longitude remain approximately constant, the Doppler shift in pulsar frequency due to the Earth's spin is given by:

(Af/f) in = v'/c = (V/c)cos 6sin$ (5.9)

where Vcos<5sin<& is the component of the velocity of the detector due to the Earth's spin, in the direction of the pulsar, V is the Earth's rotational velocity at the latitude of the detector, <5 is the declination of the pulsar, and $ is the hour angle. The optimum time for an observation is

139

when $ - 0. The source is then at its highest in the sky and atmospheric absorption is at its lowest. For most observations (Af/f) an< d (Af/f) Sp^n/d ( canneglected.

5.4.2 The Search for Periodicity in DataThe uncertainties in the values of a predicted period p

of a pulsar for a new observation are likely to render meaningless a simple folding of the data at p to obtain the light curve. A procedure to test sample periods around p is therefore required. Values for pulsed counts at the detector with an assigned significance, or an upper limit for it, can then be deduced.Outlined below is; i) an idealised procedure to achieve these aims, and ii) a discussion of the modifications needed to apply the procedure to data from the 5.3 m2 detector.

5.4.2.1 An Idealized Procedure for a Periodicity Search Two major techniques are available to search for

periodicity in a data stream? i) Fourier analysis( 181), in which the power in the sample frequencies is measured, and ii) epoch folding( 1 3 7 182), in which a x2 value is obtained from the pulse phase histogram corresponding to each sample period. Though Fourier analysis is a faster and more thorough technique than epoch folding, this advantage is reduced when the range of periods searched for is small, or the observation time is short. Furthermore the S:N ratio increases only as the 4th root of the time rather than the

140

square root (see ii) below). Also, it is often more difficult to deal with data gaps in Fourier analysis than in epoch folding. Therefore epoch folding was the technique chosen for the present purpose.i) Epoch folding and the x2~test: Epoch folding of a data stream predicted to contain pulsed counts, requires binning the arrival times of the events into N phase time bins of equal duration modulo a range of sample periods, p', around the predicted period. This results in a series of pulse phase histograms to which a x2”test can be applied. The particular form of x2”test adopted here entails testing each histogram against a constant mean in all phase bins (the "null", or "constant intensity", hypothesis). A distribution of x 2 (or x ) vs. period can be obtained. We define the x 2 to be:

2 _ v N fn; - (n/N)]N Ai = 1 ~*~o~za. 1(5.10)

where n^ is the count in the ith bin (the base rate) and n is the total counts in a histogram. n/N is the mean count per phase bin and ck 2 is its variance. For Poisson statistics ck 2 = n/N and the expectation value of x2N is (N-l) . The variance of x 2 »T (for N > 20) is -2N. In all phase histograms n will remain the same since the same data is used in each case. However, if regular pulsations are present at one of the sample periods, p1, then some of the phase bins in the corresponding histogram will rise above n/N and others will fall below. The x2-test should reveal

141

this by yielding a value of x2N higher than the expectation value for that period. The higher the x2N the more structured the pulse profile(183). No knowledge of the form of the pulse profile is needed for the technique, ii) Sensitivity of the x2~test; If we interpret a high x2jq to be the result of pulsed counts at the sample folding period, p', then we can deduce the statistical significance of the pulsed counts. Consider a hypothetical situation in which the pulse profile is rectangular and of period p, duty cycle 3 (defined as the proportion of a cycle with counts above the mean), and pulsed counts c, are present in the data. Fig. 5.4 shows the result of folding the data at

bins: 1 2 (1-3)N NFig. 5.4 Pulse phase histogram with N phase bins for data

with period p, duty cycle 3, total counts in the histogram n, and pulsed counts c.

The total counts in the histogram (Fig. 5.4) must remain n (which includes ?)• We can interpret the situation as a redistribution of the counts in each phase bin, compared to the distribution for Poissonian data, such that for phase bins 1 to (1-3)N the counts fall below the base level n^ by

142

C/N . For the remaining bins the counts are increased from n^ by [(1-8)C]/(8N) . This formalism ensures that the totalpulsed counts, defined as the count above the new base level of n^ - (C/N), is c. This is how the pulsed counts are defined in actual data (see Sec. 5.5). The x2N for the distribution is then given by:

2 _ N v ( 1- 3 ) NXN SL T (n. - ) - £ l2n z*i=i L u l N ' N JN yN [ (n. + iizilin i = 1+ ( 1 - 3 ) N i 3N

) - £ ] 2 (5.11)5 <K nexpanding Equ. 5.11 assuming Poissonian statistics yields:

(1-8)C2X2n - (N-l) + 8n (5.12) 1 > 8 > 0

(This result is in agreement with the calculations of Boldt et al. (1971) ( 172) and Buccheri et al. (1977) ( 18t+) ) .

Given that the data has been analysed and interpreted as above, the most probable value for the pulsed counts, c, is given by Equ. 5.12. The approximate statistical significance of the pulsed counts can be expressed in terms of the number, s, of standard deviations, a, of x2N that accounts for the difference between the measured x2N and the expected value of (N-l), hence, rearranging Equ. 5.12:

^ 2 - 8nsa1-8 (5.13)

If N is small then a can be estimated from probability tables of X2N« An adequate estimate of the 3a minimum detectable pulsed counts can be obtained from Equ. 5.13

143

with s = 3, though a more rigorous analysis by Pravdo et al. (1976) ( 185) suggests that this gives rather low upper limits. Boldt et al. (1971) ( 172) indicate that Equ. 5.13 is valid for any shaped pulse provided B and S are maintained. However, rectangular pulses are a good approximation to the profiles of many pulsars.

Several conclusions can be drawn from Equ. 5.12 and5.13: i) If we define the pulsed fraction as f = C/n thenthe minimum detectable pulsed fraction f . * n” 1/2. Hence,^ m m 'f . a t“ 1 / 2 where T is the duration of the observation (in minFourier analysis f . * n”1/1+ (184)). ii) The highestsensitivity (smallest fm ) occurs for pulsed data, folded with a given N, such that there is a short sharp peak in the pulse phase histogram (B is small). The sensitivity remains unchanged if a peak in the histogram is divided into several shorter peaks as long as c and B are kept unchanged,iii) If we assume a - 2N, then fm^n * N 1/1+. Thus decreasing the number of phase bins used will increase the sensitivity of the technique, provided a lower number of bins does not alter the shape of the pulse phase histogram too much.

Many pulse profiles can be approximated by a sin wave and adequately described by a minimum of about 8 phase bins. A smaller number of bins is likely to result in a less sharp peak in the pulse phase histogram. If the pulse profile is accurately known then a cross correlation of the data with the known pulse profile is likely to be more sensitive than the x2~test outlined above.

144

iii) Interpretation of the x 2 vs. period distribution; Equ. 5.12 describes the significance of a particular enhanced x2N for a particular sample folding period p' in terms of likely pulsed counts. However, if pulses are present then a complete scan about the predicted period may result in significantly enhanced x2N 's for a range of sample periods. This is because folding data at a period p* = p ± Ap, where p is the period of the pulses, is equivalent to increasing the pulse duty cycle from 3 to 3 + A3. This results in a reduced but possibly still significant X2N« The amplitude of the x2n distribution can still be related to the pulsed counts with Equ. 5.12, but the confidence level of a detection should really account for the whole x2N vs. period distribution.

1 2 3 4 5 5 7 8 9PERIOD (arb units) p'

Fig. 5.5 xN vs. period distribution for rectangularpulses with 3 = 0.5 and: (a) 1 peak per period;(b) 2 peaks per period (adapted from Buccheri ( 1Q1+) ) .

145

Fig. 5.5 shows distributions for rectangular pulseswith 1 and 2 peaks per period (adapted from Buccheri et al. (1977) ( 184) ) , (NB Fig. 5.5 and following similar Figs plot xN vs. period). They can be interpreted as follows:As the folding period p' increases or decreases from p by ±Ap, the x2n drops as expected. The number of folding operations used to obtain a particular x2N is F = T/p' =T/(p ± Ap) where T is the duration of the observation.Minimum x2N occur when F is an integral number r greater or less than the value of F when Ap = 0 (F = T/p), so:

For instance, if r = 1 then at the end of the the folding process we are folding data exactly one period out of phase with respect to the process for r = 0. Thus for half the folding process the pulses have been tending to add and enhance a peak in the pulse phase histogram. For the remainder they have been tending to reduce it. From Equ. 5.14 the position of the minima away from p is given by:

(This is in agreement with the results of Knight et al.

