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Produced by the NASA Center for Aerospace Information (CASI)
HiNAL REPORT OF RESEARCH

Performed Under NASA Grant Ao. JAU-J-14

DEVELOPI'EAT

OF A flERCURIC IODIDE SOLID

STATE SPECTROMETER FOR X-RAY ASTRO^!OI Y

September 1983

(NASA-CR-175.130) DEYELOi^MtNT CF A MERLUBIC

IUDIDE SOLID STATL SYEC;TRCMLTEb FOR X-hAY

ASTRUNOMi Final Report (MdSSdChUSettS Inst.

of Tech.) 1.19 p hC A 07 /MF A01 CSCL 03A

N84-16096

Uuclas

G3/89 15150
FINAL REPORT OF RESEARCH

Performed Under NASA Grant. No. NAGW-14

DEVELOPMENT

OF A MERCURIC IODIDE SOLID

STATE SPECTROMETER FOR X-RAY ASTRONOMY

Prepared for

NASA Headquarters

By

Dr. John B. Vallerga

For

Dr. George R. Ricker

t

. u

-w-
Since its beginning in 1962 with the discovery of Sco X-1, the

science of X-ray astronor.Lv has progressed by placing X-ray photon

detectors above the atmosphere in rockets, balloons and satellites.

This usually involves modifying the detectors to withstand extreme

environments and payload constraints (size and weight). The goal has

always been to increase the sensitivity of the detect3r3 (signal to

background ratio) either by increasing the effective collecting area

and/or observing times while decreasing the background counting rate.

Each new detector technology th p ' iz applied to X-ray astronomy must be

reoptimized to achieve thest goals using wi.deiy different strategies.

.n the '.ast few years, the balloon lab at the Center fo , Space

Research at MIT has been studying mercuric iodide (HgI.,), a solid state

room temperature X-ray detector, for use in hard X-ray astronomy. its

primary advantages are its good energy resolution at room temperatures

and a large photon cross section. Much work has been done to characterize the properties of HgI 2 that affect the operation of these

detectors for use in X-ray astronomy. The goal of this program has

been to develop space and bLlloon qualified HgI 2 detectors.

Previous to this development effort, the balloon group flew large

area balloon-borne arrays of scintillation detectors to observe hard

X-ray astronomical sources. On May 8, 1780, H. balloon gondola containing a 1500 cm` array of phoswich scintillation detectors was flown

from Palestine, Texas. This gondola, built at MIT at the balloon lab

In collaboration with the Cosmic Ray Working Group at Leiden

^r University, was the most sensitive astronomical hard X-ray (20-240 keV)
vations, tte gondola free-Yell from 3d kilometers and was destroyed

(the data were telemetered tj the ground and rat lost). The destruction of Lh:s instrument was the impetus `Dr accelerating the

development of mercuric iodide for astronomical applications to replace

the 1500 cm` array with an instrument of comparable sensitivity to

faint hard X-ray Sources.

The replacement gondola proposed was a collection of 4 Bragg

diffraction X-ray concentrators thc.t increase the effective area of 4

well-shielded low background mercuric iodide detectors. Calculations

show that this new gondola detector system can match the sensitivity )f

the old 1500 cm array if the assumed b 2kground count rate at balloon

float aititudes (crudely scaled from that of germar.i-um detectors in

orbit) is correct. In the work reported here a p r atotype detector

inside a small balloon-borne gondola has been constructed and flown at

balloon altitudes in order to test these calculations. The measured

background count rate results were then used tc confirm detector models

that can be used to predict and optimize the background of future

detector geomperies.

P R

I^

tII

L= I

1

l.
1.

d^

i

1

DEVELOPMENT OF N:.+v, TYPES OF DETF.CTURS FOR ASTRONOMICAL WORK

The ideal detector for hard X-ray astronomy would have excellent

resolution, high efficiency, large area and be lightweight (including

support equip • aent, i.e. cry)genics). This "dream" detector does riot

exist though most systems in use satisfy some of the requirements. The

st 9 ndard detector for astronomi:al hard X-rays has b!en large area Nal

scintillation crystals that operate at ambient temperatures but with

limited energy resolution (1'(o FWHM at 60 keV). The semiconductor

detectors such as germanium have recently achieved respectable areas

( - 60 cm2 ) yet their lower efficiencies require thicker detectors which

increase the background counting rate. Also the re q uirement for

cryogenic temperatures has limited their satellite lifetime to -1 year.

:-Mercuric iodide Detectors

Mercuric iodide (HgI,) solid state X-ray detectors have many

L

unique advantages in their application to X-ray astronomy. HgI,,

detectors operate at room temperature so that bulky and massive cryogenic coolers are not required. HgI 2 detectors also combine a high

quantum efficiency with good energy resolution. Because of the excellent quantum efficiency (due to the high effective atomic number of

IIgI

?

1 background contributions which scale as the thickness of the

detector (e.g. those dae to neutrons and garima rays) are lower than for

other hard X-ray detectors with equivalent stopping power. At this

time the largest area single detector that can be made is only 1 cm2,

so it is far from being ideal. Nonetheless, its properties are well

matched to a scheme that increaseG the effective area using a Bragg

diffraction concentrator.
Mercuric iodide tias been under deve^opment for ten years, and the

balloon lab at the Center for Space Research at VI': has had a working

experience with these detectors since 1979. The first use of HgI 2 in

an astronomical instrument was a small array of 12 detectors with a

combined area of 8 cm constructed here at the balloon lab at MIT and

flow► , in a small balloon-borne gondola in June, 1930. Though the

background counting rate at float was high because of a less than

optimum design, the instrument eas able to set an upper limit tc the

hard X-ray flux of Cygnus X-1 in the "high state" (Ogawara et al 1982)

HgI 2 detectors can still be considered "experimental" in that they

cannot be acquired commercially and each Bete-:tor is in some way

unique. In work carried out b;r the MIT balloon group, the HgI, crystals have been sub,;ected to extensive field testing. They have been

exposed to extreme variations in temperature, pressure, humidity,

voitage, light level and handling. As might be expected with an

experimental material, HgI. did not always pass with flying colors.

However, from this experience, general procedures for handling and

environmental limitations have been developed to extend the operating

life and reproducibility of results necessary for a field detector.

Because of its relative softness, the direct handling of HgI 2

crystals that have been cleaved from the large 20 cm parent crystal

must be minimized. The "thick" crystals (--300-1000 microns) usea at

I the balloon lab for hard X-rays have all been acquired from EG&G, Santa

Barbara. The crystals are fabricated into detectors with evaporated

palladium electrodes, mounted on ceramic or plastic substrates and

/ 
detectors is 1 cm` with energy resolution ranging from 2.8 to 15 keV

"+OHM at 60 keV. A typical energy resolution would be 4 keV at 60 kc,V.

The area limitation: requires many detectors in an array to achieve an

astronomically useful collecting area (..100 -.m 2 ). Such an array would

also require 100 measurement chains (100 preamps + voltage amplifiers)

as well as e x isting crystal growers to produce high resolution uniform

detectors in iraantity. Until the detector area can be made larger,

some method of concentrating X-rays :oust be used to increase the

effective area.

In the laboratory the detectors :oust be kept clean and dry. Small

amounts of water and contaminants (dirt, solder flux, etc.) create

current leakage paths that i n crease the overall RIMIS noise and small

_ntermi.ttent current pulses similar to a relaxation oscillator.

Encapsulated detectorf, should not be exposed to vacuum because trapped

air pockets will expand and strain the crystalline lattice, thus

ruining the charge transport properties.

The temperature effects are striking for some crystals and un1.

noticeable with others. T;Ze mobility increases with lower temperatures

for both electron and holes, but the mean lifetime of an ele^iron or

hole before it is trapped by the lattice decreases with decreasing

temperature. These competing processes ietermine a temperature where

the resolution is best (i.e. the mean free path, a = VTR, is maximized)

and is a function of the quality and purity of the crystal.

nipirically, this temperature is between -20°C and 20°C, varying with

detector (Ricker 1982). These temperatures are easily achieved at
(thermoelectric) techni,ues. Higher temperatures decrease the mobility

snd also produce a prohibitive Leakage current which can saturate DC

coupled preampl' Piers .

The bias voltage across the HgI,, must be l ipplied slowly to avoid

large induced currents. After the bias voltage is applied, trs detector "warms up" over a time period of -1 minute .o 1 hour depending on

the crystal. The pulse height for a given incident energy increases

and the resolution improves. This can be interpreted as an increase n

the depletion region ss the deep traps (energy levels in the forbidden

region caused by impurities) are emptied by the large electric field.

Resolution usually improves with increasing electric field until the

leakage current becomes too large. tU so, hi g h electric field strengths

increase the chance foe surface breakdown which intermittently injects

current pulses into the preamp and shows up as low energy events. The

frequency of occurreice of this "current infection noise" increases

with applied voltage.

'zany HgI 2 detectors perished in the balloon lab due *,o different

types of abuse.. However, by minimizing any ;candling Uhat would stress

the crystal and by keeping the detector in a clean, cool and dry

assembly, it was possible to produce a stable and dependable operating

11gi 2 X-ray detector.
f

1

Previous to the research work described here, HgI 2 showed great

theoretical potential for use in hard X-ray astronomy. The combination

of high energy resolution at room t.!mperature plus a high photon

quantum efficiency seemed id--al for many balloon and Latellite detector

appl:.cations. Yet, experien e in the field operation of Hg T was

minimal and background count rate information at high altitudes did not

exist. Development of iigi 2 specifically for use as -xn astronomical

X-ray aetector here at MIT over the last 4 yee.rs :ias changed this

situation. Procedures for handling and pac&aging detectors :lave been

developed that produce stable and dependable figl 2 detectors for use :6n

the field. HgI 2 detector assemblies have been constructed and flown at

balloon altitudes and the detectors operated flawlessly on all

occasions.

A ma,;or limitation to the application of HgI 2 in X-ray astronomy

is the inability to fabricate individual detectors larger than -1 cm2.

A scheme to overcome this limitation involves the use of X-ray concentrators to increase the effective area of a shielded, low-background

HgI 2 detector. The concentrator consists of a paraboloid shell lined

with liti.;= fluoride crystals that Bragg diffract on-axis incoming

X-rays to the focus of the paraboloid, where a small area detector is

located. To establish the expected sensitivity of this concentrator

/detector design, the background counting rate of the shielded HgI2

detector at balloon altitudes had to be determined. A balloon-borne

gondola containing a Prototype detector assembly was designed and

constructed to measure this background rate. The prototype detector

i

E 
ate acintil].ator operated in anticoincidence. :his gondola was flown

twice in the spring of 1962 from Pa'estine, Texas.

The aecond flight of this prototype instrument established a

background counting rate for a bismuth germanate-snielded mercuric

Iodide detector of 4.2!0.7 x 10 -5 counts/sec cm xeV over the energy

range o° 40-60 keV. This measurement was within 50" of the predicted

value. 11 prediction was based on a Monte ;,arlo simulation of one

detector assembly in the radiation environment at float altitudes

(-40k.m). The Monte Carlo simulation can now be used to investigate the

detector design parameters (geometry, discriminator lev-1s, etc.) to

attempt to lower the background rate even further.

Based on a background rate of 4.2 x 10 -5 cts/sec cm` keV, calcuiations indicate that with 4 concentrators and 4 detector., systems a

celestial X-ray source with 1/100 the intensity of the Crab Nebula

could be detected in 1.5 hours. This excellent sensitivity is achieved

by simultaneously reducing the background by 1"30 shielding a HgI2

detector, and increasing the effective area using hard X-ray concentraters.

The prescription for reducing the background rate by a factor of

-10 over previous balloon-borne hard X-ray detectors can be summarized

very simply: use very high quantum efficiency detectors shielded by

dc,ise, high Z active scintillators. Because they are well suited to

this experimental approach, HgI 2 and B(,O are both likely to `jecum e

ma,'or detector materials in X-ray astronomy. Furthermore, their ease

of handling and improved performance at moderately reduced temperatures

enhance their desirability for balloon and satellite use.
PUBLICATIONS

1. Ogawara, Y. et al, X-Ray Observations of the 1980 Cygnus X-1 "High State",

1982, Nature, K5, 675 A - A 50

2. Ricker, G. R. et al, New Measurement of the Fano F--tor of Mercuric Iodide,

1982, Rev. Of Scientific Inst., 53, 5. p.700 441-3272/

3. Ricker, G. R., Vallerga, J. V., Wood, D. R., A Mercuric Iodide Detector

System for X-ray Astronomy, 1983, Nucl. Inst. and Methods, Vol. 213, #1

4. Vallerga, J. V. et a_l, A Bismuth Germanate-Shielaed Mercuric Iodide X-ray

Detector for Space Applications, 1982, IEEE Trans. On Nucl. Sci., Vol. NS-29,

A l, p. 151 fie)_ 6 ///

5. Vallerga, J. V., 1982, "Studies of New Hard X - g ay Detection Metnods and

Observations of Cosmic X-Pay Sources," Ph.D. Thesis, M.I.T.