F = (T/p) + r = T/(p + Ap) (5.14)

(5.15)

(1982)(134)). The main peak of enhanced x2N covers a range of periods about p of:

^Ppeak )/T (5.16)

though this reduces with limited statistics and small 6. If there is more than one pulse per period then additional

146

minima occur and the maxima width is reduced.One technique for calculating the confidence level of a

detection using the whole x2N vs. period distribution is to introduce the binomial equation(184):

where Pr is the probability of the x2N distribution occurring randomly. It is the probability that of m independant x2N 's, k of them are greater than a threshold X2q where R is the probability of x 2 N > x Q 2 « R can be estimated from x2N probability tables or by assuming a Poisson distribution of X2N»

Equ. 5.17 is really only valid if the data obeys normal statistics when no pulses are present and if the x2N ate independant. For N phase bins the x2N ate independant if the steps between folding periods APstep > P 2/(TN). This is because if the data is folded at a period p'' = p' + (p2/(TN)), rather than p', then at the end of the folding process we are folding data out of phase with respect to the folding process for p' by one phase bin. The allocation of events to phase bins will ensure that half of the data used in the folding at period p'' will fall in the same bins as they would for folding period p'. The rest will fall in different bins. We can choose a minimum step size to be used in a period scan as:

Pr ” ) Pk(l-R)m_k (5.17)

Apstep = <P2/(TN)) (5.18)

147

Steps much larger than APstep may result in loss of information about the shape of the x2N distribution.Smaller steps are sufficiently interdependant as to provide no new information and thus waste computing time.

5.4.2.2 The Search for Periodicity in Data from the 6.3 m 2 Detector

Discussed below is the development of modifications to the analysis processes, needed to enable their application to data from the 6.3 m2 detector, and how these procedures were implemented.i) Epoch folding data from the 6.3 m2 detector: Analysis of data from the 6.3 m 2 detector proceeded with computer programs developed to implement the epoch folding described in Sec. 5.4.2.1. The process involved binning the counts in the data time bins from the telemetry format into N phase time bins. It was ensured that each phase time bin had the same number of data time bin entries so that the counts in each phase time bin were accumulated for the same amount of time (correction for phase dependent dead time). A consistent procedure was adopted to deal with the counts in data time bins that fell on a boundary between phase bins. Most invalid counts in data time bins occurring due to telemetry errors were eliminated by ensuring the presence of a correct frame synchronisation pattern, and that bit parity checks were passed. For each observation a x2N vs. period distribution was obtained for a range of sample periods, p', in steps of 1-3 APstep/ covering ~4 APpea]< centered on a

148

predicted period p (see Equ. 5.18 and 5.16). 8 phase binswere generally used (see Sec. 5.4.2.1).

For the maximum timing resolution of 4 ms, it was necessary for the analysis program to use data summed from all detector slabs. This inevitably meant spurious data would be encountered (see Sec. 5.3.2). However, the presence of a regular period p in such data can still be confirmed using epoch folding, provided the spuriosities are random. This is because their significance in the pulse phase histogram will tend to be reduced by epoch folding at p, relative to the pulsed counts. However, this conclusion is likely to be true only if the spurious data last for short times compared with the folding period and occur with frequency much greater than T/N , or have time characteristics much larger than the folding period. Thus it was likely that the presence of the infrequent and short spurious data identified by Beurle (1983) ( 16l+) , would result- in misleading pulse phase histograms.

It was possible to eliminate most of the spurious data by ensuring that only counts per data time bin of amount less than a number s of standard deviations from the mean count rate were accepted for analysis. The mean was taken for a time At covering preceding data time bins. Both s and At were program variables adapted to suit the particular spuriosities found in each data run. Fig. 5.6 shows an example of a vs. period distribution for 4000 s of databoth with and without reduction of the spurious data using

149

the technique. (There should be no periodicity in the data used to obtain these plots at the periods analysed). Though the modified x2N ate distributed closer to the expectation value, and are found to have a more Poissonian distribution, there is still residual non-statistical behavior. Beurle (1983) ( 161+) suggests such behavior may be due to subtle telemetry errors.

89235 89240 89245 89250PERIOD (us)

Fig. 5.6 A xN vs. period for 4000 s of data: (a) withoutdata smoothing; and (b) with data smoothing.

ii) Sensitivity and interpretation of epoch folding appliedto data from the 6.3 m2 detector: The non-Poissonian natureof the data necessitates modifying the evaluation of thestatistical significance of the X2N- This is because Equ.5 .1 2 , relating x2N to the pulsed counts, can not be directlyapplied if the data is too non-Poissonian. However, therelationship can be determined empirically by forcing intothe data time bins different numbers of pulsed counts atperiods equal to the folding period. Fig. 5.7 shows anexample of a plot of forced pulse counts vs. x2N obtained

150

in this way. Each empirical point and error bar represents the mean and s.d. of 1 2 x2N 's obtained with the same data for 12 similar but independent (see Sec. 5.4.2.1) folding periods. Square pulses of 8 = 0.5 and 0.25 were used. Comparison is made with results from Equ. 5.12. The effect of the non-statistical noise is to increase the expected x2N for any particular pulsed counts, and to increase the spread o f X 2 n .

Fig.

______i______I______» - I______ «0 5 10 15 20 25 30

FORCED PULSED COUNTS (ct s“ 1)5.7 Forced pulsed counts vs. x2N with 8 phase bins:

(a) result of Equ. 5.12 with 12000 s of data and 8 = 0.25, (A) empirical result; (b) result ofEqu. 5.12 with 4000 s of data and 8 = 0.25, (■)empirical result; (c) result of Equ. 5.12 with 4000 s of data and 8 = 0.5, (•) empirical result.

A 3a x2N level (x2 3g. ) can be taken as 3 X2N s.d.'s above the empirical x2N for no pulsed counts. The number of pulsed counts that are found empirically to be the

151

most likely to produce x 2 3 a is then the 3 a upper limit pulsed counts. For high numbers of pulsed counts the empirical mean x2N is approximately equal to that calculated with Equ. 5.12. In this case Equ. 5.12 can be used to find the approximate statistical significance of a detection provided the empirical s.d. is used. X 2 3 a can be used in the binomial distribution of Equ. 5.17 assuming the probability of x 2 > X 2 3 a is 0.01.

The partially empirical techniques described above were applied to each run of data, chosen to be 4000 s long (see Sec. 5.5). The values of 3 used for the forced counts were adjusted to suit the particular shape of the pulses predicted to be present. For each chosen pulsar 3 consecutive data runs were used, timed to cover the period when the pulsar was highest in the sky. Each run results in a x 2 N vs. period distribution. Use of several independent runs allows a further check to be made on the consistency of the distributions( 13 7).

Unfortunately, long continuous data runs were impracticable with the data available. Thus in order to obtain maximum sensitivity it was necessary to combine the X2n from each separate 4000 s run. In principle if the relative phases of each histogram for a given period in a given run are known, then the histograms can be added and the combined x2N calculated. However, this was not the case. Also, if the relative phases are known and the data obeys Poisson statistics then the x2w values for each

152

data set can be converted into probabilities using statistical tables. The product of the probabilities gives the probability of occurrence of the period in all the independant data sets. However, since the x2N data is not truely Poissonian it was decided to add the x2N foe the same period in each data run and test the significance of the result empirically using the method of forcing pulsed counts into the data.

5.4.3 Interpretation of Pulsed Count RatesThe data analysis procedure outlined in Sec. 5.4.2

will result in values for detected pulsed counts due to a candidate pulsar or an upper limit for them. To interpret either of these results it is desirable to convert the pulsed detector counts into a pulsed flux from the source incident at the Earth, or an upper limit for it. In this way the sensitivity of the technique can be assessed in terms of its sensitivity to particular sources. Source parameters can be obtained for comparison with the results of other observations or theories.

A major drawback of the 6.3 m2 detector, used to find values for the pulsed flux from a source, is the provision of only one spectral energy band of 30 - 2000 keV (see Sec. 5.3). It is therefore necessary to assume some likely spectrum of pulsed events from the candidate source and to convert this flux into a pulsed count rate from the detector ( 151+ 186) . This predicted count rate can

153

then be compared with the actual result. (In detectors with complex response functions this "backwards unfolding" is often required even if good spectral resolution is available (1 5 **)) . The equation adopted to achieve the prediction is:

20 0 0D = / F(E)T(E,d,i|>) e(E)A(\|>)dE (5.19)3 0where D (ct s” 1) is the predicted detector pulsed count rate and F (ph cm”2s” 1keV“ x) is the assumed pulsed photon number flux spectrum. T is the transmission of the atmosphere (a function of detector altitude d and the source altitude angle \|>) , e is the efficiency of the detector at energy E (see Fig. 5.2). A (cm2) is the effective area of the detector for a source at altitude angle obtained-using Fig. 5.3. D will be a slow function of time through the changes in and d during the observation.

5.5 RESULTS AND INTERPRETATIONDuring the 1978 flight of the 6.3 m2 detector almost

the whole southern sky was scanned at least once. Therefore the majority of known X-ray pulsars passed through the field of view. The analysis techniques (see Sec. 5.4) could be applied to examine the data for periodicities due to any of these sources. However, the analysis of 3 isolated X-ray pulsars associated with SNR's, and 3 binary X-ray pulsars, was considered sufficient to illustrate the results achievable with a variety of pulse characteristics and

154

periods. The use of the technique for a general pulsar search could then be assessed. Furthermore, it was considered that the particular sample of sources selected should also be worthy of study from an astrophysical standpoint.

One selection of pulsars that satisfies the requirements is; PSR0531+21, PSR0833+45, and PSR1510-59 (examples of isolated pulsars), and GXl+4, 4U1626-67 and 0A01653-47 (examples of binary pulsars). Discussed here for each source are; i) a review of their hard X-ray pulse characteristics and, ii) the results and implications of the analysis, from both instrumental and astrophysical viewpoints. Pulsed emissions were observed from the Crab and GXl+4. Upper limits were obtained for the remaining sources.

5.5.1 Observations of Isolated X-ray Pulsars Associated with SNR* s

Since the discovery of 33 ms X-ray pulsations from the pulsar in the Crab SNR by Fritz et al. (1969)(187), the system has become the most extensively observed object in the X-ray/y-ray sky(175 183 188). it is one of the most luminous known sources at these energies and the structure of the pulses are unusually similar in all wavebands( 13 170 183). All except the radio pulse profile show a characteristic double peak separated in phase by 0.42 ± 0.05 ( 1 3 3 18). The emission is found also to be very constant with time( 176 183 188). Thus it is perhaps not surprising that the Crab is

155

not only observed by astronomers seeking answers to specific astronomical questions but also for assurances about the functioning and calibration of instrumentation( 189). Hence, as an example of a non-binary pulsar, the Crab seemed a mandatory choice for the present work.

The interest in sources with Crab-like characteristics (see Sec. 2.2.4) meant examples of other such sources were also a natural choice for investigation here. The best studied other examples, are Vela and the newly discovered MSH1510-59. Outlined in this section are the results obtained for the analysis of data for the Crab, Vela and PSR1510-59. (A summary of the general characteristics of isolated X-ray pulsars is given in Ch. 2.)

5.5.1.1 The Crab pulsar PSR0531+21 (NP0532) i) Hard X-ray characteristics of the Crab: Although the pulse profile of the Crab at X-ray energies is remarkably similar to those at other energies (see Fig. 5.8), there are certain unique characteristics: i) There is a particularly strong emission between the two pulses, though the flux drops to zero for the remainder of the cycle as at other energies^120 133 175 188). ii) The secondary peak, defined by phase, is found to be more asymmetric and to dominate the primary(176). iii) The primary pulse becomes 20% narrower in the energy band of about 100-400 keV, though increasing again in neighbouring energy bands(176).

156

CRAB. PSR 0531*21

-<%P-Gamina>2Q0MeV

50 -2Q CMe V

1 -2Q Me V

X - R a y10 0 - ^ O O ke V

!3 -1 63 ke V

’ .5 - I O k e V

Opficat

Radio£.30MHz

Fig. 5.8 Crab pulsar light curves at different photon energies. Schonfelder (1983) (203).

The total time averaged emission spectrum and the time averaged pulsed emission spectrum of the Crab between 10 keV and 1 GeV can both be described adequately by single, but different, power laws^190)- For a compilation of observations above 50 keV Graser et al. (1982)(19°) fit a power law for the pulsed spectrum with a photon number flux index of a = -2.2. However, towards lower energies the spectrum begins to flatten. Power law fits have given spectral indicies ranging from -2.2 (4-200 keV)(182), to -2.0, (20-100 keV)(191). Recent results from HEAO-3(188)(50-500 keV) (see Fig. 5.9) gave:

F (E) = (1.04 ± 0.04) x lO-'M-j-p'o- k v')~2- 4 °-09 (5.20)

where F is in ph cm” 2s“ eV"1.

157

Fig. 5.9 Phase averaged spectrum of the Crab pulsar. Heavy solid line is the best fit of HEAO-3 data to a single power law. Previously measured spectra are also shown. Mahoney et al. (1984) ( 188).

ii) Results: Data from the 6.3 m2 detector for the timeinterval 16h23 to 17h29 UT 16^ Nov. 1978, was analysed forindications of the Crab. Table 5.1 gives details of theobservation and analysis parameters. The predicted pulsarperiod, p, of 33223456.0 ns was obtained using the pulsarparameters provided by Melott et al. (1982) ( 192), for anX-ray observation of the Crab 116 days later. The techniquesdescribed in Sec. 5.4.1 were used. The data was analysedin one continuous run using the epoch folding technique.

158

C r a b V e l a P S R 1 5 1 0 - 5 9 G X 1 + 4 4 U 1 6 2 6 - 6 7 O A O 1 6 5 3 - 4 0

O b s e r v a t i o n s t a r t ( 1 6 / 1 1 / 8 5 ) (UT) 1 6 h 2 3 1 8 h 2 0 0 1 h 0 2 0 4 h 3 5 0 2 h 2 1 0 2 h 0 8O b s e r v a t i o n f i n i s h ( 1 6 / 1 1 / 8 5 ) (UT) 1 7 h 2 9 2 1 h 4 0 0 4 h 2 2 0 5 h 4 2 0 5 h 4 2 0 5 h 4 2M e r i d i a n p a s s a g e o f s o u r c e (UT) 1 6 h 5 3 2 0 h 0 0 0 2 h 4 3 0 4 h 5 4 0 3 h 4 7 0 4 h l 5M i n . z e n i t h a n g l e to s o u r c e (deq) 47 26 40 12 47 30A l t i t u d e of b a l l o o n (mbar) 6 . 3 5 . 6 4 . 0 3 . 6 4 . 0 3 . 6P r e d i c t e d p u l s a r p e r i o d (s) 0 . 0 3 3 2 2 3 4 5 6 0 . 0 8 9 2 4 3 9 0 . 1 5 0 0 4 6 1 1 3 . 9 5 7 . 6 7 7 8 9 3 8 . 2 1 2S p i n d o w n c o r r e c t i o n (us) - 4 . 2 9 2 1 . 1 9 5 - 3 7 . 2 3 2 - 0 . 2 6 s - 6 6 0 - 2 0 0D o p p l e rc o r r e c t i o n (us) - 0 . 9 5 1 - 4 . 2 1 - 9 . 6 9 6 - 0 . 0 1 s - 3 8 0 - 1 0 0No. of p h a s e b i n s 7 8 8 8 8 8P e r i o d s t e p t i m e (us) 0 . 0 4 0 . 7 4 1 . 5 0 . 8 3 0 s 0 . 0 0 7 8 s 0 . 1 9 5 sNo. of e p o c h f o l d s 1 2 0 4 0 0 1 3 4 5 0 0 8 0 0 0 0 35 1 5 6 0 315R a n g e of p e r i o d s ±30 ±10 ±23 ± 113 ± 0 . 0 59s ± 1 . 4 6 s

Table 5.1 Observation and analysis parameters for the 6 pulsars studied.

The 4 ms timing resolution of the detector meant only 7 phase bins could be used. Epoch folding periods of upto =30^5 away from p were employed. This is ample enough to account for any pulsar glitches. Fig. 5.10 shows the resulting xN vs. period distribution. A rise in the distribution can be seen with a peak at a period of 33223380 ns, (only 76 ns from p). Using the method of Koo et al. (1980) ( 193), the period limits which decrease the peak x2n by unity from the maximum, suggest an error in this period of ±50 ns. Therefore the peak x2N is within 2 s.d's of p (the FWHM is ±145 ns). Fig. 5.11 shows the peak region on an expanded scale covering 2 period shifts either side of the peak. (Arrows mark the predicted period and the expected position of x2j minima using Equ. 5. 15) .

30

25 -

20 -

as 15 -CJ

Fig.33200 33210 33220 33230 33240 33250

PERIOD (us)5.10 xN vs. period for the Crab using data for the

time 16h23 to 17h29 UT lfith Nov. 1978.

33222J5 3 3 2 2 3j0 33223.5 33224.0PERIOD (us)

Fig. 5.11 The peak region of xN vs. period from Fig. 5 on an expanded scale.

X2 away from the peak were found to be somewhat non-Poissonian (s.d. = 5), due to the spurious data as

160

discussed in Sec. 5.3. The 3a sensitivity limit is, X 2 3 a

= 24 (c.f. 17 expected for normal statistics with 6 degreesof freedom). The sensitivity of the technique to Crab pulses in the data was tested by forcing rectangular pulses with duty cycle 3 = 0.5 into the data, and epoch folding at periods away from p (see Sec. 5.4.2.2). (Knight et al. (1982) ( 13I+) suggests 3 = 0.4 for pulsed X-ray flux from the Crab. A 3 of 0.5 was considered an adequate approximation here, especially since the increased interpulse emission may increase 3 for energies above about 100 keV). Using the resulting empirical relationship between x2N and the forced detector counts (similar to Fig. 5.7), the x 2 3 a was found to be equivalent to 21 pulsed ct s” 1. (4 ct s”1 more thanwould be predicted with Equ. 5.12 for Poissonian data).Using the X 2 3 a as x Q with Equ. 5.17 yields an expectation probability of random occurance of the distribution of Fig. 5.11 as <10"lt+. The peak x2N is 223a above the background level.

The apparently high level of significance of the x2N peak, and .its proximity to the predicted position, is evidence for a genuine detection of Crab pulses. This is further supported by the shape of the xN vs. period distribution (see Fig. 5.11). The width of the peak is -0.5 us. This is in good agreement with the width of 0.55 us predicted with Equ. 5.16. There is also evidence for small side lobes, and minima at ±APmj_n (see Equ. 5.15). The presence of small extra minima due to the double peak

161

structure of Crab pulses, as suggested by Fig. 5.5, are not seen. This is probably because of the small number of phase bins used and the presence of interpeak emission. At periods much shorter or larger than 33223380 ns the x2N values decline to a low level as expected. However, this level is higher than the expectation value of 6 for the data. This can be explained by the slight non-Poissonian nature of the data.