6. Vallerga, J. V., Vanderspek 't. K., Ricker, G. R., A Mercuric Detector

System for X-Ray Astronomy, 1983, Nucl. Inst. and Methods, Vol. 213, al

r
Submitted to Nature on 15 October 1981

X-RAY OBSERVATIONS OF THE 1980 CYGNUS X-1 "HIGH STATE"

Y. Ogawara and K. Mitsuda

Institute for Space and Aeronautical Science

4-6-1 Komaba, Meguro-ku

Tokyo, Japan

K. Masai

Department of Astrophysics, Nagoya University

Chikusa-ku, Nagoya, Japan

J.V. Vallerga, L.R. Cominsky, J.M. Grunsfeld, J.S. Kruper, and G.R. Ricker

Department of Physics and Center for Space Research

Massachusetts Institute of Technology

Cambridge, Massachusetts 02139

CSR-HEA-81-25

This work supported in part by NASA grants NGL-22-009-015 and NSG 7339.
 q r 

4.

Observ,ition,. of Inten v,ity transitions in tti p X-ray errission from th

bIact- hnle r8n f ildato Cy Ojnu-. X- 1 (see refers-nce 1 and refrr^nc( —, tIwrrtn

marr c- by the HAuchn s,,tell it o

(ILIrin(I the surimer o )(I f', U and are reporte

ho re. In addition, we report t ► ,e results of hard (20 -1[U k(, V) X - riv

ohservations of (:Vq X-1. These data were ohtained (III ► • inn the "high” int 

state b y hailoon-home rtorcuric iodise detectors, and ere connorn-i with 

ohtainrd at similar enemies approximately one month earl ier 3 Oi,rin(l

i the 

intensity state.

Cyq X-1 was observed to be in a "high state" by *.he Uhuru smitcllit

December 1970 (its first ohservation of the source) to April 197 1 dnd a

during 197f, by the Ariel V 1-7 , and OSO-f; instritment.s a ,, well as the

Astronomical iotherlands SatelIite Q . Suhsenuont ohservations at many

di f ferr v -ra y enoraic-s h:.ve inricated that rye) X-1 has been i n an exte

Ic,4' Mate since early 1976. In his extensive review of Cy r] X-1 (,-ef

Oda h as summarized th n light curves and spectral information which were

nuhl ished as of 1977. During tho transition to a "high" state, the sof 

kPV) and hare{ (>10 keV) X-ray intensities appear to he anti cor-relat.ed, 

"high" state below 10 keV is associated with a substantial decrease in t

X-ray intensity above 10 keV.

For the work reported here, soft X-ray ohservations (1-30 rc :V) we

with the scanning Y-axis detectors on board the Hakucho satellite. The 

detector system consist ,, of a pair of XPnnn- filled proportional count.er

(SFX-V1 and SFX-V2) with a 1`.7x32° FWHM field of vier. Each of the two

counters has an effective are,, of 32 cri2 . nnl.Y the SFX-V1 data was use

this analysis since the SFX-V2 field of -.iev. contained Cygnus X-3 during 

of the ohservations, the satellite spins at 6 rpm, — that :3000 transi
, I i vem sourer are ava i 1.:hle duri rl'I the 5 da i 1 y 1 0 mi nut r lt,OV(- the

tra r kina stat iorr at Ucir inoura, Jal,trn. The source 1.1 • an!.its are superpmed and

corri-cted for aspoct tronsr ► is,.ion and Garth occultation. [lacUiromid is also

sub Iractf, (^. (For furthrr details on tlic 11al.ucho satellite see reference 10.)

Figure 1 .hoot, a portion of the X - ray light carve obtained wi 11 the

llakucho Y-axis detectors dur-rnn 1980. Each point repr., scots the d, , ily avc-racie

X- ray intensity fror-r Cyn X-1. lr-ro "high" statr:s are cloar • ly visih1c. The

first lasted at least fifteen Oys, and wa,; fc,llo-::ed by a decreatir, over

appror.irrately 25 drys to the pre-transi t.ion level. A second increase was

observed approximately four months later. llris increase was quite brief, with

the "high" stat^ persisting for only two days. The erratic variahility seen

ri ng tho 1980 obs(, rvat i ons is cirri 1 ar to that ohs^rved in 1975 and rc.portc,d

by Holt et ar . using the Ari el V All Stay lloni t.or 7 .

Also indicated in Figure 1 arc the dates when hard Y.-ray (20-200 keV)

g alloon observ, ons of Cya X-1 took place. The "low" state ohservations (8

May 1980) wcre nerferr)od by t.hc, t'll /Lrei den Ral 1 oon Group and will he reported

separately 1 , the "high" state oh;rrvations (28 Jone 1980) constituted the

first ase of a mercuric iodide X-ray detector- in an astrono^jical expr-r'inlnnt.

The detector consisted of 11 H(ri 2 crystals with a total effective area of 7.G

crrr . The detector asscr,t, l y (crystals and prearrps) was enclosed within a

cylindrical Nal scintillator wall (active anti coincidence shield) which

rojocted charged Particles and n,ira.;a rays. The u ►°efl iriht enerc!y resolution of

tho detectnrs ramrd from 2-6 kcV FbHII at 00 kcV and remained stable during

the 13-horn • fl i rlht from ('cal esti n Texas. The dondol a contai ni nn tho detector

was larrnehod on an 11 x 106 ft.3 halloon at 001.5 U1' on 29 June 191;0 and floated
3

• .t for 13 hours at ;:n average al ti tudu of 40 Im P 20 9M CS),, F c y irruI'r

exposure to Cyq X" 1 oc" ccn re p ! at 043N Of 24 ,iohc MO. lhn dolvrto • s wore

calihral d ";inn an on-buird Am 241 source for tern winw % nut of evor y hour.

No cnunt.s W each Aop, rote Hq1 2 dntuctor vere binned in 64-Mond int.orvWs

in coarso e fern channels, 10 keV wide. The dnteetor count; were than summed

to give too total cnun;s per enerey (hannol ncr 64 seconds.

Total expnsure to fhr X-ray source was calc"lat".ed as a function of thee.

the calculat i ons Wok into nrco"nt aspect NOW, W, W i l"AP, sidvrcal t'i'le,

and field of view of individual H91 2 detectors), atmonpWre and doii ctor

window transmission, detcctnr efficiency, and d^ad tine. The maxima" expo;ur

to Cygnus X - 1 at transit was 37 percent over the 50-7n heV energy range. For

each vnorgy channel, an effective time on source, Left =

0

. t i r i , was

i

calculated and used to convert observed counts to a photon flux at the top of

the atmosp b ore. Here, t i refers to an individual t.imr bin and c is the tota

expos"re including it o nsrrheric and aspect corrections as discussed 6hovo.

There was less than a 2 sipa detection in all energy hands.

Cygnus X-1 parsed through the field of view in 95 minutes result inn ii 8

minutes of "on source" data (i.e., 5 minutes were nccupied by cal iterations an

7 minutes were lost due to telemetry dropouts). Forty minutes before• and

after this period were averaged and ",rd as background.

Finure 2 shows the 2-sigma upper limits we have dorived for the hord

X-ray spect rum during the 28 June 1980 Cyg X-1 "high" state observation. Als

indicateo in figure 2 are the results from the £t May 19HO hard X-ray

ohservations by Scheepmaker et a1. 3 which wore obtained vhen Cvq X.-1 was in

the "low" state, as well as representative herd X-ray measurements during the

1 71 and 1975 "high" states. The 1-12 keV and 12-10 QV X-ray intensities

) IU'11^ 
4

fr(n the Ii kU r In nboorval inns have henr y avor aged nvnr the hi ► Ih state and 

state. As can he seen in l ipire 2 tho hord X-ray inten%ity is o"ticorr

with the :oft X-ray intensit y ; the pivnt point occurs for an energy bel

ktN. O q i "q dKe emr, on U0 ..3, Li, ng et al .1 1 , l? also recorded a cant i 

decrease in the hard X-ray intrn4ity (Q-43 00 from 19 May 19RO un

least. 2 June, lh , npeetral anticorr p 1 atio" repnrt( d here was also obse

previnun low try hi g h state transiLionn 8,13 &nd seers to he a conFist.ent

feature of the hirrOO hohavior of Cygnus X-1.

I rWous observations have established that the spectrum varies wi

during the di ffnrent intensity suites and is not well - represented by a 

relation (c^.(t ., nower law or exponential). This has led to the idea th

soft component of the spectrum ray oriCinate in the outer optically thic

region of an accretion disk surrounding the central compact object, whi

hard component arises in an optically thin, 'rot inner region" The ob

time scales for the transition to and decay from a "high" state are cons

with theoretical calculations for the evolution of a standard accretion

disk , and sugg p st fluctuaticns in the accreted stellar wind cr change

the optical depth of the g as at tho outer edge of the disk (see referenc

and 16 and references therein). Spectral anticnrralation has also been 

to explain the energy dapordence of the short term variability (mvec to

sec) 11 . A nu"Wcal model proposer' by Ichicraru 18 shc. , s that two physic

distinct states cen exist in the accretion disk depending upon whether 

radiative 1o ;s is smal ler thin or larger than the vi scous heatinn near 

outer hnund;,ry (i.e. a transition i s expacted to occur for A /T ?

 _010 3 

K-2 where A is the accretion rate and T is the gas tv"veraturo at the ou

edge of t're disk) Thus, th^ middle part of the disk may hecone optica

thick (high state) or optically thin (low state).
I ri 

3

OWGIP!AL FAGE i:i

`;F POOR QUALITY

in conclusion, 1W havv es t oblinhed th..t Cy4nrrs X-I r(', d to its hQ

1 v in 1000 al ter an opparono 4 .5 year Mtvi vdl in thv 10','i `tirt e. It 0

intcrestinn to note th rt this tire perind in similar to 1he n year inturv

het; r" the f i rn t p and second nh , -rvod t rans i t i on;. n ' f . We have est„51 0 I rid

that the soft and hard X-ray inUnsitivs are Micorr p latvd in the transit

hctwcc n rho tun Mist inct states and wp hol ivve that any theoret.icol r a 41 

the bimodal behavior of Cygnus X -1 must account for this spectral antieor

lation.

AC ► HOULEDGEMEMIS

The authors thank of l the m mbers of the fl, I;ucho tear, for th p i r

contributions to the observations. the assistance of Drs. H. !none, K.

Makishima and 1. Murakani in the data analysis and preparation of the

manuscri pt is aloo acknowledged. The authors are also grateful for the

assistance of: Hobert Kuhara and the launch crew at the National Center f

Atrmsnh p ric Research; Dick Lynn, Carol Ortalo, Wayne Schnepple, and Jahn

Warren at the Advanced Masur p ments Laboratory pit EG&G (Santa Parbira). A

Scheepmaker kindly provided data from the MIT/Leiden 1980 t,al l oon flight 

advance of publication. L.C. acFnowledg ^s support from 7onta Internationa

during p art of this wore this research was sponsored in part by

the Rational Aeronautics and Space Administration under grants

r N GL-2.2-009-115 and NNG7339.

1
6

ORIGINAL PACE 19

OF POOR QUALITY

R, fL1!Lt10E S

1. 0(1;(, P Space Science Pvvi-vvv;, 20, 75/- . 813 (1977).

2. Oda, 14. and the Hakucho Tears, IAU Cirrulor 3401 (1081(1).

3. Schei•vmator, A. (A al. nanuscript in pr(.pdration.

4. Tanonha ► c : i, II. vt A. Ast.ropliy-,. J. 177, 1-5-1 W (1071).

5. Hult, S. S. et al. t!ature 1 1 )(., 1083-10 (1975).

6. Sanford, P. et X71. Ha turp ?56, 109-111 (1975).

/. 11ol I ., S.S. et al. Ho turc 2.(11, 213-115 (1916).

S. Dolan, J.F. et al. A.stronhys. J. 236, 551 ` r 7 (1')7").

9. Heise, J. et al. tlature 250, 107--1083 (]h75).

10. Yondo, I. et al. Space Science In strur(entation, 5, 211 (19£?1).

11. Ling, J.C. et al. S .A.A. S. 13, t1n. ?, 1). 549.

12. Ling, J.C. (priv,te communication).

13. tlatteson, J.L. et al. in X -Ray Bi naries, eras, E. Boldt. and Y. Koi " dn (NASA

SP-389) , 407-424 (197S).

14. lhornr, Y.S. and [ , rice, R.H., Astrophvs. J. 195, 1.101-LI05 (1975).

15. Lightman, ,%.P. A , . t-runhys.J. 194, 419-421 W)74).

16. Eardley, D.M. et al. Comrients on As tro(► hysics (1978).

17. Doi , K. , 1Jature, 2-87, 210-212 (1980).

18. Ichi rn:aru, S., Astrophys.J . , 214, 840-855 (1977) .

i D
I^

7

F1 1L11 , CAP IIWIS

Fi(pirr: 1: So ft X - ray 0-1? VrV) 1 i0t curve of Cya X-1 obtairwd with the

llakucho catOlite. Two "high" states are evident:. Tho first bogan

before 10 June 19£30 nnct thr, second transient event occurred in

hovcwber 1980. Also shown sire the djteS of siruil ianeous

ohs,vrvatiow, in hare, X - rays by injruric'nts.