Analysis of different and shorter lengths of data also produced x2N distributions significantly deviating from that expected if no pulses are present. The size of the peak x2N was found to increase with observation time as expected. In the light of the significances obtained, and for the purposes of the assessment of the technique, it was considered unnecessary to analyse further data runs. It was concluded that the Crab had been detected with a period at the detector of 33223380 ± 50 ns.iii) Interpretation; Fig. 5.12 shows 2 cycles of the pulse phase histogram at the measured period (comparison is made with Wilson et al. (1983)(176), (45-360 keV)). The totalcounts for one cycle in the new histogram are 4.3 * 108.The pulse phase histogram clearly shows the expected two peaks separated in phase by 0.50 ± 0.10. The second peak appears to be higher than the first and is followed by a sharp drop in intensity lasting for -1/4 of the cycle.Between the peaks the drop in count rate is less. This is consistent with previous observations. However, it should

162

be noted that the poor energy resolution of the 6.3 m2 detector means the phase of the epoch folding relative to the Crab pulses may etffect the shape of the pulse phase histogram.

the Crab at the measured period of 33443380 ns, compared with; (b) the profile of Wilson et al. (1983) ( 176).

If we assume that the mean count per bin in the series of phase bins with the lower counts corresponds to the background level when the pulsar is "off", then we can deduce a pulsed count rate from the detector of 165 ± 20 ct s"1. Using Equ. 5.12 and the s.d. of x2N values for folding periods away from p we deduce that the peak x2N is equivalent to pulsed counts of 173 ct s-1. This is in good agreement with the previous result. Thus the pulsed fraction for the observation is 0.15 ± 0.02%.

Fig. 5.13 shows the result of binning the light curve

163

from Wilson et al. (1983)(176), for 45-360 keV, into 7 phase bins and comparing the profile with the present measurement. The pulsed count rates in each profile have been equalized and the phases altered to align the peaks. The two histograms are found to be quite similar, allowing for errors of the present observation. However, there appears to be relatively more counts in the second peak and a higher interpulse emission for the new observation. This may be because the6.3 m2 detector is more sensitive to the interpulse emission which is greater at energies above a few 100 keV.

Fig. 5.13 Pulse phase histograms for the Crab: (a) presentresult (30 - 2000 keV); (fc>) profile adapted fromWilson et al. (1983) ( 176) (45 - 360 keV).

Adjusting the pulsed emission spectrum for the Crab given by Equ. 5.20 (with oi = 2.14), for atmospheric attenuation and the detector response using Equ. 5.19, yields an expected pulsed detector count rate, for 30-2000 keV, of 100 ± 10 ct s"1. The spectrum of Graser et al. (1982) ( 19°) with a = 2.2, yields 90 ± 10 ct s” 1. That of

164

Hameury et al. (1983) ( 191) with a = 2.0, yields 120 ct s"1. Even allowing for a possible under estimate of the error in the measured count rate, due to the non-Poissonian nature of the data, and for the spread of the results above, there is still some disagreement between these count rates and the new value of 165 ± 20 ct s"1. This may suggest some systematic errors.

One likely source of error could be an inaccurate estimate of the detector altitude due to uncertainties in the temperature characteristics of the barometer (see Sec. 5.3). This could result in errors in pressure of up to ±1 mbar. A systematic over estimate of the residual pressure by 1 mbar would increase the pulsed count rate by 30 ct s" 1, assuming the spectrum of Equ. 5.20. Furthermore, it is suggested by Wilson et al. (1983)(176) that atmospheric scattering of photons into such a large area detector could be a significant effect. This would tend to increase the measured count rate relative to that predicted using Equ. 5.19. Therefore the present results are almost certainly consistent with the accepted spectrum of pulsed X-ray flux from the Crab.

The empirically estimated 3a sensitivity of the detector of 21 pulsed ct s”1 from above, for a 4000 s observation of Crab-like pulses, can be used to find the ultimate sensitivity of the detector to a Crab-like pulsed spectrum. If we assume a value for the residual atmosphere to the source of 3 gm cm"2, and an effective detector area

165

of 6.3 m2, then applying these values with Equ. 5.19 and the Crab spectrum yields a predicted count rate of 260 ± 20 ct s” 1. This leads to an ultimate 3a sensitivity of the detector to a Crab-like spectrum, in a 5 h observation, of =50 pulsed mCrab. If the data had been Poissonian, then, using Equ. 5.12, this sensitivity would be =30 pulsed mCrab. Fig. 5.14 compares the sensitivity to that predicted for the FIGARO and OSSE detectors (see Sec. 1.5). The sensitivity of the 6.3 m2 detector compares favorably with these instruments despite the presence of non-statistical noise in the data and the very wide aperture of the 6.3 m2 detector.

ENERGY (MeV)Fig. 5.14 Detector sensitivities for a 5 h observation:

(a) Phase averaged spectrum of the Crab; (b) minimum detectable pulsed flux of FIGARO; (c) minimum detectable pulsed flux of OSSE; (d) minimum detectable pulsed flux of y-burst detector.

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5.5.1.2 Vela PSR0833+45The Vela pulsar is the strongest known y-ray source in

the sky and was the first to be associated with -a "SNR( 19*0 . Pulsations were first discovered in the radio waveband with a period of 89 ms(195) (the third fastest at these wavelengths). Though Vela is more like the Crab than any other source^195) there are some important general differences: Vela is older and 4 times nearer than the Crab and lies in a complex region containing other X-ray sources(189 197 198); the main nebula is of a different class, and there is a small 0.1 pc nebula around the pulsar not seen in the Crab‘(13°); the structure, relative phase, and number of pulses per period, varies with energy(130 133 134 1 9 5 1 9 9 20 0). and there is evidence for time-variability of the emissions( 1 89 1 95 20 1 20 2).i) Hard X-ray characteristics of Vela: At X-ray energies there is only conflicting evidence for pulsed emission. 3a detections from some observers are contradicted by upper limits found by others^ 1 31* 1 97 1 98 20 1 ). Turner et al.(1984)(2°°) have recently been the first to detect pulsed y-rays from Vela in the energy ran-ge 0.3-30 MeV. Fig. 5.15 reproduces their pulse phase histograms. They observed two peaks of emission separated in phase by 0.43 ± 0.03. Pulsed emission is also observed between these peaks but the remainder of the cycle is at background level(185 20°). The ratio of interpulse emission intensity to first peak intensity is -8 times greater in the range 50-300 Mev(200).

167

Fig. 5.15 Light curves for Y-rays from Vela: (a) 50-6000MeV, COS-B; (b) 0.3-30 MeV, Turner; (c) 0.3-1.5 MeV, Turner; (d) 2295 MHz radio light curve.(Turner et al. (1984) ( 20 0)) .

Fig. 5.16 shows the time averaged pulsed emission spectrum from Vela at X-ray/y-ray energies compiled from recent results. Kanbach et al. (1980) ( 195) fit a power law to COS-B data for 50 MeV-6 GeV with an index of a = 1.89 ± 0.06. However, there are indications of a break in the spectrum at a few 100 MeV. Knight et al. (1982) ( 1 310 suggest a best fit power law for 300-3000 MeV of a = 2.0 ± 0.2

168

and for 50-300 MeV of a = 1.77 ± 0.15. The observations of Turner suggest a continuing tendency for the spectrum to bend away from the COS-B power law. They fit a spectrum for 1-30 MeV of:

dN/dE = (3.6 ± 1.3)x|o'^ “ ( 1. 6 ± °-2) ph cm" 2s” xMeV“ 1 (5.21)

Fig. 5.16 The energy distribution of pulsed Y-rays from Vela. (Turner et al. (1984) \ 2 0 0 )). Also shown are upper limits from the 6.3 m2 detector assuming a power law spectrum with a = 1.3 and:(a) 6 = 0.5? (b) 3 = 0.2.

Extrapolations of Turners measurement at 0.3-1.5 MeV back to the optical flux at 3 eV yields a power law with a = 1.08. The lack of a definite X-ray detection suggests a spectral turnover in this region(124 13l+ 2 0 3). The upper limits obtained with HEAO for 15-175 keV suggest that a

169

decrease in a of >0.6 from the COS-B power law must occur between 100 keV and 10 MeV. Theoretical models that are consistent with the observed spectrum have been developed by Fawley (see Turner et al. (1984) (20 °)). These predict a spectral turnover and a plateau below -1 MeV (see Fig. 5.16). ii) Results; Y-ray counts for the observation time 18h20-21h40 UT 16th Nov. 1978 were epoch folded for signs of pulsations from Vela. Table 5.1 summarises the observational and analysis parameters. The predicted period, p, was obtained using radio measurements from Manchester et al. (1983)(128) for data taken 110 days earlier. Fig. 5.17 displays the vs. perioddistributions for 3 4000 s data runs obtained by scanning over 5 period shifts about p.

; (-) 20h34-21h40 (arrows mark predicted positionof peak and minima).The mean and s.d. of the x2N away from p, for each

4000 s data run, were found to be 8.5 ± 4.5. This implies a 3a x2n of 22. The 3a x2N level for the combined 12000 s x2N

170

distribution was found to be 38. Both these x2N values are higher than would be expected ideally, suggesting some residual non-Poissonian noise despite the data smoothing.An empirical relation between simulated pulsed counts due to Vela and the resulting combined x2N was obtained by the method of forcing rectangular pulses into the data (see Sec. 5.4.2.2). A duty cycle of 8 = 0.5 was initially chosen on the tentative assumption that it may be the same as for the Crab X-ray pulses. Using the empirical relation, and the 3a X2n value of 38, yields a 3a sensitivity to Vela of 14 pulsed ct s"1 from the detector (c.f. -10 ct s"1 predictedwith Equ. 5.12 for a continuous 12000 s observation with Poissonian data and the observed background count rate).

From Fig. 5.17 it can be seen that no single x2N rises above the 3a level for any 4000 s data run. This was also true for the combined 12000 s of data. Results obtained using 20 folding bins, rather than 8, yielded even lower sensitivities as expected from the discussion in Sec. 5.4.2.1 (28 ct s” 1 rather than 18 ct s” 1, using 4000 s of data). The x2N distribution with 20 bins was found to be more non-Poissonian than for 8 bins indicating that the amplitude of the additional spurious data may increase with increasing frequency. This conclusion suggests an increased importance for using small numbers of bins with this data.

The lack of any x2N deviating significantly from that expected suggests that pulsed counts from Vela were not detected. This is further supported by the form of the x2N

171

distribution which has no peaks or minima coincident with positions expected for them. Thus a 3a upper limit to the detection of Vela was taken to be 14 ct s” 1. iii) Interpretation: The upper limit can be compared with the various spectra and upper limits to the pulsed emission by multiplying the spectra through the detector response using Equ. 5.19. An extrapolation of Equ. 5.21 with a = 1.6f produces an expected count rate from the detector for 30-2000 MeV of 29 ± 10 pulsed ct s-1. This value will be subject to possible systematic errors due to the inaccuracies in altitude information (see Sec. 5.5.1.1).In this case an over estimate of the residual pressure by 1 mbar would increase the detected counts by - 6 ct s"1. Even allowing for this error the predicted count rate is still inconsistent with the present upper limit and therefore agrees with the results of Turner et al. (1984)(20°). Their flux measurement for 0.3-1.5 MeV is also below the a = 1.6 power law. Thus the possibility is confirmed that a power law spectrum in this energy domain must have a < 1.6.

Applying the power law with a = 1.08(200) yields an expected detector count rate of -9 ct s"1. This is consistent with the present upper limit. Spectra with a much less than this would have flux levels in the range 0.3-1.5 MeV that are inconsistent with the result of Turner, thus it is likely that a power law in the range 30-2000 keV will have 1 < a < 1.6. For instance a = 1.3 could beconsistent with both results.

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The spectrum of Fawley with a plateau below 1 Mev(200) (see Fig. 5.16) would give a pulsed count rate at the detector of -9 ct s-1. This is consistent with the present measurement. Alternatively if we assume a spectral form which follows the a = 1.6 power law down to an energy below which there is a plateau, then for the upper limit of 14 ct s’ 1 to be consistent, the spectral break must occur above 400 keV. However, to be consistent with Turner the break must occur above about 900 keV.

The assumption has been made that the pulsar duty cycle for Vela is 3 = 0.5. However, in the optical energy band 3 - 0.4 and in the high energy y-ray band 3 = 0.2(131*). Using the latter value the empirical upper limits to the pulsed counts reduce to -8 ct s"1. This suggests an even flafter spectrum for Vela in the range 30-2000 MeV. The spectra of Fawley and the power law with a = 1.08 are still consistent with this upper limit, allowing for the errors involved. A spectral turnover and plateau would require a break at E >900 keV. Table 5.2 summarises the upper limit fluxes obtained for various values of a and 3 assuming a power law spectrum. The values for 3 = 0.5 and 0.2 with a power law with a = 1.3 are also plotted in Fig. 5.16.

In conclusion it appears that despite the problem of non-statistical noise and possible systematic errors, the6.3 m2 detector has successfully provided upper limits for the pulsed emission of Vela. They are consistent with

173

previous results and confirm that the spectrum must be flatter in the 30-2000 MeV range than at higher energies.

F (ph cm-2 s"1)

a = 00o»1—I 1111a = 1.30 ii

iia = 1.60

3 = 0.5 1.6 X 10“ 31111 1.9 x 10“ 3

Iiii

2.3 x 10“ 3

3 = 0.2 8.1

J"1or—1 X 1111in•a> x 10-4

iiii1.2

ro1Ot—iX

Table 5.2 3a upper limit pulsed flux F from Vela (30-2000 keV) for power law fits dN/dE = KE”a with two pulse duty cycles 3.

5.5.1.3 PSR1510-59A compact X-ray source, now designated PSR1510-59, was

first associated with the SNR MSH15-52 by Cowley et al.(1972) (2QLf) . Seward et al. (1982) (205) discovered that theX-ray emission was pulsed at the short period of 150 ms.The same period was subsequently observed at radio and optical wavelengths( 1 32). A possible Y-ray identification has not been confirmed, and the relative phase of the pulses in each waveband is unknown(205). The source is observed to have a smooth and high dp/dt(178), (the highest for any radio pulsar( 1 78 205)) and a high X-ray luminosity. This indicates that the pulsar is not an accreting binary( 132 205). The presence of a luminous nebula around the pulsar, and a hard featureless spectrum, indicates synchrotron emission powered by the pulsar spin down( 1 3 0 1 70 1 78 2 0 5). Thus PSR1510-59 appears to be a Crab-like system. However, like Vela, it lies in a complex emission region and is not as luminous as the Crab(206). The small 0.1 pc nebulosity

174

present around Vela is not observed around PSR1510-59.i) Hard X-ray characteristics of PSR1510-59: To date hardX-ray pulses have not been observed, but emission in the energy range 0.2-4 keV is found to be >80% pulsed, with one pulse per period of FWHM covering 0.25 of the cycle(135 178) (c.f. 0.1 for the radio pulse). This is quitedissimilar to the Crab or Vela at these energies.

The time averaged spectrum of the pulsed emission in the X-ray band is even less well defined than the pulse profile. Seward et al. (1984)(17°) suggest power law fits in the energy range 0.2-4 keV with a = 0.5-1.0. This index must steepen at higher energies to maintain a plausible total luminosity. Results from ARIEL V in the energy range 2-10 keV, taken before pulsations were discovered, suggest a = 1.7 though this includes some of the diffuse emission^135 205).ii) Results: Data for the period 01h02-04h22 UT 16fch Nov.1978 was analysed for pulses from PSR1510-59. Table 5.1 summarises the observational and analysis parameters. The predicted pulsar period of p = 0.150046 s was calculated using parameters obtained from X-ray observations by Weisskopf et al. (1983) ( 1 78), 280 days after the present observation. Fig. 5.18 shows the xN vs. period distribution obtained for each of the 3 separate 4000 s data runs (calculated as described in Sec. 5.4.3.2 and applied as for the analysis of Vela).

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150030 150040 150050 150060PERIOD (us)

Fig. 5.18 xN vs. period for PSR1510-59 with 8 phase bins for observations at: (O) 01h02-02h08; (•) 02h08-03hl5; (-) 03hl5-04h22 (arrows mark predictedposition of peak and minima).

The 3a x2N level for each 4000 s run and for the combined x 2 * T distribution, were found to be 19 and 36 respectively. Data spuriosities were removed in the usual way (see Sec. 5.4.2), and allowance was made for the residual non-Poissonian behavior. Forcing rectangular pulses into the data and epoch folding at periods away from p yielded a 3a sensitivity to pulses from PSR1510-59 of 9 pulsed ct s” 1, for the complete observation. A duty cycle of 3 = 0.25 was chosen to be consistent with that observed for the low energy X-ray pulses. (A duty cycle of 3 = 0.5, like that of the Crab at hard X-ray energies, gives 15 ct s"1, c.f. 10 ct s”1 for Poissonian data). It can be seen from Fig. 5.18 that no x2N deviates significantly from that expected if no pulses are present, nor is any x 2 N for a

176

given period consistently above average in each sepetrate 4000 s observation. This was also true for the combined x2N distribution. Therefore a null detection of PSR1510-59 with an upper limit of 9 pulsed detector ct s"1 was assumed,iii) Interpretation: Table 5.3 displays upper limits to the pulsed flux incident at the earth from PSR1510-59 in the range 30-2000 keV. They were obtained using Equ. 5.19 and assuming power law spectra with the various photon number spectral indeed . discussed earlier. It is found that power law spectra that are consistent with the detector upper limits for 3 = 0.25, only yield luminosities in the range 0.2-4 keV consistent with those of Seward et al.(1982)(205) if a > 1.5. However, this extrapolation is only tentative.- Its validity will depend on factors such as the Hydrogen column. If we assume a spectral form for PSR1510-59 like the Crab, and allow for the larger distance to PSR1510-59, then the present upper limit to the detector pulsed counts corresponds to a flux from the source equivalent to -200 pulsed mCrab (3 = 0.25) or -340 pulsed mCrab (3 = 0.5).