Figure ?: X-ray spectra taken hc, fore and after the "hi g h" state transition of

Cyq X-1 in late play/early June 19M. Data points v,m.'re ohtr^inrd

with the 11M,.ucho sated 1 i to and v. i th balloon borne• hard X--)-Oy

detectors. The Hakucho data r,oin +.s for high and low states are

averages over 1? June - 17 July, 19£10 and 1 - 19 Pay 1980,

respectively. Also inclOed are representative hard X-ray spectra

taken during previous high states.

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2

ST'JDIES OF NEW HARD X-RAY DETECTION nTHODS

AND

OBSERVATIONS OF COSMIC X-RAY SCURCES

by

JOHN VINCENT VALLERGA

Submitted to the Department of Physics on august 19, 1962

in partial fulfillment of the requirements for they

Degree of Doctor of Philosophy in Physics

ABSTRACT

To establish the expected sensitivity of a new hard X-ray te_escope

design, an experiment was conducted to measure the background counting

rate at balloon altitudes (40 km) of mercuric iodide, a room temperature

solid state X-ray detector. A balloon-borne gondola containing a

prototype detector assembly wac designed, constructed, and flown twice in

the spring of 1982- from Palestine, Texas. The prototype detector

consisted of thin mercuric iodide (HgI,) detectors surrounded by s large

bismuth gp rmanate (Bi.,:,e,012) scintillator operated in anticoincidence.

The bismuth germanate shield vetoed most of the background counting rate

induced by atmospheric gamma-rays, neutrons and cosmic rays.

The second flight of this instrument established a differential

background counting rate of 4.2!0. 7 x 10 " ' counts/ sec cm' keV over the

energy range of 40 to 80 keV. This measurement was within 50% of the

predicted value. The prediction was based on n Monte Carlo simulation of

the detector assembly in the radiation environment at float altitude.

Two results are also presented on the hard X-ray (20-200keV)

observations of the COS-3 gamma-ray source, CG 135+01, and the Galactic

Center using a balloon-borne array of phoswich detectors with a total

active area of 1200 cm'. The observations took place during a high

altitude balloon flight on May 8, 1980 from Palestine, Tx. CG 135+01 was

detected with a 4 sigma level of significance while upper limits vere

placed on the flux from the Galactic Center, confirming the reported

variability.

Thesis Supervisor: Dr. George R. Ricker

Title: Principal Research Scientl9t

'RE?CE;I)INC, PA(,E BLANK NOT HUTM

[I 

ft^

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4

Dedicated to Helene

i

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s ` ID

4

ACKNOWLEDGMENTS

I first wish to thank my advisor, George Ricker, who taught me the art

and science of high altitude ballooning. I also want to thank those people

who helped in the construction and preparation of the many 'jalloon gondolas:

Jim Ballintine, Darren DePo,y, John Doty, Doug Franklin, Dan Griscom, John

Grunsfeld, John Kruper, Pete Tappan, Roland Vanderspek, and Darien Wood. I

also would like to thank Wayne Schnepple and Carol Ortale for all their help

with mercuric iodide. Robert Kubara and the crew at the National Scientific

Balloon Facility did an outstanding job launching our payloads. I am also

grateful for Brenda Parsons' help in producing this and many other documents.

s^

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641

5 qj

TABLE OF CONTENTS

Page

LIST OF FIGURES ...................................................... 7

LISTOF TABLES ....................................................... 8

1.0 INTRODUCTION .................................................... 9

2.0 ASTRONOMICAL HARD X-RAY DETECTORS IN CURRENT USE ................ 12

2.1 Proportional Counters .................... ................. 13

2.2 Scintillators .............................................. 14

2.2. 1

Phoswich Detectors .................................. 17

2.3 Solid State Detectors ...................................... 19

3.0 DEVELOPMENT OF NEW TYPES OF DETECTORS FOR ASTRONOMICAL WORK..... 23

3.1 Mercuric Iodide Detectors ........... I ..... I ................ 23

3.1.1 Detector Production and Characteristics ............. 25

3.1.2 Physics of Detectors and Energy Resolution, ....... 26

3.1.3 Experimental Development for Astronomical Use....... 32

3.2 Bismuth Grrmanate as Active Shielding ...................... 35

4.0 RESULTS FROM MAY 1980 BALLOON FLIGHT OF LARGE AREA (1500 cm 2)

PHOSWICH DETECTOR ............................................... 38

4.1 Instrumer.c ................................................. 38

4.2 The 1980 Flight ............................................ 41

4.3 Detection in Hard X-rays of the COS-B

Gamma Ray Source CG 135 + 1 .................................. 42

4.4 Upper Limits on the Hard X-ray Flux of

the Galactic Center ........................................ 49

4.4.1 May 1980 Drift Scan ................................. 53

5.0 DESIGN CONSIDERATIONS FOR A NEW HARD X-RAY TELESCOPE............ 57

5.1 Sensitivity ................................................ 57

5.2 Bragg Diffraction Concentrator ............................. 60

5.3 Low Background Balloon-Borne Telescope Design .............. 60

6.0 SHIELDED MERCURIC IODIDE BACKGROUND MEASUREMENT EXPERIMENT...... 65

6.,1 Prototype Detector System .................................. 65

6.1.1 Primary Mercuric Iodide Detector .................... 66

6.1.2 Bismuth Germanate Shield ............................ 68
f 6

Page

6.2 Balloon Gondola and Associated Electronics ................. 70

6.2.1 Event Logic Electronics ............................. 73

6.2.2 Data Recording ...................................... 75

6.2.3 Ground Support and Data Analysis .................... 77

7.0 BACKGROUND SOURCES AND MONTE CARLO ANALYSIS ..................... 79

7.1 Gamma Ray Induced Background ............................... 81

7.1.1 Monte Carlo Simulation .............................. 81

7.1.2 Results of Simulation ................ 6.............. 89

7 .2 Neutron Induced Background ................................. 89

7.3 Charged Particle Induced Background ........................ 96

7.3.1 Spallation and Bremsetrahlung ....................... 98

8.0 RESULTS OF BACKGROUND MEASUREMENTS WITH A SHIELDED MERCURIC

IODIDE DETECTOR AT SEA LEVEL AND AT BALLOON ALTITUDES........... 101

8.1 Sea Level Background ....................................... 101

8.2 Energy Calibration ......................................... 104

8.3 Balloon Flights ............................................ 107

8.4 Flight Results and Comparison to Predictions ............... 108

9.0 CONCLUSIONS ..................................................... 117

APPENDICIES.......................................................... 120

REFERENCES ........................................................... 129

r
S l} ^^ . Y fi s'H f^:Y

7

LIST OF FIGURES

Page

2.1 Phoswich Detector Asssembly ........ ...... ................... 18

3.1 Linear Attenuation Coefficient of Mercuric Iodide

and Germanium .................................................. 24

3.2 Mercuric Iodide Detector Charge Collection Efficiency

Simulation........ ......... ................................ 31

3.3 Linear Attenuation Coefficient of Bismuth Germanate

and Cesium Iodide .............................................. 36

4.1 Soft X-ray Sources and Field of View of Collimators

near the Gamma-ray Source CG135 + 0 1 ............................. 43

4.2 Hard X-ray Spectrum of CG 135+01 Assuming Two Possible

Locations .............. 45

4.3 Hard X-ray Spectrum of the Galactic Center ....... ............. 52

4.4 Collimator Response vs. Time to the Galactic Center............ 54

5.1 Sectional View of Proposed Hard X-ray Diffraction

Telescope ......................... I............................ 61

5.2 Bismuth Germanate-Shielded Mercuric Iodide Prototype

Detector Assembly .............................................. 62

5.3 Photograph of Gondola Containing Prototype Detector............ 64

6.1 Detailed View of Mercuric iodide Detector Assembly ............. 67

6.2 Schematic of Diffration Telescope Aperture Shield .............. 71

6.3 Transmission of Graded Shield (Pb, W, Sn) ...................... 72

6.4 Event Logic Diagram ............................................ 74

6.5 Cassette Tape Data Format ...................................... 76

7.1 Ambient Spectrum of Gamma-rays at Balloon Altitudes

and Transmission of 2.5 cm. of Bismuth Germanate ............... 82

7.2 Two Dimensional Monte Carlo Ray Trace Plot ..................... 86

7.3 Predicted Valid Background Sp--tram due to Gamma-rays.......... 87

7.4 Predicted Veto and Bottom Background Spectrum due to

Gamma-rays ..................................................... 88

7.5 Predicted Valid Background Spectrum due to Neutrons............ 93

8.1 Sea Level Valid Background Spectra at Various Locations........ 102

8.2 Valid Background Spectra with Anticoincidence Shield

Detectors On and Off .. ..... ...... ........ ..I............ 105

8.3 Mercuric Iodide Flight Detector Calibration Spectrum

of '" Am Source .................. ....... .......I............. 106

8.4 Veto and Bottom Spectra from Flight 1282P ...................... 109

8.5 Veto and Bottom SpA ,^tra from Flight 1292P ...................... 110

8.6 Valid Spectrum from Flight 1292P .................. ..... ...... 111

8.7 Count Rate vs. Time of Bottom Detector during Flight 1292P..... 113

ri

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1QW, J
8

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LIST OF TABLES

Page

2.1 Properties of Commonly Used Scintillators ...................... 16

2.2 Properties of Commonly Used Solid State Detectors .............. 21

3.1 Charge Transport Parameters of HgI 2

and Ge ..................... 30

4.1 Flux and Derived Luminosities for Two Possible

CC135 +01 Candidates........ .. .

2

..................... 47

7.1 Int p grated Radiation Flux at 3 gm/cm and

Magnetic Latitude 420 N (Palestine, Texas) ..................... 80

7.2 Predicted Background Count Rate Due to Gamma Rays .............. 90

7.3 Predicted Background Count Rate Due to Neutrons ................ 94

8.1 Measured Surface Background Count Rate at

Various Locations ................ ...... ...................... 103

8.2 Comparison of Measured and Predicted Background

Count Rate at Balloon Altitudes ................................ 116

Oerl' i.^
9

1.0 INTRODUCTION

I

Since its beginning in 1962 with the discovery of Sco X-1, the

science of X-ray astronomy has progressed by placing X-ray photon

u

detectors above the atmosphere in rockets, balloons and satellites.

This usually involves modifying the detectors to withstand extreme

environments and payload constraints (size and weight). The goal has

always been to increase the sensitivity of the detectors (signal to

background ratio) either by increasing the effective collecting area

and/or observing times while decreasing the background counting rate.

1 Each new detector technology that is applied to X-ray astronomy must be

reoptimized to achieve these goals using widely different strategies.

In the last few years, the balloon lab at the Center fir Space

Research at MIT has been studying mercuric iodide (HgI 2 ), a solid state

room temperature X-ray detector, for use in hard X-ray astronomy. Its

primary advantages are its good energy resolution at room temperatures

j and a large photon cross section. Much work has been done to ci.aracterize the pro perties of HgI 2 that affect the operation of these

detectors for use in X-ray astronomy. The goal of this program has

been to develop space and ^alloon qualified HgI 2 detectors.

Previous to this development effort, the balloon group flew large

iarea balloon-borne arrays of scintillation detectors to observe hard

X-ray astronomical sources. On May 8, 1980, a balloon gondola containing a 1500 cm: 2 array of phoswich scintillation detectors was flown

from Palestine, Texas. This gondola, built at MIT at the balloon lab

g:

in collaboration with the Cosmic Ray Working Group at Leiden

University, was the most sensitive astronomical hard X-ray (20-240 keV)

I. b
ie

detector in existence. After 16 hours of very high quality observations, the gondola free-fell from 38 kilometers and was destroyed

(the data were telemetered to the ground and not lost). The des-

'ruction of this instrument was the impetus for accelerating the

development of mercuric iodide for astronomical applications to replace

the 1500 cm array with an instrument of comparable sensitivity to

faint hard X-ray sources.

The replacement gondola proposed was a collection of 4 Bragg

diffraction X-ray concentrators that increase the effective area of 4

well-shielded, low background mercuric iodide detectors. Calculations

show that this new gondola detector system can match the sensitivity of

the old 1500 cm array if the assumed background count rate at balloon

float altitudes (crudely scaled from that of germanium detectors in

orbit) is correct. In the work reported in this thesis a prototype

detector inside a small balloon-borne gondola has been constructed and

flown at balloon altitudes in order to test these calculations. The

measured background count rate results will be used to confirm detector

models that then can be used to predict and optimize the background of

future detector geometries.

This thesis begins (Section 2) by describing the currently popular

detectors used in hard X-ray astronomy. Section 3 goes into detail

concerning the two new materials used in the prototype detector:

mercuric iodide, the primary solid state X-ray detector; and bismuth

germanate, a scintillator with a very large photon cross section, which

is used as an anticoincidence shield. In section 4, two of the most

iC significant results from the 1980 flight of the 1500 cm gondola are

QJl
O

11

discussed. These results show the p otential of a new gondola design

that can achieve this excellent sensitivity. The design considerations

for the new gondola detector are explained in Section 5 while Section o

is a detailed description of the detector prototype gondola and associated electronics. The prototype detector model used in predicting the

background count rate due to gamma rays, neutrons and charged particles

is described in Section 7. The results of two balloon flights of this

prototype detector are compared with predictions in Section 8. A summary and conclusion is presented in Section 9.

i^

1.
c

'J'

12

2.0 ASTRONOMICAL HARD X-RAY DETECTORS IN CURRENT USE

Detectors used to count photons from celestial X-ray sources

differ greatly from their laboratory cuunterparts. In general, high

sensitivity to faint sources is the goal in astronomical detector

design. Because the atmosphere is opaque to hard X-rays (20-200

keV) for column densities greater than 20gm/cm 2 , studies of cosmic

sources must be carried out from satellites or high altitude balloons.