F (ph s"cm 2M

a = 1.0 iiiia = 1 . 5 iiii

a = 1 . 7 iiiia = 2 . 2

( C r a b - l i k e )

3 = 0. 25 8 . 7 x l O " 4iiii 1 . 1 x 10~3

iiii 1 . 3 x I Q " 3iiii 1 . 7 x 10“ 3

3 = 0. 5 1 . 5 x 10~3iiii 1 . 9 x 10“ 3

iiii 2 . 1 x I Q " 3iiii 2 . 8 x I Q " 3

Table 5.3 3a upper limit pulsed flux F from PSR1519-59 (30-2000keV) for power law fits dN/dE = KE_ot with two pulse duty cycles 3.

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5.5.2 Observations of Accretion Powered Binary X-Ray PulsarsA summary of the general characteristics of binary X-ray

sources is given in Ch. 2. Discussed below are results of the analysis of data from the 6.3 m2 detector for GXl+4, 4U1626-67 and OAO1653-40. The unusual characteristics of GXl+4 make it physically one of the most interesting X-ray pulsars (see Sec. 5.5.2.1). More importantly, the strength of the pulsed flux and the distinctive sinusoidal pulse profile indicate that GXl+4 is an ideal choice for assessing the use of a wide aperture detector for studies of X-ray binary pulsars, particularly since the pulse characteristics contrast strongly with those of the isolated pulsars. For instance the pulse period of GXl+4 is about 300 times that of the Crab (see Sec. 5.5.1).

4U1626-67 and OAO1653-40 have recently been shown to share some characteristics in common with GXl+4, as well as interesting and unique properties of their own. More importantly, both sources are suitable for testing the proposed pulsar search technique. Though, like GXl+4, their periods are quite long (see Sec. 2.2 and Table 2.1) they are not long enough to result in inaccuracies due to a small number of folding cycles, or so long that their luminosities are small (see Table 2.1). However, their periods are long enough to be about 2 orders of magnitude different from the periods of the isolated pulsars.

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5.5.2.1 GXl+4It was not until 1970 that Lewin et al. (1971) ( 207)

used a proportional counter with a field of view small enough to confirm the presence, in the complex alactic centre region, of the discrete hard X-ray source GXl+4. The flux from GXl+4 was observed to be regularly modulated at a period of -2 min and showed many of the signs of accretion on to a magnetized neutron star (see Sec. 2.2). Now known to be the brightest X-ray source in the region at energies >20 keV, GXl+4 is found to exhibit unusual spectral variability, with reversals of the relative phase of the pulses in successive soft X-ray energy bands, and a high pulsed fraction. An unusually high spin up indicates an evolutionary period of only -50 y(193). No radio emission has yet been observed, and the possible Y-ray detection by COS-3(208) is not confirmed. However, GXl+4 has been associated with the low energy X-ray source 3U1728-24, and an infra-red source associated with a rare symbiotic star of 19m (209).i) Hard X-ray characteristics of GX 1+4: Many of the X-ray characteristics of the pulsed emission from GXl+4 are still uncertain. Although the originally measured period has been confirmed by most observers, there have been some indications that the period is actually 4 min with 2 pulses 0.5 apart in phase(5). Some observers have detected no pulses at all(20 8 2 1 0 2 1 *) . However, recent analysis of HEAO data suggests the shorter period seems the more

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probable(5). The period derivative is found to be the largest of any pulsar except 3A0535+26, but varies by 2-5% per year, and is correlated to the X-ray luminosity( 17k).The pulse profiles also appear to be highly variable with both energy and epoch, though predominantly for energies below 20 keV(5 210 212). For 25-60 keV the profile from HEAO-1 A2 shows nearly sinusoidal behaviot with the peak defined to be at phase 0(5). The pulse profile obtained with ARIEL V in April 1980 also shows a sinusoidal form above 25 keV but with a second smaller peak at phase 0.8(210).

Earlier pulse profiles observed with SAS-3 by Dotty et al (1981)(17tO in Feb. 1976 for E > 25 keV had peaks at the same phases but with a behaviour less sinusoidal than observed with ARIEL V. The second peak dominating the first by a factor of 2. Dotty observed a constant profile for energies above 8 keV. However, the results of Strickman et al. (1980)(208 213) for 20-70 keV obtained 3 and 21 months later indicate a systematic difference in intensity and pulse shape between odd/even pulses. This suggested a 4 min period and 8 = 0.4, though considerable non-periodiceffects were also observed. Non-periodic effects seen by Dotty, and Ricketts et al. (1982)(21°), were sufficient to account for any difference between successive pulses. An observation by Koo et al. (1980) ( 193) in April 1974 for 20-60 keV, also indicated a 4 min period, with 8 = 0.29 ± 0.10,

The intensity and spectral shape of the total X-ray emission from GX1+4 (pulsed plus non-pulsed), also shows

considerable variability with both time and pulse phase.The complexity of the X-ray spectrum prevents unique fits by power law, exponential, or blackbody spectra, with or without high energy cut-offs^5 208) (see Sec. 1.2.1). However, over limited energy ranges, such spectral fits have been proposed. At energies >20 keV, power laws with a = 2.4-4.l(207 21 1 2 1I+)/ or temperatures of kT = 18-42 kev(17l+ 20 7 20 8 21l+ 215), have been suggested. Ricketts et al. (1982)(21°) fit a multicomponent blackbody spectrum for 1-50 keV. The variety of fits perhaps reflects the long term variability of GXl+4(208). Pulse phase spectroscopy has revealed* that different spectral fits and intensities are required at different phases(5 208 21°).

Dotty et al. (1981) (17tO suggest that the total X-ray emission may be a mixture of radiation from different processes^ in the system. The relative similarities of the pulse shape for different energy bands above 8 keV at any particular time implies that the modulation process at these energies may be independent of energy. Therefore the pulsed component of the emission can be treated separately. Dotty fits a blackbody spectrum with kT = 8.5 ± 1 keV to the pulsed component, yielding a luminosity of -5 x 10 37 erg s"1 for 8-55 keV. An exponential spectrum with kT = 34 keV can also be fitted if a large degree of Hydrogen absorption is allowed. Koo et al. (1980) ( 1 9 3) have recently suggested that a fit to the pulsed spectrum could have the same slope as the total time averaged emission spectrum at hard X-ray

181

energies indicating power law fits with a = 2.6 for 20-64 keV. This is in agreement with Lewin et al. (1971)(207). They suggest a = 2.6 ± 0.2, but also find kT = 18 ± 2 gives an adequate fit.

Mechanisms to explain the X-ray characteristics of GX1+4 are unclear. Dotty et al. (1981) (17I+) suggest that the average spectrum is consistent with emission from a hot spot of size 2 x 10~5 dj^c2 km2, where d^Dc is the distance to GX1+4 in kpc. The more recent measurements, though indicating no underlying change in the mechanisms, are more consistent with a fan beam interpretation,ii) Results: Data from the 6.3 m2 detector for the time interval 04h35 to 05h42 UT Nov. 1978, was epoch foldedmodulo a range of periods about a predicted pulse period for GX1+4 of p = 113.95 s. This period was calculated using data from Ricketts et al. (1982) (21°), for an observation 235 days later, and Strickman et al. (1980)(213), for an observation 357 days before the present one. A period derivative consistent with both these observations was used. Table 5.1 shows the observation and analysis parameters.

Fig. 5.19 shows the x^ vs. period distribution for a continuous 4000 s run of data covering 4 period shifts (4 times the predicted width of the x2N peak). Arrows mark the expected period and position of the x2N minima calculated with Equ. 5.15. A peak in the distribution can be seen at a period of 114 s. Using the method of Koo et al.(1980)(193) (see Sec. 5.5.1.1) the error in this period is

182

±0.1 s (FWHM = 4 s). The position of the peak is therefore consistent with that expected for pulses from GXl+4.

105 110 115 120 125PERIOD (s)

Fig. 5.19 xN vs. period for GX1+4 with 8 phase bins for an observation at 04h35 to 05h42. (Arrows mark the expected period and minima).

X2n values for epoch folding periods away from p were found to be considerably non-Poissonian (mean of 15, s.d. of 9). This is possibly a consequence of the small number of folding cycles achieved (see Table 5.1). The resulting effective 3a x2N sensitivity limit of 43 (c.f. 18.5 forstatistically normal data) indicates that the x2N peak is 20 s.d.'s above the mean. Applying the calculated s.d. in Equ. 5.17, and using independent x2N values, yields an expectation probability for random occurrence of the distribution of <2 x 10"8. An empirical value for the 3a sensitivity of the data to rectangular pulses with duty cycle 3 = 0.5, was found to be 25 pulsed detector ct s"1.

183

iii) Interpretation: The apparent significant deviation of the x 2n distribution from that expected for non-pulsed data, and the general shape of the distribution, are clear indications of a genuine detection of GX1+4. As for the Crab data (see Sec. 5.5.1.1), there is evidence for side lobes in the expected places. Further evidence for a detection is provided by the vs. period distributionobtained with 20 phase bins for a shorter observation of 3200 s. Fig. 5.20 compares this distribution to the "background" distribution obtained when GX1+4 was below the horizon. A significant peak is only apparent when the source is in view. The peak is in the same place as for the 8 bin data, but is of lower significance as expected (14.5 s.d.'s from the mean) (see Sec. 5.4.2.2).

view; (b) with GX1+4 below the horizon.The period measured for GX1+4 fits well into the trend

of decreasing period with time. Fig. 5.21 shows the period in comparison with those measured by other observers.

184

Fig. 5.22 shows 2 cycles of the folded pulse phase histogram obtained for the period of 114 s using 20 phase bins. The total number of counts in the data for 1 cycle of this pulse phase histogram is 3.86 x 108. The pulse profile shows an almost sinusoidal form with one peak per cycle, as found at lower energies by other observers. It agrees well with that of White et al. (1983)(5) for 25-60 keV. There is no evidence for a precursor at phase 0.8. The duty cycle, defined as the proportion of bins with counts greater than the mean per bin, is 0.5.

Fig. 5.21 The pulse period of GX1+4 for 1970 to 1980. (from Ricketts et al. (1982) ( 2 1 0)). (a) marks theresult from the 6.3 m2 detector.

PHASEFig. 5.22 Pulse phase histogram of GX1+4 for p = 114s.

185

If we assume that the mean counts per bin for the 1/3 cycle covering phase bins 15-20 (see Fig. 5.22) is a good approximate value for the non-pulsed contribution to the count in each bin, then the total pulsed contribution in the bins outside this region above the mean is 72 ± 8 pulsed detector ct s-1. (The error quoted reflects the variance of the counts in the bins chosen to calculate the mean). This result is in good agreement with the value of 67 pulsed detector ct s”1 obtained using Equ. 5.12 with the peak X2 / a duty cycle of 0.5, and the measured x2N s.d.The measured pulse fraction is 0.075 ± 0.007%.

Table 5.4 gives the total pulsed flux from GXl+4 (obtained using Equ. 5.19) incident at the detector in the energy range 30-2000 keV for the various spectra suggested earlier. Fig. 5.23 shows a plot of some of these spectra.The best fit blackbody spectrum has a temperature of 10 ±1.0 keV, assuming a hot spot of the same size as suggested by Dotty et al. (1981) (171+) . Their value of 8.5 ± 1.0 keV for 8-55 keV is in reasonable agreement with this. The slight discrepancy could be accounted for by the systematic errors involved in the calculation (see Sec. 5.5.1.1). The suggestion by Dotty of an exponential spectrum with kT = 34 keV yields a luminosity in this energy range 2.5 times lower than they measure. This may be expected since no allowance has been made here for Hydrogen absorption. Power law fits will only give a luminosity comparable with those of Dotty

186

for photon number indacgs >3.B B ( 1 0 . 5 keV) T B ( 1 8 * 5 keV) T B ( 3 4 . 0 keV) a - 2 . 2

( C r a b - l i k e )a » 2 . 6 a - 2 . 9

P h o t o n N o . f l u x ( 3 0 - 2 0 0 0 keV) (ph c m " 2s “ *)

1 . 8 x 1 0 - 2 2 . 4 * 1 0 - 2 1 . 7 x 1 0 - 2 1 . 1 * 10-2 1 . 3 x 1 0 “ 2 1 . 5 x 1 0 - 2

E n e r g y f l u x ( 8 - 5 5 k eV) 2 . 0 x 1 0 37 5 . 3 * 1 0 37 1 . 9 x 10 37 1 . 3 x i o 37 2 . 5 x 1 0 37 4 . 0 x I Q 37( erg s " 1)_______ __________ 1__________ 1 ___________________ 1

Table 5.4 Pulsed flux from GX1+4 (±10%) for fits to ablackbody spectrum (BB), thermal bremstrahlung spectra (TB), and power law spectra dN/dE = kE“a, with 3 = 0.5.

Fig. 5.23 Time averaged pulsed emission spectra of GX1+4 that are consistent with the results from the 6.3 m2 detector: (a) power law with a = 2.2; (b)blackbody with kT = 10.5 keV; (c) thermalbremstrahlung with kT = 34.0 keV.

187

5.5.2.2 4U1626-67Pulsed emission from 4U1626-67 was first discovered

when the soft X-ray experiment on SAS-3 performed a pulsar search using the technique of studying sources with known hard X-ray spectra (see Sec. 5.2)(166). Pulses were later confirmed in the optical(216). Like GX1+4 the object has an X-ray luminosity >1037 erg s"1 and has now been shown to display considerable variability of spectral shape and intensity with both pulse phase and time, as well as reversals of the relative phases of pulses between successive low energy X-ray bands(217). Except for the occasional glitch^177) the pulse period of -7.7 s is remarkably stable(216). It fills a gap in the distribution of pulsar periods. Like GX1+4, the object is atypical of optically identified X-ray pulsars since the primary is not a luminous OB star^218). However, an observed period of41.5 min(218), confirms that the system is a compact lowmass binary involving accretion onto a neutron star(216 218 2 1 9 220) i) Hard X-ray characteristics of 4U1626-67: As yet there have been no X-ray observations of the pulsed emission from 4U1626-67 above about 50 keV. In the range 1-20 keV Eisner et al. (1983)(177) deduced a constant pulsed fraction of f = 0.17 ± 0.01. This agrees quite well with the result of Rappaport et al. (1977) ( 166) of f - 0.15 for E > 12 keV.However, the pulse shape has been observed to vary on time scales of minutes^177 218). Above 14 keV the pulse profile

188

is found to be nearly sinusoidal with one peak per cycle(5 166 177 217 218).

Unlike GX1+4, the total average X-ray spectrum (pulsed plus non-pulsed) drops off sharply above 20 keV and can be fitted by a continuum spectrum with a high energy cut-off (see Sec. 1.2.1). White et al. (1983)(5) fit a power law with a photon number flux index of a = 1.4 below 20 keV.This is in agreement with the results of Pravdo et al.(1979) (2 17) and Eisner et al. (1983) ( 177). Above 20 keVWhite uses a cut-off with Ec = 20 keV, and Ep = 6 keV. A blackbody fit with the same cut-off has a temperature of kT = 0.65 keV. However, there are sudden and complex departures from this mean spectrum as a function of phased5 ill 217) an£ other changes on time scales of minutes. These are not correlated with the observed intensity changes, which can be by factors of up to 2 below 50 keV on time scales of hours to months^177 217). Quasi-periodic flares, unique to 4U1626-67, occur at a period of about 1000 s(219).

Explanations of some of the properties of 4U1626-67 have been given by Joss et al. (1978) ( 22°). The blackbody fit by White et al. ( 1 9 8 3 ) suggests a hot spot with area1.1 x 1012 cm2 if the star is 6 kpc distant,ii) Results; Data for the period 02h21 to 05h42 16^ Nov. 1978 was analysed for pulses from 4U1626-67. Table 5.1 summarises the observational and analysis parameters. The predicted pulsar period of p = 7.67789 s was calculated using the period given by McClintock et al. (1980)(218) for

189

an optical observation -170 days after our flight. A period derivative from Peterson et al. (1980) C 2 1 6 ) of 4.95 x 10-11s s was also used. The x2N vs. period distribution for 12000 s of data, was obtained by combining the distributions resulting from epoch folding 3 separate 4000 s data runs (see Sec. 5.4.3.2) and gave a 3a x2N level of 32. Account was taken of the non-Poissonian nature of the distribution.

No x 2n was found to deviate significantly from the values expected if no pulsations are present. This was also true for individual 4000 s data runs. Thus a null detection of 4U1626-67 was assumed. The procedure of forcing rectangular pulses into the data using a duty cycle of 0.5 yielded an overall 3a sensitivity of 16 pulsed ct s” 1.iii) Interpretation: If we assume, as is possible for GX1+4, that the spectrum of the pulsed emission from 4U1626-67 lies parallel to that of the time averaged total emission, then the power law spectrum suggested by Pravdo et al.(1979) (2 17) with a = 1.4 and a cut-off at 20 keV, may be appropriate. A flux of this form incident on the detector, normalized to give the calculated upper limit detector pulsed counts, yields an upper limit photon number flux in the range 30-2000 keV of 0.021 ph cm“2s_1, though this is subject to possible systematic errors (see Sec. 5.5.1.1).The same spectrum results in an upper limit pulsed energy flux in the range 0.5-60 keV of 2.2 x 10 38 erg s" l.. This is consistent with the total energy flux measured by Pravdo, which is 18 times Sma/tfr in this energy range. If we assume

190

that the pulsed fraction of =15% given by Eisner et al.(1983) ( 1 77) and Rappaport et al. (1977) ( 166) can beapplied, then the expected pulsed X-ray energy flux is perhaps 2 orders of magnitude harder than suggested by Pravdo.