Thus, weight and volume of the instrument are at a premium. The

detectors should have large areas and a high quantum efficiency to

collect as many source photons as possible and also be well shielded to

reject background events induced by primary cosmic rays and their

secondaries. Good energy resolution increases the sensitivity to

features in a source spectrum and can "resolve out" line features in

the background. Some of these goals are incompatible and cannot all be

found in one type of detector. The final choice is usually determined

by the scientific objective of the observation. For example, timing

observations of X-ray pulsars require large areas to acquire many

photons from the source and can tolerate a somewhat higher background

rate because of the time signature of the source. On the other hand,

energy resolution might have a higher priority in observations of

possible cyclotron line features in magnetized neutron stars.

It takes an appreciable detector mass to stop hard X-rays. The

most sensitive hard X-ray telescopes have been flown from high altitude

balloons where payload weights can reach 2000 kg with sensitive areas

of 3000 cm 2 . A typical long balloon flight lasts only -30 hours, so

that satellite-borne hard X-ray experiments can make up for reduced

t

t^
13

areas with longer observing times. The most common astronomical X-ray

detectors used on satellites and balloons have been proportional

counters, scintillators and cooled solid state detectors. A detailed

review of these detectors can be found in Peterson (1975) and references therein. What follows is a summary of the properties of various

detectors, emphasizing the advantages and disadvantages of each type

for use in X-ray astronomy.

2.1 Proportional Counters

A proportional counter consists of a gas filled chamber traversed

by thin anode wires (_.01 mm diameter) carrying a potential of several

thousand volts. A photon entering the chamber may interact with the

gas liberating a photoelectron that ionizes the gas, producing additional electrons and ions. As the electrons move toward the anode

they can further ionize the gas resulting in a "gas multiplication" of

-10 3 to 105 . This electron charge pulse is collected at the anode and

is proportional to the energy of the X-ray. The most quantum efficient

counters contain xenon (Z-54) with mixtures of a "quench" gas such as

methane or carbon dioxide to absorb UV radiation and de-excite metastable states of the positive ions.

Xenon proportional counters require a large thickness to achieve

reasonable efficiencies and have mostly been used for energies below 60

keV (Clark at al. 1973)• The photon mean free path in xenon at STP at

50 keV is -15 cm. The small efficiencies of proportional counters at

high energies can be somewhat compensated by the very large areas that

can be constructed and deployed such as the High Energy Detectors (HED)

of the A2 experiment on the first High Energy Astronomical Observatory

4
a4

14

I

,

(HEAD 1) (2500 cm geometric area, Rothschild et al. 1979)•

Background reduction is accomplished by using ground plane wire

i grids that section a common gas volume into many elements, each with

a

its own anode. This makes the "wall-less" counters with the outer

i

elements operated in anticoincidence (Serlemitsos 1971). Non-local

interactions, such as high energy charged particles, can be rejected

because of simultaneous signals in many anodes.

The method of charge collection in proportional counters can be

used to advantage in locating the position of the photon interaction.

A charge pulse produced locally floors in two directions along the anode

wire and the ratio of the charge collected at both ends can give the

one dimensional position of the interaction along the wire. Mal tiwire

counters using this position sensitive technique can be used to determine the two dimensional position of interaction. These position

sensitive detectors are used at the focus of grazing incidence X-ray

focusing telescopes such as the Imaging Proportional Counter (IPC) on

the HEAO 2 (Einstein) satellite (Ciacconi et al. 1979).

Though large areas can be constructed in a straightforward manner,

proportional counters have not been very effective in hard X-ray

astronomy above 60 keV because of the very low quantum efficiencies at

these energies for reasonable volume detectors.

2.2 Scintillators

A photon interacting in an inorganic scintillator will release an

c^

These states de-excite to the ground stat•o by emitting optical or

energetic electron that excites elevated lattice states by collision.

ultraviolet photons. The number of photons released is proportional to
15 T

the energy of the incident X-ray (deviations from linearity in light

yield vs. energy do exist; a complete description of the theory and

experimental techniques of scintillators can be found in Birks, 1964)•

The optical photons that exit the crystal impinge on the photocathode

of a photomoltiplier tube. The photocathode releases photoelectrons

which are multiplied by the dynodes producing a charge pulse proportional to the X-ray energy deposited in the crystal.

Inorganic scintillators such as sodium iodide (NaI) and neeium

iodide (CsI) have been the "work horses" of hard X-ray astronomy

because of their high stopping power. (The mean free path for a 100

keV photon in NaI is 1.8 mm.) Organic scintillators, made up mostly of

hydrogen and carbon, are almost transparent to hard X-rays yet can

detect charged particles very efficiently. Thus, they are frequently

used as anticoincidence shields to reject cosmic rays because of their

relatively low cost and availability in all sizes and shapes:

The properties of some scintillators are shown in Table 2.1.

Thallium d.oped NaI is popular because it has the highest light yield

of all scintillators and large areas can be acquired. NaI(Tl) is

hygroscopic and brittle, so it must be hermetically sealed and handled

gently. The high light yield of NaI(Tl) gives it the best energy

resolution of all scintillators: approximately 11 keV full width at

half maximum (7d HM) at 60 keV (Ballintine 1981). Sodium doped CsI has

had more use as an anticoincidence shield, either as a large shield

^I

with a "well" for the primary detector or as the shield crystal in a

"phoswich" detector (Peterson 1975)• Bismuth germanate, (Bi 4Ge 30 12 or

"BGO"),is a recent addition to the scintillator family. Its most

y
Light Output

(relative to NaI) 100

Density ( gm/cm3 ) 3.67

Decay Time Constant

(nanoseconds) 230

Wavelength of Maximum

Emission ( nanometers) 415

Index of Refraction 1.85

Thickness Required to

Attenuate 500 keV photons

by 90% (centimeters) 6.7

e

85 10

4.52 7.13

630 300

420 480

1.8 2.2

5.3 2.4

is

0-47 ^^

Table 2.1•

Properties of Commonly Used Scintillators

INORGANIC

NaI(TI) CsI(Na) B14Ge3012

ORGANIC

Plastic

(NE102)

28

1.03

2.4

423

1.6

23

*Adapted from Farukhi (1978)

^^ I

I;
17 T

outstanding characteristic is its photon stopping power due to a high

effective atomic number and density (see section 3.2).

2.2.1 Phoswich Detectors

A "phoswich" detector (phosphor sandwich) consists of two dissimilar scinti.11ators such as NaI and CsI optically coupled together

and viewed by a single photomultiplier (Fig. 2.1). It takes advantage

of the different luminescence decay times of the two crystals so tha t a

decay time discriminator can differentiate between pulses originating

in the NaI or CsI. It is an efficient background rejection scheme

providing "4n" ster^adians shielding of penetrating radiation as well

as removing the phototube mass (and its associated background production) away from the primary detector.

To achieve a low background in a phoswich configuration the shield

crystal (usually CsI(Na)) must be quite thick (-3 cm) while the primary

crystal thickness is determined by the energy range of interest. The

HERO A4 experiment (Matteson,1978) contained 3 phoswich detectors

optimized for 3 energy ranges (10-200 keV, 80-2000 keV, 300-10,000 keV)

with effective areas of 220 cm 2 , 170 cm and 120 cm 2 , respectively. To

date, this instrument has been the most sensitive broad band hard X-ray

detector in orbit.

Larger area phoswich systems have been flown from balloons by

groups from UCSD, Max Planck Institute and here at MIT. The Balloon

Group at the Center for Space Research in collaboration with the Cosmic

Ray Working Group at the University of Leiden, designed, constructed

and flew a 1500 cm phoswich array. This instrument (see Section 4 and

Ballintine 1981) was the most sensitive (per uiit time) astronomical

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ALUM SUPPORT STRUCTURE 

Mu METAL SHIELD 

ALUM CASE PRE AMPliFIE

Pb SHIE

(MODULE 

~ 

OPTICAL COUPLING 

QUARTZ - WINDOW 

PHOTOMUL TIPUER .RCA 4830-12 

Be WINDOW ( I mm THICK) 

Nol (TI) PHOSPHOR (96mm It 96mm It 3mm) 

Csl(No) PHOSPHER (96mm)(96mmfront ,84 mm.Dfo. Rea

Figure 2.1 Phoswich assembly. This sectional view shows the construction 

of an individual phoswich detector (16 were used in the MIT 

1500 cm 2 instrument last flown in May 1980). See section 

2.2.1 for a description of its operation. 
19

hard X-ray detector until its demise in May 1980.

2.3 Solid State Detectors

The application of silicon and germanium solid state detectors to

X-ray astronomy has occurred only recently due to the slow progress in

growing large volume crystals. Such a solid state detector can be

thought of as a solid ionization chamber in which incoming radiation

produces energetic electrons that ionize the lattice and raise electrons from the valance band to the conduction band. These conduction

electrons (and the corresponding holes) diffuse under an applied

electric field to the electrodes and are collected as a charge pulse

I`I

proportional to the energy deposited in the crystal. Due to the lack

of competing energy loss mechanisms and the small band gap, most of the

deposited energy goes into electron-hole production, giving many charge

carriers and a correspondingly smaller relative statistical variance

(see Section 3.1.2). This small variance in generated charge gives the

solid state detector its excellent energy resoluti.on (0.5 keV FWHM at

60 keV, Knoll 1979 pg. 500), which is limited by the preamplifier

electronics.

The advantage of good energy resolution is the ability to resolve

possible features in an energy spectrum that would be lost in the

continuum using a broad band instrument such as NaI. Possible features

expected in astronomical observations include nuclear gamma ray lines

from radioactive decay of elements produced in supernova explosions, a

positron annihilation line from the Galactic Center, and cyclotron

lines from magnetized neutron stars. Though it is possible to de-

^

r

convolve the monoenergetic response of a broad band detector, the

a

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NAIV
k' R

,

20

stability and the number of photons required is prohibitive in most

astronomical hard x-ray observations.

The properties of various solid state detectors are listed in

Table 2.2. Germanium has been the usual choice for hard X-ray astronomy because of the materials available in large sizes, it has the

highest photon stopping power. The most important limitation of

germanium detectors is the requirement that they be operated at liquid

nitrogen temperatures to reduce thermally-generated leakage current.

This involves massive and bulky cryogenic equipment and, on satellites,

a limited operating life. The gamma ray spectrometer on HERO 3

(Mahoney at _al. 190) made with 4 Ge detectors 8 (effective area, 64 cm 2

and sensitive from 45-10,000 keV) hac 3 planned operating life of 1

year. Of all the other detector materials listed in Table 2.2, only

HgI 2 combines a high photon stopping power with satisfactory room

temperature operating ability. Room temperature solid stata detectors

i such as HgI 2 work because of a large bandgap that suppresses thermally

generated leakage current. Cadmium telluride's smaller bandgap results

in larger room temperature leakage currents that degrade the energy

resolution. Gallium arsenide (GaAs) has a bandgap similar to CdTe

(1.43eV) and therefore suffers from the same leakage current problems

at room temperature. The development of GaAs as an X-ray detector is

not as advanced as is the development of the materials listed in Table

2.2, with a maximum detector thickness limited to 100 microns.

Silicon can be made to work at room temperature with suitably doped

con.acts, -et its hard X-ray stopping power is insufficient. Since the

( development of "flight qualified" HgI 2 detector system is a major focus
'i

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21 f^

Table 2.2"

Properties of Commonly Used Solid State Detectors

Si Ge HgI2 CdTe

(300°K) (77°K) (300°K) (300°K)

Atomic Number 14 32 80-53 48-52

Density (gm/cm 3 ) 2.3 5.3 6.4 6.0

Band Gap (eV) 1.12 0.66 2.13 1.47

Energy Required per

electron-hole pair (2V) 3.6 3.0 4.2 4.4

Thickness Required to

Attenuate 60 keV by

90% (mm) 57 2.1 0.6 o.6

*Adapted from 'kited (1979)
22

of this thesis, extensive details of mercuric iodide as a detector are

discussed fully in Section 3.

a

V
fi 23

K

3.0 DEVELOPMENT OF NEW TYPES OF D"."CTORS FOR ASTRONOMICAL WORK

The ideal detector for hard X-ray astronomy would have excellent

resolution, high efficiency, large area and be lightweight (including

support equipment, i.e. cryogenics). This "dream" detector does not

exist though most systems in use satisfy some of the requirements. The

standard detector for astronomical hard X-rays has been large area NaI

scintillation crystals that operate at ambient temperatures but with

limited energy resolution (17 0% FWHM at 60 keV). The semiconductor

detectors such as germanium have recently achieved respectable areas

(64• cm 2 , Mahoney at al. 1980) yet their lower efficiencies require

thicker detectors which increase the background counting rate. Also the

requirement for cryogenic temperatures has limited their satellite

lifetime to -1 year.