Power law spectra consistent with the upper limit can only provide an energy flux as high as that given by Pravdo et al. (1979)(217), assuming a 15% pulsed fraction, if the photon number flux index is a> 2.1. Table 5.5 displays the upper limit photon number fluxes for various power laws.If, as is possible for GX1+4, the pulsed components obey a blackbody law, then assuming the hot spot has an area as given by White et al. (1983) (5), the temperature must be <3.5 keV to be consistent with the upper limit. Thus the value of 0.65 keV suggested by White for the total emission spectrum is consistent, whether or not a cut-off is included. Blackbody fits with higher temperatures would require high energy cut-offs.

a = 2.2 (Crab-like)

iiiia = 2.4 iiii

ct = 2.6 iiiia = 2.9

Photon No. flux ii ii ii(30-2000 keV) (ph cm”2s“ 1)3.8 x 10“3 iiii

5.1 x 10“3 iiii5.5 x 10“ 3 iiii

5.6 x 10“3Energy flux (0.5-60 keV) (erg s"1)

6.6 x i o 36iiiiii

1.6 x 1 0 37iiiiii3.3 x 1037

iiiiii9.6 x i o 37

Table 5.5 3a upper limit pulsed flux from 4U1626-67 assuming power law spectra dN/dE = KE”a with pulse duty cycle 6= 0.5

5.5.2.3 0A01653-4QPolidan et al. (1978)(221) discovered a new X-ray

source, designated OAO1653-40, in a crowded region of the

191

galactic plane, using COPERNICUS. 5 months later an • observation with HEAO-1 on 4t 1 Sept. 1978 revealed a 38 s period in the emissions, the first pulsar period to be found in the 10-100 s ranged175). The period derivative is as yet unknown. Like GX1+4 and 4U1653-67, OAO1653-40 has been confirmed to have a luminosity higher than would be deduced from its long period(5 175) (see Table 2.1) . The light curve of OAO1653-40 shows evidence for considerable long term variability especially at low X-ray energies(5). No optical identification has been achieved though 2 very red stars in the region are candidates(126 167). Details of mechanisms to explain the system are not developed, however, the secondary is thought to be a compact object. No binary behavior has been observed.i) Hard X-ray characteristics of 0A01653-4Q: Above 8 keV the X-ray pulse profile is sinusoidal with no evidence of an interpulse(126). The pulsed fraction appears to increase with energy^175). A value of 30% for the peak to mean amplitude in the range 8-40 keV was observed by White et al. (1979) ( 126). Data from HEAO-1 indicate a constant pulse shape in the range 13-80 keV.

The total intensity of the X-ray emission from OAO1653-40 shows some variability on time scales of days to months, however, White et al. (1979) ( 126) observed no variability in amplitude or period of the pulses during their 8 h observation. Unlike GX1+4, there appears to be little variation in the shape of the total emission spectra

192

as a function of phase. The phase average total emission spectrum observed by White can be fitted by a power law with photon number index a = 1.2 in the range 1.5-30 keV. As for GX1+4 the high energy cut-off was not obvious, however the best fit power law up to 60 keV gives a = 1.4 and indicates some steepening above 20 keV. Further work by White(5) suggests a gradual cut-off to this power law with Ec = 18 and Ep = 26, whilst an exponential fit gives a temperature of kT > 43 keV. For the energy range 13-189 keV, Byrne et al. (1981) ( 175) fit a power law with the steeper index of a = 2.7 ± 0.4, or kT = 26 ± 3. They suggest that this is consistent with the results of White provided there is a cut-off between 20 keV and 80 keV.Byrne fits an exponential law to the time averaged component alone with kT = 44 keV, and a power law in the range 13-80 keV given by:

II- = (2.3 ± 0.2) x 10-1* (E/30) "( 2- 0 ± 0.3) ph cm”2s” xkev-1dE (5.22)

however, it is concluded that since the pulse fraction increases with energy to =50% for 40-80 keV, this spectrum must steepen with energy like the total emission spectrum,ii) Results: y-ray events for the time interval 02h08 to 05h42 UT 16^ Nov. 1978 were epoch folded for signs of pulses from OAO1653-40 (see Table 5.1). The predicted period of p = 38.212 ± 0.005 s was obtained using data from Byrne et al. (1981) 175) from an observation 73 days before our flight. The spin up correction was estimated using a

193

theoretical prediction of dp/dt by Armstrong et al.(1980) C 1 6 7 ) . Both this correction and the Doppler corrections are smaller than the precision to which the period is known. This uncertainty is well covered by the width of the period scan used.

Despite the smoothing operation the data for this observation was found to be particularly non-Poissonian.The experimental 3a x2N for the combined 12000 s observation has the high value of 53. This is found empirically to be equivalent to 18 pulsed detector ct s“ 1 assuming a rectangular pulse profile with duty cycle 3 = 0.5. No individual x 2 N value -was found to deviate significantly from the mean for the combined or separate data runs. Therefore an upper limit to the pulsed counts from OAO1653-40 was fixed at 18 ct s“ l. In view of the nature of the data this upper limit is somewhat tentative.iii) Interpretation: Folding the power law spectrum of Equ. 5.22 through the detector response gives an expected detector pulsed count rate of 48 ct s"1. Even allowing for the errors in Equ. 5.22 this is inconsistent with the present upper limit. This may confirm the possibility of a spectral cut-off in the range 20-80 keV.

For a power law photon number spectrum to be consistent with the upper limit and with the pulsed flux in the range 40-80 keV given by Byrne et al. (1981)( 1 75)f a must be greater than 3.0. Such an index would be consistent with a cut-off at -50 keV. An exponential spectrum with kT - 44 keV,

194

as suggested by Byrne, would yield -22 pulsed ct s” 1.This is just consistent with the upper limit provided by the6.3 m 2 detector suggesting that an extrapolation above 20 keV of the power law fits to the low energy data is incorrect.The spectrum of pulsed counts from OAO1653-40 must steepen above 20 keV.

5.6 CONCLUSIONIt is a unique characteristic of X-ray sources that

their X-ray flux is so often modulated with a regular period. The analysis of data from the 6.3 m2 y-burst detector has shown the possible viability of using a wide aperture instrument to exploit this fact. The wide aperture of the detector allows the possibility of observing nearly all the X-ray pulsars in about half of sky in 24 h. Specifically, it has succeeded in measuring the pulsed X-ray flux from the Crab and GX1+4 and providing upper limits for Vela, PSR1510-59, 4U1626-67 and OAO1653-40. This was achieved despite the presence of non-statistical noise in the data.

The data has highlighted the fact that the pulsed emission spectra of many X-ray pulsars steepen above -20 keV. The estimated sensitivity of the detector of =50 pulsed mCrab (3a) for a 5 h observation is comparable to the sensitivities of the planned OSSE and FIGARO X-ray detectors. Both of these have fields of view less than half that of the 6.3 m 2 detector and use more efficient detection

195

media.It is possible to search in the data from the 1978

flight of 6.3 m 2 detector for periodicity due to any other pulsar that passed through the field of view, known or unknown. This could best be achieved with more sophisticated program techniques than were employed here.For instance, the technique of "Fast folding" could be used. This could considerably reduce the computing time involved. However, a purpose built wide aperture instrument would certainly be more successful. Such an instrument could be provided with moderate energy resolution, perhaps by using an array of Nal crystals. It seems likely that a future instrument with such capabilities could provide useful information on the pulsed X-ray flux from X-ray pulsars. It could also double as a y-burst detector.

196

ACKNOWLEDGEMENTS

I am greatly indebted to my supervisor for his wise guidance and many "quick minutes" of discussion and advice, and to Professor P.C. Hedgecock for the opportunity of working in the Cosmic Rays and Space Physics group (now part of Astrophysics) at Imperial College.

I wish to thank; Dr. John Quenby and Dr. Geoff Rochester for their support, Dr. Tim Sumner and Dr. Phil Harper for their advice and tolerance of my computing questions, Arther Bewick for his electronics skills and advice, Dr. John Greenhill for some useful discussions, and John Spragg for providing me with the ARIEL VI detectors.

For their friendship and encouragement I thank Manu Joshi, Pete Smith, Dr. Kevin Beurle, and the rest of the CRSP/Astrophysics group, and Nick Haskell. Special thanks to Kevin for his moral support during the ill-fated Sicily campaign and for some interseting discussions about scientific colaborations. Thanks also to Dr. Rob Lewis.

I am also indebted to; Dr. Peter Sanford for looking after me at UCL, Dr. John Leake and his staff for looking after me at Harwell, and the staff of SIL Ltd., and Dr. D. Vencatesen for looking after me in Saskatoon.

Thanks also to; my tutor at Southampton, Bill Myers, for his support, Dr. Tony Dean for his enthusiasm for hard X-ray astronomy, Steve Shadbolt for his tolerance and frienship while working with me on our 3rc year undergraduate project in X-ray astronomy, and Chris Collins for his companionship and knowledge for physics.

Finally, thanks to the SERC for the provision of a grant and to my parents for their financial support throughout my years as a student.

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