3.1 Mercuric Iodide Detectors

Mercuric iodide (HgI 2 ) solid state X-ray detectors have many

unique advantages in their application to X-ray astronomy. HgI2

detectors operate at room temperature so that bulky and massive cryogenic coolers are not required. HgI 2 detectors also combine a high

quantum efficiency with good energy resolution. Because of the excellent quantum efficiency (due to the high effective atomic number of

HgI 2 , see fig. 3.1) background contributions which scale as the thickness of the detector

lower than for other

power. At this time

is only 1 cm2 , so it

erties are well matc:

(e.g. those due to neutrons and gamma rays) are

hard X-ray detectors with equivalent stopping

the largest area single detector that can be made

is far from being ideal. Nonetheless, its propied to a scheme that increases the effective area

^^.
Ji

^yrvwP,z[ •.•

f

24

C . ` (,iy

LINEAR ATTENUATION COEFFICIENT

104

n 1

Mercuric

Iodide

photoelectric _ _ __ Germanium

pair

production

\ _ , Compton

,-^ 103

i

102

0

101

0

c.^

0

^-' 100

a

m

.^ 101

102 1

10 10 2

103 104 105

Energy (keU)

Figure 3.1 Linear attenuation coefficient vs. energy for mercuric iodide

and germanium. For energies above 33 keV, the thickness of a

detector required to stop a given percentage of incident

photons is much less for mercuric iodide than germanium.
r+)4 .

1

25

`

using a Bragg diffraction concentrator (described in more detail in

Section 5.2).

I 3.1.1 Detector Production and Characterization

— i

Mercuric iodide is a recent development in the search for high Z

room temperature solid state detectors. The first real energy spectrum

presented using HgI 2 as a detector was of the radioactive ource 241 Am

0

in 1972 ( Malin 1972). In the last ten years many international groups

have been growing HgI 2 crystals but the major efforts have been at

4

EG&G's Department of Energy laboratories in Santa Barbara and at the

r.

ft Hebrew University of Jerusalem, Israel ( Whited 1979; Van den Berg

UI

fw i 1978). Detectors studied at MIT have either been grown here or acquir1

^.i

ed from Santa Barbara.

Mercuric iodide must be grown using vapor growth techniques

because it undergoes a destructive phase change at 127°C when the room

temperature red colored tetragonal structure transforms reversibly to

the yellow orthorhombric structure. This temperature is below the

250°C melting point of HgI 2 , preventing growth through solidification.

'; f The most widely used growth technique is the temperature gradient

n

i

reversal process ( Van den Berg 1977) producing crystals of -20 cm 3 from

7

which thin ( . 05 cm x 1 cm 2 ) detectors are then cleaved.

k

Recently, a new method of growing crystals has been developed

using an organic polymer ( Faile 1981). This method produces thin

Lcrystals that can be directly fabricated into detectors thereby bypassing the cleaving process that might damage the soft HgI 2 . Detect-

-

ors fabricated from crystals grown with the Fails organic polymer

method show excellent charge transport properties. The reasons for

4

:J
I

}

(40)'I

26 41

such good performance have not yet been established. One hypothesis

(Burger 1982) is that the polymer affects the stoichiometry of the

final crystal (ratio of iodine to mercury) decreasing the number of

trapping centers. At this time crystals grown with the Faile process

are too thin (-50-100 microns) to be of use in hard X-ray astronomy.

However, as low energy X-ray (1-10 keV) astronomical detectors, they

show great promise.

3.1.2 Physics of Detectors

Radiation interacting with semiconductor detectors converts most

of its energy into electron-hole pairs by exciting electrons into the

conduction band of the crystal. This charge is collected by applying a

constant electric field across the detector. The resulting current

pulse is proportional to the energy deposited in the crystal. The

energy required to produce an electron-hole pair in HgI 2 and Ge is 4.2

eV and 3.0 eV, respectively (e.g. a 42 keV X-ray will produce an

average of 10,000 electron-hole pairs in HgI 2 ). To see these small

charge pulses, the statistical variation in the leakage (DC) current

and therefore the current itself must be small. For example, a 500

picoampere current corresponds to 3120 electrons per microsecond with a

statistical fluctuation of (3150) 1/2 or 55 electrons. If the charge

generated by the 42 keV X-ray is collected in 1 microsecond then the

variation in the collected charge is dominated by the statistical

variation in the generated charge and not the statistical variation of

the leakage current. Because of its relatively large bandgap (2.13 eV)

the HgI 2 leakage current at toom temperature is typically in the range

of 10-500 picoamperes. This allows room temperature operation in
27

contrast to germanium (bandgap .66 eV) which requires liquid nitrogon

temperatures to achieve these low leakage currents.

In the absence of quantum statistics, electronic noise, and

variations in the efficiency with which charge is collected, discrete

mono-energetic events should give a delta function line spectnim. In

actuality, all these effects act to broaden the line. Due to the

stochastic nature of the charge generation process, the variance of the

number of electron hole pairs is:

a FE/e

where o 2 is the variance in the number of pairs generated

c is the energy required to create an electron-hole pair

F is the Fano factor

E is the energy deposited in the crystal

The Fano factor is introduced as an adjustment factor to relate the

actual variance with the Poisson predicted variance (Knoll 1979, pg•

373)• It is determined empirically, although qualitative models do

exist that attempt to explain the non-unity value (Klein 1968). In

semiconductor detectors and gas counters, the statistical variation is

smaller than expected if the lattice ionization events are uncorrelated

(Poisson, F-1). Of course, the ionization events are correlated

(electrons are excited into the conduction band by other electrons) and

it is the random nature of competing energy loss mechanisms (phonons)

that determine the variance. The smaller the Fano factor, the better

the energy resolution. Upper limits to the Fano factor are determined

I
r

28 r

by accounting for all other line broadening mechanisms and assuming the

residual line width is due to charge generation statistics alone.

Recently, a new lower value for the HgI 2 Fano factor of .19*-.03 was

determined at the balloon lab at MIT ( Ricker et al. 1982a; see appendix). This continued the trend of decreasing the measured Fano factor

that was also observed in silicon and germanium as crystal growing and

fabrication techniques improved (Knoll 1979, pg. 374). The full width

at half maximum ( FWHM), aE, due to charge generation statistics is

given by:

AE-2.355 (FeE)1/2

For F-.19, the formal value of AE in HgI 2 is .5 keV FWHM at 60 keV.

This energy resolution represents a theoretical intrinsic limit which

has yet to be observed.

A further limit to energy resolution is imposed by the fact that

the efficiency of charge collection does not equal 100 percent for

typical detectors of -500 microns thickness. The mean free drift

length, a , of an electron (hole) can be expressed as:

a-PTE

where 3 ; is the electric field

V is the electron (hole) mobility

T is the mean free time before the electron (hole) is

"trapped" (i.e. removed from the conduction band by

defects or impurities in the lattice.)

^e

The drift length characterizes the mean distance a charge carrier

will travel before being lost or " trapped" by the lattice. For good

collection efficiency these lengths should be large compared to the

0

r ;'p\ 0
29

I

thickness of the detector. Expressed in terms of a e and a h , the charge

collection efficiency, n, as a function of the position x of the

initial ionization in the crystal is given by:

X

n(x)

s (

1 -exP(-x/a * as (1—exP(-d*x)/ah)

where d is the thickness of the detector (Day 1967). Table 3.1 lists

the charge transport parameters of good HgI 2 crystals, as well as

germanium crystal parameters for comparison. Two points must be made.

First, the electric field cannot be increased indefinitely (thereby

increasing a) because electrical breakdown of the crystal at high

voltages will result the destruction of the detector. Typical

voltages used in the lab are -1000 V for a 500 micron detector, giving

electric fields of 2x10 4 V/cm. Second, the large ratio of the AT

product for electrons and holes (-100) creates an asymmetry in the

collection ability of the detector. If the incident radiation ionization sites are distributed uniformly throughout the detectors (high

energy gamma rays) then the spread in pulse heights for monoenergetic

radiation is the difference between 

Amax 

and Amin' If the radiation

interacts closer to one electrode (low energy X-rays) the polarity of

the electric field should be applied in such a way as to make the

distance the holes must travel as small as possible. Fig. 3.2 shows

the result of a simulation of a 390 micron thick HgI 2 detector irradiated with monoenergetic beams of 20, 60 and 150 keV photons. The

charge collection parameters of Table 3.1 and an applied electric field

of 16000 V/cm were used as well as a 2 keV FWHM Gaussian folding
I T C+)4 ^

30

Table 3.1•

Charge Transport Parameters for HgI 2 at Room Temperature

and Ce at Liquid Nitrogen Temperature

HgI 2 (300°K) Ce (77°K)

Mobility, u in cm2/volt-sec

electrons

holes

Mobility-]lifetime product,

PT, in cm /volt

electrons

holes

Maximum electric field

90 3.6x104

3 4.2x10

5x10

-4

>5x10-2

5X10 -

6

>5x10-

50,000 V/cm 2000 V/cm

*Adapted from Iwancyzlc (1981)

F

G,

F-

•1
r

I: 750 microns

m ORIGINAL PAGE 19

^c

OF POOR QUALITY

Z

(61 ov

.w

c

C

m

c

C

0 20 40 60 80 100 120

Energy (keV)

140 160 160 200

to

Y

390 microns

Ny

c

7

lO^ Q

v

T

c

c

CD

c

0 20 40 60 80 100 120 140 160 180 200

Energy (keV)

Figure 3.2 Mercuric iodide charge collection efficiency simulation.

Shown are the resulting simulated energy spectra for two

different thicknesses of detector: a) 390 microns and b) 750

microns. Both detectors were irradiated with three different

beams of monoenergetic X-rays with energies of 20, 60, and

150 keV. The "low energy tails" of the higher energy peaks

is due to the small mean free path of the holes resulting in

incomplete charge collection for interactions far from the

negative electrode.
32

^'

T

routine to simulate the electronic noise. The high energy interactions

have the characteristic "low energy tail" due to incomplete charge

collection. For high resolution spectrometers, hole trapping is the

major limitation of detector thickness. A thicker (fig-3.2b) and thus

more efficient detector of X-rays must sacrifice some of the energy

resolution of the detected radiation. The energy at which this tradeoff occurs is a function of the quality of the detector.

Besides charge generation and charge collection efficiency, the

electronic noise introduced by the preamplifier can broaden monoenergetic line features. At low energies (<6 keV) electronic noise is

the dominant noise source and much progress has been made at room

temperature with low noise preamplifiers (AE electronic"225 eV FWHM)

(Iwanczyk, 1981). At higher energies, preamplifiers with an electronic

noise contribution of 1 keV FWHM are common and the line width is

dominated by the charge collection efficiency, especially since the

uncorrelated noise sources add in quadrature.

3 . 1.3 Experimental Development for Astronomical Use

Mercuric iodide has been under development for ten years, and the

balloon lab at MIT has had a working experience with these detectors

since 1979. The first use of HgI 2 in an astronomical instrument was

small array of 12 detectors with a combined area of 8 cm constructed

here at the balloon lab at MIT and flown in a small balloon-borne

gondola in June, 1980. Though the background counting rate at float

was high because of a less than optimum design, the instrument was able

to set an upper limit to the hard X-ray flux of Cygnus X-1 in the "high

state" (Ogawara at al 1982; see Appendix).

i;.

k;

4Y

s u.
Im

33

r

HgI 2 detectors can still be considered "experimental" in that they

cannot be acquired commercially and each detector is in some way

unique. In work carried out by the MIT balloon group, the HgI 2 crystals have been subjected to extensive field testing. They have been

exposed to extreme variations in temperature, pressure, humidity,

voltage, ligh` level and handling. As might be expected with a.n

experimental material, HgI 2 did not always pass with flying colors.

However, from this ex perience, general procedures for handling and

environmental limitations have been developed to extend the operating

life and reproducibility of results necessary for a field detector.

Because of its relative softness, the direct handling of HgI2

crystals that have been cleaved from the large 20 cm 3 parent crystal

must be minimized. The "thick" crystals (-300-1000 microns) used at

the balloon lab for hard X-rays have all been acquired from EG&G, Santa

Barbara. The crystals are fabricated into detectors with evaporated

palladium electrodes, mounted on ceramic or plastic substrates and

encapsulated with Dow Corning 3140 RTV. The typical area of these

detectors is 1 cm with energy resolution ranging from 2.8 to 15 keV

FWHM at 60 keV. A typical energy resolution would be 4 keV at 60 keV.

The area limitation requires many detectors in an array to achieve an

astronomically useful collecting area (-100 cm 2). Such an array would

also require 100 measurement chains (100 preamps + voltage amplifiers)

as well as existing crystal growers to produce high resolution uniform

detectors in quantity. Until the detector area can be made larger,

some method of concentrating X-rays must be used to increase the

effective area. (An X-ray concentrating technique using Bragg diffracI

J
T,•r MR^3I _..

+'

^e

34 1'

U

tion is described in Section 5.)

In the laboratory the detectors must be kept clean and dry. Small

amounts of water and contaminants (dirt, solder flux, etc.) create

current leakage paths that increase the overall RMS noise and small

intermittent current pulses similar to a relaxation oscillator.

Encapsulated detectors should not be exposed to vacuum because trapped

air pockets will er tnd and strain the crystalline lattice, thus

ruining the charge transport properties.

The temperature effects are striking for some crystals and unnoticeable with others. The mobility increases with lower temperatures

for both electron and holes (u a -T '9; Ph- T _

!

.7 ; Minder 1974) but the

mean lifetime of an electron or hole before it is trapped by the

lattice decreases with decreasing temperature. These competing processes determine a temperature where the resolution is best ( i.e. the mean

free path, is maximized) and is a function of the quality and

purity of the crystal. Empirically, this temperature is between -20°C

and 20 0 C,varying with detector (Ricker 1982a, Grunsfeld 1980). These

temperatures are easily achieved at balloon altitudes or on satellites

by passive ( radiative) or active (thermoelectric) techniques. Higher

temperatures decrease the mobility and also produce a prohibitive

leakage current which can saturate DC coupled preamplifiers.

The bias voltage across the HgI 2 must be applied slowly to avoid

large induced currents. After the bias voltage is applied, the deteci

tor "warms up" over a time period of -1 minute to 1 hour depending on

the crystal. The pulse height for a given incident energy increases

and the resolution improves. This can be interpreted as an increase in
t T^v MR X - - ^^

the depletion region as the deep traps (energy levels in the forbidden

region caused by impurities) are emptied by the large elec*ric field.

Resolution usually improves with increasing electric field until the

leakage current becomes too large. Also, high electric field strengths

increase the chance for surface breakdown which intermittently injects

current pulses into the preamp and shows up as low energy events. The

frequency of occurrence of this "current injection noise" increases

with applied voltage.

Many HgI 2 detectors perished in the balloon lab due to different

types of abuse. However, by minimizing any handling that would stress

the crystal and by keeping the detector in a clean, cool and dry

assembly, it was possible to produce a stable and dependable operating

HgI 2 X-ray detector.

3.2 Bismuth Germanate as Active Shielding

A recent development in scintillation crystals is bismuth

germanate (Bi 4Ge 30 12 or "BGO") (Nestor 1975). BGO is a clear, inert,

non-hygroscopic scintillator with a light output -10% that of NaI.

Because it contains bismuth (Z-83) and has a density of 7.1 gm/cm 3 , it

i

has a large attenuation coefficient for photons (Fig. 3.3). These

characteristics make it ideal as an anticoincidence shield where

stopping power is more important than energy resolution.

The stopping power of BGO is at least a factor of -2 greater than

CsI, another scintillator commonly used as an anticoincidence shield

(Fig. 3.3). The shield geometries can be more compact, enabling the use

of smaller photomultipliers. Another advantage of a smaller size is a

[ decrease in the count rate, leading to reduced dead times. Since BGO

f
.-. 103

i F—

v

102

a^

U

Q) 101

C

v

a

100

-^ 101

I 

Bismuth

Germanate

Cesium

Iodide

photoelectric

_ _ _ _ _ Lompton

pair

production

J'

38 R

ORIGINAL PAGE 19

OF POOR QUALITY

LINEAR ATTENUATION COEFFICIENT

10 

Pj

10 10

1

102 103 104 10t

Energy (keV)

Figure 3.3 Linear attenuation coefficient vs. energy for bismuth

germanate and for cesium iodide, a more commonly used

anticoincidence shield.
t ^^

37

i

is non-hygroscopic, it does not need to be hermetically sealed. This

allows a simpler, less expensive container that can be modified and

repaired in the laboratory. Unlike NaI, BGO is very hard and rugged

and can be machined like glass. It is also resistant to thermal

shocks.

Even with all these advantageous properties, BGO will probably not

be used as a primary detector for astronomical hard X-rays because of

its low light output (as compared to NaI) which determines its energy

resolution (30 keV FWHM at 60 keV). What is more important for use as

an anticoincidence shield is the minimum detecishle energy above the

photomultiplier noise so that the anticoincidence lower level discriminator can be set as low as possible. In the laboratory the BGO

crystal/photomulti.plier combination described in Section 6 can detect 6

keV X-rays. The light output , f BGO increases with decreasing temperatures (Filtingsrud 1979) b y 25% from 30°C to 0°C. This fact combined

with the lower noise characteristics of photomultiplier tubes at low

temperatures argues for a slightly cooled BGO shield (0°C), easily

achieved at balloon altitudes.

c^

JAI
39 `

4.0 RESULTS FROM MAY 1980 BALLOON FLIGHT OF A LARGE

AREA (1500 cm 2 ) PHOSWICH DETECTOR

Until its demise on May 8, 1980, one of the most sensitive hard

X-ray telescopes existing was the balloon-borne 1500 cm phoswich array

constructed at MIT by the Balloon Group in collaboration with the

Cosmic Ray Working Group, University of Leiden, the Netherlands.

A detailed description of the orientation system, phoswich detectors and associated electronics has previously been reported (T.

Ballintine 1981) and a short summary follows. Observations of two

celestial objects using this telescope will be reported: the COS A

gamma ray source CG 135 + 1 and a drift scan across the Galactic Center.

These observations have been analyzed as part of the investigations

leading to this thesis. The exceptional sensitivity of these observations motivates the search for a replacement instrument, an example

of which is described in later sections of this thesis.

4.1 Instrument

The gondola flown on May 7 and 8, 1980, basically consisted of a

detector system with associated electronics and an orientation system

that could point the detectors at any celestial hard X-ray source above

the horizon. The data telemetered back to ground included: detector

events (energy, time, pulse shape); redundant aspect information of the

azimuth and elevation of pointed detectors vs. time measured by various

sensors and detectors (including a charge-coupled device (CCD) television camera aligned with the X-ray axis and a sun sensor); and

general housekeeping information (voltages, currents, temperatures,

etc.).
I

39

The detector assembly consisted of an array of passively collimated phoswich scintillation crystals sensitive to 20-240 keV X-rays.

The sixteen 92 cm square detectors were divided into 2 groups of 8,

module A and module B, each surrounded by plastic scintillators to veto

high energy charge particle events. The phoswich assembly (Fig. 2.1)

had as the primary X-ray detector a 9.6 cm x 9.6 em x 0.3 em thallium

activated NaI crystal bonded to a 9.6 cm x 9.6 cm x 3.8 cm sodium

activated CsI crystal. The crystal was heremetically sealed in an

aluminum case with a 1 mm beryllium entrance window and a 3 mm quartz

window optically coupled to a photomultiplier tube. Only events

occurring in the NaI crystal (as determined by the risetime) were

considered valid events.

In front of the detector were a pair of slat collimators that

defined the field of view. Each collimator consisted of flat parallel

plates made up of layered metal. The horizontal collimator restricted

the field of view in elevation and the vertical collimator restricted

the field of view in azimuth. The slat collimators had a triangular

response to off axis X-rays. Module A had a field of view of 3 000

FWHM while Module B had 3° elevation x 1.5' azimuth FWHM field of view.

The finite geometric size of the collimators reduced the sensitive area

of 1475 cm to 1125 cm effective area for on-axis X-rays. The vertical collimators were mounted in such a way as to be rocked right and

left by 4.5 0 relative to the nominal viewing direction. This allowed a

background measurement to the right and left of a faint X-ray source

without moving the entire gondola. For faint sources where half the

observation was spent on background, the collimators pointed two

V
Tl

40

f '

d

minutes on the source, two minutes to the left of the source, two

minutes source, two minutes right, and the cycle repeated. This

"chopping" of the field of view reduced systematic errors in background

subtraction due to long term variations in the background courting

rate. For drift scans the collimators were left in the neutral (onaxis) position, and the source was allowed to drift through the field

I

of view as the earth rotated.

The aspect determination of the gondola is important in determining the exposure of the detectors to a celestial source. Once the

position on the celestial sphere of the collimator axis and the X-ray

source itself is determined, then the collimator transmission function,

r

T, can be calculated usina:

cos(II source) (3-III E1 I) (F- Az Az I)

Tsource axis source axis

3xF

where (Az source' El source ) and (Az axis' axis) are the azimuth and

elevation of the X-ray source and collimator axis respectively. F is a

constant equal to 3 for module A and 1.5 for module B.

The method of establishing the X-ray axis of the collimators

during flight are described in detail elsewhere (Ballintine 1981). The

aspect during the galactic center drift scan came from two independent

sources: the "coarse" aspect based on the orientation system (angle of

elevation axle, azimuth of magnetometer, roll and pitch sensors); and a

direct measurement with the "star TV", a CCD camera whose optical axis

is calibrated against the X-ray collimator a_eis on the ground. The

f, star TV has an accuracy of -.1 0 (calibration error) though the coverage

P l

it 1

._.- )
^4 i D

n

9.N

41

of the drift scan is incomplete due to lack of bright stars in the

field of view of the CCD (3 0 x4°). The coarse aspect is complete

throughout the flight though its accuracy is only 7-3 0. The CG 135+01

observation occurred during the day so the coarse aspect was used along

with an azimuth-sensitive sun sensor accurate to .1 0 . Based on the

collimator transmission function, an error of 0.3 degrees in either

azimuth or elevation is a 10% error in intensity for module A and a 20%

error in intensity for module B if the error is in the azimuthal

direction. These errors in absolute intensity are small compared to

the statistically determined errors of faint sources such as CG135+01

and the Galactic Center.

4.2 The 1980 Flight

The gondola was launched at 00:30 UT on May 8, 1980. The 1.2x106

m 3 balloon from which it was suspended reached an average float altitude of 38 km (-4 gm/cm 2 atmospheric column density). Sources observed

during the flight included NRAO 140, Her R-1, Cyg R-1, 30273, Crab

Nebula, CG 135+1 and the Galactic Center. Results on the first five

have been reported elsewhere (Scheepmaker 1981; Ricker 1981). Two

drift scans of the Galactic Center were made during the period 08:25 UT

to 9:21 UT on May 8, 1980. The X-ray detectors were pointed at -1800

azimuth and 29 0 elevation (gondola position: 31.3 0 N latitude and 94.30W

longitude). The source drifted across the field of view at -.2°/min.

The 2 hour observation of CG 135 + 1 began at 13.06 UT on May 8. The

collimators were pointed at the most recent COS B error circle

(Swanenburg 1981) and the chopping collimator technique was used. The

average atmospheric column density to the source was 9.2 gm/cm 2 for the
K

42

Galactic Center and 4.25 gm/cm for CG 135+01.

4.3 Detection In Hard X-rays of the COS-B

Gamma Ray Source CG 135+1

With the publication of the second COS-B catalogue (Swanenburg et

al. 1981) only 4 of the 25 detected gamma ray (E>100 MeV) sources had

been positively identified with sources at other wavelengths. Two are

associated with the Crab and Vela pulsars through their timing signatures and the other two are identified by positional coincidences with

the quasar 30273 and the cloud complex o Oph. Because of the large

error circles of COS-B (radius -1 0 ) and the fact that most of the COS-B

sources lie on the galactic plane, the identification of additional

sources with lower energy counterparts is ambiguous at best.

The two "best" candidate sources that have been proposed as the

lower energy counterparts to CG135+1 are the QSO 0241+62 discovered by

Apparao et al. (1978) in soft X-rays and the radio star GT0236+61

(LSI+61 0 303) discovered by Gregory and Taylor (1978) in a survey of

variable radio sources in the Galactic plane. LSI+61 0 303 is also a

soft X-ray source and has been observed by the Einstein HRI and IPC

instruments (Bignami 1980). Recently a 26.5 day period was observed in

the radio emission of the radio star (Taylor 1982) and has been interpreted as a binary period of a compact object orbiting a BO-BO.5 main

sequence star. Since the only identified 100 MeV gamma ray sources are

the Crab and Vela pulsars, the QSO 30273, and the cloud complex P Oph,

the QSO 0241+62 and the radio star system LSI +61 0 303 are both strong

contenders for the identification with CG 135+01. Fig. 4.1 shows the

location of these 2 candidates with respect to the COS-B error circle.
---------•

I

I

I

FWHM MOD. A

I

I

I

I

I

I

I

2

1

U')

W

W 0

Of

V

Ld

71

—.l

43

ORIGINAL PAGc 1j

OF POOR QUALITY

2•J

l

I

,I

I

4

I I

^.

I

-2 FWHM MOO. B

-3 ^. y q I_ I I F

M cal m N m

DEGREES

1 LSI 61 303

2 OSO 0241+62

CG 135+01 ERROR CIRCLE

Figure 4.1 The sky near the COS B gamma-ray source CG 135+01. Shown are

the positions of two soft x-ray sources proposed as X-ray

counterparts, as well as the COS B 901,16 confidence circle for

the gamma-ray source. Also shown is the average field of

view of the slat collimators over the 1500 cm = phoswich

array. The FWHM of the triangular response is plotted for

both Module A (3 0 x3°) and Module B (1.5°x3°).
RJ (^+D4

t

44

Note that the QSO 0241+62 is just outside the reported 904M confidence

contour (Swanenburg 1981).

The observation on May 8 detected a hard X-ray source with a 4.1

sigma level of significance from the region of the sky containing the

gamma ray source C0135+01. Both candidates, the QSO 0241+62 and the

radio star LSI+61 0 303 were in the field of view of the detectors

(Module A: 3°x3° FWHM; Module B: 1.5°x3° FWHM) and therefore it

cannot be positively determined from which source the X-rays were

detected (Fig. 4.1). The intensity can be determined assuming it was

one source or the other and that intensity can be compared with previous observations carried out by other observers in different (sometimes overlapping) energy bands.

Figure 4.2 shows the hard X-ray spectrum plotted two ways. In 4.2

the source of emission was assumed to be the QSO 0241+62, which was on

the average -1.6 0 away from the center of the field of view; while in

4.2b the source assumed was LSI+61°303, which was 0.65° away from the

center. Thus, due to aspect corrections, the intensity for the source

in (a) is greater than in (b). This intensity difference is due to the

triangular response of the slat collimators. Plotted along with these

measurements are results reported in the literature from previous

observations of this region (OSO 8, Worrell et al. 1980; OSO 7, Marachi

et al. 1978; SAS-3 Apparao et al. 1978; MISO, Perrotti et al. 1980; MPC

on HERO 2, Halpern 1982).

As can be seen in Fig. 4.2a, if the hard X-ray source is the QSO

0241+62 then the intensity in the range 20-50 keV (overlap with OSO 8)

has increased by a factor of -5 over a timescale of -2 years. This

n
45

z

U

m

N

N^

U

C

C

OJ

mz

U

m

N

NU

N

C

O

J

O

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i+a

'yA

ORIGINAL PAGE IS

OF POOR. QUALITY

10-2

CG135+1:

candidate

1

:m-'

I D MPC ^"

r OSO 7 h T 

C 10 5 C — OSO 8

- SAS 3

-4- MIT/L EICEN

to 

1ma

a lm _

CG135-1:

LS , a

candidate

lm' -

— MPC ±I

lm" g ^ t OS^u / —^^ I ^ I

-- M I SO

i

-4- MIT/LEIOEN

: ma :a4 lmZ :m3

Energy (keV)

Figure 4.2 X-ray spectra for two possible candidates for the gamma-ray

source CG 135+01. The MIT/Leiden points are plotted assuming

the source is either the QSO 0241+62 [upper panel (A)I or the

radio star LSI+61 0 303 [lower panel (B)]. For further details

see text.

la4 :aZ 123

Energy (keV)

1114)
p

 n - , ,,.

..

0

46

QSO is known to be highly variable in both the radio and optical bands

(Feldman 1980). There is also evidence for variability in X-rays by

approximately this same factor based on the OSO 7, OSO 8 and SAS-3 data

points. The box shown for the MPC on the Einstein Observatory covers

the range of fluxes measured in 4 separate observations over a period

of 7 months (Feb. to July, 1979).

Plotted in Figure 4.2b are the MIT/Leiden data points assuming the

observed hard X-ray source was the radio star LSI+61°303• Of the

previous results plotted for this source, only the MPC instrument has a

small enough field of view to avoid contamination from the QSO. The

OSO-7 data points ive been corrected for aspect assuming the radio

star was the soft X-ray source although the 1 sigma error box does not

contain the radio star (Markert 1982). Based on the factor of -20

lower flux measured by the MPC, the OSO-7 points are more consistent

with the QSO interpretation of their origin. The MISO points also

suffer from possible source confusion (collimators: 3'x3° FWHM). The

points are plotted only for the radio star because the QSO lies just

outside their 1-sigma error region determined from their drift scan

results. An extrapolation of the MIT/Leiden data points is inconsistent with the MISO'points in the .2-1 MeV range by a factor of 12.

Also the MIT/Leiden data do not require a spectral break in the power

law extrapolation to explain the COS-B gamma ray intensity as do the

MISO points (Perotti 1980). A low energy extrapolation would also seem

to be inconsistent with the very low flux measured by the MPC.

The parameters of a power law fit to the hard X-ray data are shown

in Table 4.1 for both assumed sources. The hard X-ray luminosities

rX1117p A-
`, J^

47

ORIGINAL PAGE 19

OF POOR QUALITY

CAWr

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48

rx

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i

derived from this flux are similar to the measured gamma ray (100 MeV1 GeV) and soft X-ray (2-10 keV) luminosities for the QSO. Other

quasars (e.g. 3C"'7 >) also exhibit luminosities in this range throughout

the spectra from radio to gamma rays. Models accounting for this

energy generation in quasars are usually based on accretion disks

around supermassive black holes (Eardley et al. 1978) where the R- and

gamma-rays are created by inverse Compton scattering of photons off

relativistic electrons.

A similar process (inverse Compton scattering) for producing high

energy photons in the LSI+61 0 303 system is proposed by Marachi and

Treves (1981). Their model proposes a young pulsar losing energy in a

relativistic wind at a rate of 10 37 ergs/sec. The wind is randomized

in a shock region at the interaction boundary of the pulsar and stellar

winds. The stellar photons are inverse Compton scattered to R- and

gamma-ray energies. This model produces relativistic electrons continuously while competing accretion disk models can modulate the production to better explain the parodic synchrotron radio emission (Taylor

and Gregory 1982). The radio emission peaks between the phases 0.3 to

0.8 of the 26.52 day period (phase 0.0-JD 2443366.775) and so should

the hard X-ray and gamma rays if particle acceleration is linked to an

accretion process. The MIT/Leiden observations occurred at the radio

phase -0.75. The MISO observation occured at the radio phase -0.01,

inconsistent with this interpretation of the radio emission. The MISO

derived luminosity, L (.2- 1MeV) -10 37 erg/sec, also pushes the Marachi

and Treves model to its limit requiring 100% conversion of rotational

energy into 1 GeV electrons. There is also a large discrepancy with
the DLpC soft X-ray luminosity.

Althoua:, this detection of CC 135+1 in the energy range 20-180 keV

cannot with certainty determine the low energy counterpart to the gamma

ray source, it does extend the knowledge of this source region into

another spectral decade. Whether it shows variability of the quasar

0241+62 or establishes a hard X-ray flux for LSI+61°303, this detection

should be helpful in unravelling this interesting part of the sky.

4.4 Upper Limits on the Hard X-ray

Flux of the Galactic Center

In the study of astrophysical processes near galactic nuclei, the

closest example, our own galactic nucleus, has been hidden at optical

wavelengths due to intervening gas and dust in the galactic disk

(visual extinction -27 magnitudes). The -10 kpc distance to the

nucleus is a factor of 60 smaller than the distance to the nucleus of

M31. The galactic nucleus is, however, observable both in the infrared

and radio bands and also at X-ray energies above about 2 keV.

A review by Oort (Dort 1977) summarizes the principle features of

the galactic center region observed in radio and infrared Within the

central 30 arc min of the galactic center (100 pc at 10 kpc) the mass

density distribution, inferred from the infrared light distribution,

!i shows an increase from 100 MO pc -5 at R-100 pc to 4x10 5 MO pc -3 at R-1

pe, Radio observations show large H II regions and molecular cloud

complexes. The radio source Sgr A West is thought to be coincident

wi;n the dynamical center. Contained in Sgr A West is a compact

non-thermal radio source (Balick and Brown 1974) with VLBI determined

Eμ ^•

d scale sizes at 3.6 cm of 07015 (150 A.U.) (Lo at al. 1981). The 2.2um

49
50

T

infrared sourco IRS16 is commonly accepted as the counterpart to the

non-thermal radio source. This position is also the dynamical center

of a group of compact clouds seen in Ne II infrared fine structure

emission (Lacy 1980). Recently, deep exposures of this region using

charge-coupled devices have discovered a far red (9200A) source close

to the reported location of the non-thermal radio source (Ricker

1982b).

Detection of X-rays (1-10 keV) from the Galactic Center began with

the UHURU satellite which found evidence for an extended source (or

source complex) about the galactic center designated GCX (Kellogg

1971). Many instruments using modulation collimators and coded masks

produced images of the galactic center with angular resolutions from

1 2.5 arc min to 15 arc min (see Skinner for summary of results, Skinner

1979). Six sources were discovered within 1.5 0 of the galactic center.

The brightest, A1742-289, was a transient source with a peak luminosity

of 4x1038 ergs/sec (1-10 keV) in February 1975, though it faded to

2x10 36 ergs/sec by June 1976. The source is 1 arc min from the galactic center and is considered a typical X-ray transient not necessarily

associated with the nucleus (Skinner 1979)•

1

The second High Energy Astronomical Observatory, HEAO-2 or

Einstein Observatory, observed the galactic center region in 1979 using

a focussing X-ray telescope over the energy range .5 to 4.5 keV (Watson

at al. 1981). They found 12 point sources and diffuse emission within

w

0.5 0 of the galactic nucleus. One of the sources was coincident with

Y

the galactic center (1 arc min resolution) and had a luminosity of

^m

1.5x1035 ergs/sec. This luminosity may be in error by as much as a

,

,
51

factor of 10 due to assumptions about the spectral shape and low energy

cutoff that are used in converting measured counts to source luminosity. The variability was less than 2% over 6 months. The Einstein MPC

also detected a flux (2 to 6 keV) from the galactic center region

(Watson et al. 1981) although the field of view contained all of the 12

point sources as well as the diffuse emission. The measured flux is

consistent with UHURU results for OCX and is plotted in Fig. 4.3.

In the hard X-ray regime, source confusion due to large fields of

view has been the main hindrance to observational progress. Balloon

observations with 1.5 0 angular resolution discovered several sources

(20-50 keV) within 10 0 of the galactic center (Ricker et al. 1976)

though emission from the galactic center was not definitely observed

since 0X1+4 was also in the field of view. The OSO-8 hard X-ray

detector with a 5' angular resolution detected a strong source in

September 1978 whose 95% confidence contour contained the galactic

center (Dennis et al. 1980). A photon power law spectrum from 21 to

197 keV is best fit to their data with a power law photon index of

-2,3±,3,

The low energy detector (LED) on the A4 experiment on the HEAO-1

satellite observed the galactic center region 3 times separated by -6

months. Using data fitting procedures on the 1.6 0 angular resolution

scanning exposures, a source located within 0.5° of the galactic center

was detected on all three observations in the energy band 13-100 keV

(Matteson 1982). This source had the hardest spectral index of all

i

sources within 10° of the nucleus, accounting for 14% of the flux in

the 13-50 keV energy band and 23° of the total flux from 50-100 keV for

n 

..
.••.•..•• MPC 9/79

HERO I

1 9/78

2 3/78

- - - - - OSO 8 9/78 1

MIT/Leiden 5/80

951 conf. upper limits

2

L

101 102 103

Energy (keV)

10-5 ^"'

100

I

1

V

I

52

ORIGINAL PAGE IR

OF POOR QUALITY

10-1

GALACTIC CENTER SPECTRA

K^ '

10-2

a^

U

N

10-3

U

0

10-4

Figure 4.3 X-ray spectra of the Galactic Center. The MIT/Leiden 95%

confidence upper limits confirm the variability of this hard

X-ray source.

C
53

the observation in September 1977. The spectral fits for two observations from HEAO-1 are shown in Fig. 4.3 along with the OSO-8 results.

The variability in the higher energies over the 6 month time period are

apparent. The variation at 100 keV is a factor of 3 while only 15% at

20 keV. The luminosity from 10 to 100 keV for the 9/78 HEAO 1 observation is 7x10 37 ergs/sec.

At energies above 100 keV most observations have very large fields

of view containing many X-ray sources. At 300 keV HEAO-1 and HEAO-3

showed variability of a factor of 4 and 8 decrease over 6 months,

respectively (Riegler at al. 1981). Although this cannot be positively

identified with the lower energy Galactic Center source, the large

variability implies a strong single emitter with a peak luminosity of

-3x1038 ergs/sec (Matteson 1982) in the 100 to 1 MeV range or 30 times

the Crab Nebula luminosity.

4.4.1 May 1980 Drift Scan

The MIT/Leiden gondola scanned across the galactic center region

twice on May 8. Because the Galactic Center only reaches an elevation

angle of -30° at transit, the atmospheric column density to the source

was 9.2 gm/cm 2 which gives an atmospheric transmission of 25% in the

60-120 keV band.

The response function of the collimators vs. time, R(t), was

calculated based on the gondola aspect (azimuth, elevation, longitude,

latitude, roll, and pitch); the source (RA, Dec, time); and the calibrated collimator axes with respect to the gondola. The response of

Module A and B to the galactic center is shown in Fig. 4.4 (the collimator response is normalized such that an on-axis source has n respjai
215 225 235 245 255

TIME (minutes CDT)

W

an

z

0

a

U1

ce 0.4

0.0

0.2

0.0

205

&8

1.0

a

fi

54

Ur1u1t

Collimator Response to Galactic Center

Figure 4.4 The 1500 cm' phoswich array collimator response vs. time

during 2 drift scans across the Galactic Center. The response

is normalized such that an on-axis source with a response of 1

would be observed with 592 cm' effective area. The glitch at

minute 222 was due to the elevation drive suddenly moving

0.2°.

Vi
f[p[n`yV^^`Il

55

onse of 1 with the effective area equal to 592 cm2).

The measured count rate vs. time, F (counts/minute), for each

module was fit with two free parameters: a constant background rate

B(cts/min); and a constant source flux A (eta/min) times the independent variable R(t).

FW-B + A • R(t)

A chi-squared minimizing technique was used for 3 energy bands: 40-60

keV; 60-120 keV; 120-200 keV. The source flux was then corrected for

detector and window efficiency and atmospheric transmission resulting

in a source photon flux at the top of the atmosphere (Ballintine 1981).

Though small positive detections occurred in the 2 lower energy bands

for both modules A and B, the data were consistent with a source flux

equal to zero. Upper limits (95o confidence) to a source at the

galactic center are plotted in Fig. 4.3. The 95% confidence limits

were determined using the X 2 minimum + 3.8 criterion as described by

Avni (1976).

This observation indicates a continued variability of the hard

X-ray source near the Galactic Center. The variability is at least a

factor of 2 at 60 to 120 keV over a period of 1.6 years. This strengthens the identification of this 10-100 keV source with the (>100 keV)

a highly variable source. The next logical step of identifying this

bright, variable hard X-ray source with the Galactic Center is on less

firm ground. The best argument is based on extragalactic observations

r

of active galactic nuclei (Seyferts, quasars, N galaxies) where highly

luminous objects vary on short time scales and are located close to the

dynamical center of the galaxy. The identification with the Einstein

re
-r:g++^' V i1 spy .• I i'•

56

source at the galactic center that has a constant luminosity in soft

X-rays cannot be proven, though the NEAO-1 A-4 results vary less with

time at lower energies. Qualitative source models have beer. proposed

(Matteson 1982) in which high temperature (10 10K), high luminosity

(-1038 erg/sec), short timescale (6 months), energy outbursts occur in

a central region providing the hard X-ray flux. Outer regions (which

are heated to only 107 K, have luminosities -10 35 ergs/sec, and lifetimes of years) provide the constant soft X-ray luminosity which is the

integrated effect of many impulsive events.

The above results indicate that what appears to be a stable low

intensity source in one energy band might be a highly variable and

luminous object in another energy band. Sensitive hard X-ray telescopes with good angular resolution are needed to resolve the candidate ,

Galactic Center sources and monitor their variability. Such an

instrument is described in Section 5.

'T

I^ I

i

{

4

`^ 90

^u
n

57

5.0 DESIGN CONSIDERATIONS FOR A NEW HARD X-RAY TELESCOPE

After the demise of the 1500 cm balloon-borne hard X-ray telescope due to a free-fall from float altitude, the design of a completely new telescope became necessary. Though starting from scratch

was not desired, it did allow a rare freedom to rethink and redesign

all ma6or systems in a future gondola, especially the main X-ray

detectors. The constraints imposed upon a design for a balloon-borne

detector are considerable. The goal was to match the sensitivity of

the 1500 cm detector in hard X-rays with a gondola that would weigh

less than 400 kgs. This would allow possible long duration balloon

flights (-30 da-s) at higher altitudes.

5.1 Sensitivity

Signals from cosmic X-ray sources are received as photons and

counted individually by the detector. The limitations of measuring

observable properties such as intensity, location and spectra are

statistical in nature and X-ray telescope parameters are optimized to

minimize statistically determined errors. The basic parameters of an

1 X-ray telescope are: collecting area, A, in cm 2 ; differential background counting rate, B (counts/cm 2 sec keV); energy bandwidth, AE, in

keV; observation time, t; and detection efficiency c (including atmosa.

pheric transmission). The number of signal counts, N s , from a source

with differential photon flux, F (photons/cm 2 sec keV), observed with a

counter of effective area, A, efficiency e, over a time t, bandwidth AE

is:

N FAtAEc

s

I During this same time the number of background counts, N B , observed is:

1'

V

MV
n'

58

.i

I

NBmBAteE

(B is measured at the detector so e'1)

The total number of counts observed is N s +N B so a separate measurement

of the background of length t must be made and then subtracted to yield

the signal counts:

Ns=(Ns+NB)-NB

i

and the variation:

ANs°Ns+NB)2+AN,)112 (d( 

assuming F and B are uncorrelated and .`iat changes, AB, in the background rate are small compared to statistical fluctuations. ( Changes

in the background counting rate with time is a major problem and many

1

techniques have been developed to reduce this systematic error. See

Section 4.1.)

l For a faint source (i.e., eF<B) the statistical fluctuations in B

r•

Fdominate the error. For NB>>Ns:

pNsa AN 2NB = 2BAtdE

Y^ (This is gust the standard deviation of a Poisson distributed count

rate.) The minimum detectable flux, Fmin, can be defined as that flux

where the signal counts are a certain number, S, of standard deviations

? above the background

Ns Fm. 

AteAE

Sa __

ANs2BAtaE

S /TB—

Fmin ^ e ttGE

The number used for S depends on the number of separate measurements of

D

W
`I

1

59

V

length t and the significance required. A "four sigma detection",

(S-4), has a probability of randomly occuring once in 30,000 measurements (Bevington 1969).

To make Fmin as small as possible (i.e. sensitive to very faint

sources), the goal in designing a hard X-ray detector is a large

effective area of efficient, low background detectors that can observe

for a long time. Because these goals are sometimes incompatible,

tradeoffs among the design parameters must be made. For example, to

take advantage of a satellite's long observing time, effective area

(and its corresponding mass) is usually sacrificed. Large areas are

also difficult to shield so higher background rates are reluctantly

accepted.

The formula for Fmin does disguise the fact that A is the photon

collecting area, not the detector area, while B refers to the background counts observed at the detector divided by the collecting area.

Thus, if the collecting area could be increased with X-ray mirrors or

concentrators while keeping the detectors the same size, not only would

A increase but B would decrease and Fmin would be proportional to A

rather than A-1/2. (This is not true if the background, B, is dominated by diffuse sky flux. In that case an increase in collecting area

would.increase the background and Fmin would be proportional to A-1/2).

This technique of separating the collecting area from the detector

has been used in optical astronomy since Galileo (Galileo 1610) and

recently has been applied to all of the electromagnetic spectrum up to

and including the soft X-ray ( 1 -5 keV) imaging telescope of the

Einstein Observatory (HEAD-2) (Giacconi 1979).

4 C
r

60

.r

U

5.2 Bragg Diffraction Concentrator

7

In the energy regime accessible with balloon-borne detectors

(greater than 20 keV), grazing incidence reflection of X-rays does not

work (reflection efficiencies <<1%). But X-rays of this energy can be

Bragg diffracted with reasonable efficiencies using thin crystals of

lithium fluoride (Lytle 1969). The angle of diffraction has a finite

spread ("rocking curve" -.2 0 for LiF) so any focal point would be

"blurred" and concentration of X-rays to a single detector rather than

imaging would be the result. However, the inability to image X-ray

sources using a Bragg diffraction concentrator relaxes some construction and pointing accuracy constraints. A concentrator need not be

perfectly rigid with arc-second tolerances and can be made relatively

lightweight and inexpensive.

5.3 Low Background Balloon-Borne Telescope Design

The preceding design considerations have lead to the concept of a

hard X-ray diffraction telescope consisting of a lithium fluoride lined

paraboloid that concentrates hard X-rays onto a small, well-shielded

HgI 2 detector (Figs. 5.1 and 5.2). This design has been accepted as

the basis for the Hard X-ray Diffraction Telescope (HXDT) to be flown

on a future Space Lab mission. The size and shape of the paraboloid

determines the energy bandwidth of the photons concentrated to the

detector at the focus. '.`n-axis X-rays entering the aperture will Bragg

diffract from the LiF crystals and proceed to the focus only if their

energies are correct for that position on the paraboloid. The Bragg

l

diffraction serves only to concentrate on- or near-axis X-rays; no fine

iI_ energy resolution is derived from the concentrator, just the solid

Li _,
LMila

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ORIGINAL PAGE `g'

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state detector.

To limit the size of the concentrator (and therefore a possible

balloon gondola), two nested paraboloids optimized over 1 octave in

energy (40-80 keV) has been the basis of a new gondola design. This

energy range was chosen because it lies above the energy of significant

atmospheric absorption at balloon altitudes. The upper energy boundary

is limited by the size of the paraboloid. For a small, lightweight

balloon gondola, 80 keV seemed a reasonable compromise. Some interesting astrophysical processes that have signatures in this energy

range include cyclotron emission or absorption features from magnetic

neutron stars and spectral turnovers in active galactic nuclei power

law spectra. Gamma-ray line emission from long lived radioactive

t

isotopes such as 44T1 is expected to be observed in supernova remnants

i

in this energy range. Observing these narrow line features will

require excellent energy resolution.

The detector system behind these concentrators consists of thin

HgI 2 detectors surrounded by a large bismuth germanate scintillator

operated in anticoincidence (Fig. 5.2). A prototype detector system

has been designed, constructed and integrated into a small balloon

gondola (Fig. 5.3) to measure the actual counting rate background at

rate, B(cts/see cm keV)

as well as a check on

This will allow further

rejection and overall

balloon altitudes. This background countin;

will be used in the sensitivity calculation

background models of the detector systems.

modification and improvements in background

gondola design.

de/
6

JAI

ORIGINAL PAGE IS

OF POOR QUALITY

I
4 65

I

r,

r+'Y .

6.0 SHIELDED MERCURIC IODIDE BACKGROUND MEASUREMENT EXPFRInNT

In research work leading to this thesis, a small gondola (Fig.

5.3) containing a prototype detector for a hard X-ray diffracting

tvtoscope was designed, constructed and flown at balloon altitudes.

Its purpose is to measure the background counting rate at the eventual

telescope operating environment (>40 km altitude, balloon or satellite). Zie radiation environment (X- and gamma rays, high energy

charged particles, neutrons) is quite different at ground level due to

atmospheric shielding and ground radiation. The measurements are

compared with predictions made by a detector Monte Carlo simulation.

Once the simulation is checked, design parameters ( geometry of detectors, shielding components and discriminator settings) can be varied in

the simulation to optimize (lower) the expected background counting

rate.

The small gondola can be conceptually divided into two subsystems:

the prototype detector consisting of a bismuth germanate shielded

mercuric iodide detector (Fig. 5.2) described in Section 6.1; and the

detector support systems consisting of electronics, power systems, data

recording, and temperature control all described in Section 6.2. The

radiation environment and detector computer model is described in

Section 7.

x 6.1 Prototype Detector System

The overall design goal for the detector behind the X-ray concentrator was to minimize the number of background events in the energy

range of 40-80 keV while keeping the response to the concentrated

t

w ^e

signal X-rays as high as possible. The proposed concentrator requires
1

66

a detector with an area equal to the size of its focal "spot" (-1 cm 

for a lithium fluoride concentrator optimized for 40-80 keV). The

detector must also have a field of view larger than the solid angle

subtended by the concentrator to "see" all diffracted signal X-rays

( p-.12 steradian for same concentrator). The active anticoincidence

shield should surround the detector except the aperture and if not stop

incoming radiation then at least detect and veto it. Passive mass

viewed by the detector (i.e., structural material such as aluminum that

cannot be used to detect radiation) should be minimized because secondary radiation generated in this mass might not be vetoed.

The prototype detector system consists of two HgI 2 detectors

mounted back to back and operated in anti coincidence. The two detectors are placed inside a bismuth germanate (Bi 4Ge 30 12 or "BGO") scintillating shield along with two hybrid charge sensitive preamps. The

BGO shield is also operated in anticoincidance with the main (top)

detector. The axially symmetric "well" containing the detectors is gas

tight (.5mm Al entrance window) to retain one atmospheric pressure at

balloon altitudes. Such a gas tight enclosure eliminates the problem

of high voltage corona breakdown encountered at low atmospheric pressures and avoids the need for potting the detector assembly.

6.1.1 Primary Mercuric Iodide Detector

The "main" HgI 2 detector that will eventually be used to detect

X-rays from a celestial X-ray source is the top detector in the back to

back geometry (Fig. 6.1). The actual HgI 2 detector used in the test

flights has an effective area of 0.8 cm (defined by the size of the

palladium electrode) and is 390 microns thick. The bottom detector has
Top HgI2

Detector

RexoIite

Substrate

Paper Based

PC Board

^-- 1 inch —^

llnl

I

V/

ORIGINAL PAGE 13

OF POOR QUALITY

Hg I z DETECTOR ASSEMBLY (DETAIL)

1

Bottom HgI2

Detector

Stainless Steel

Support Structure

Figure 6.1 Detailed view of mercuric iodide detector assembly showing the

components that fit in the axially symetric well inside the

bismuth germanate shield.
Bismuth germanate was chosen as the shield material because of its

i

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1 68

the same area and is -1000 microns thick. The bottom detector's main

purpose is to provide active shielding from radiation generated in the

passive material underneath (printed circuit boards and preamps) as

well as better protection from aperture charged particle flux. The

HgI 2 detectors are separated by a thin (13um) mylar insulator and

mounted on a Rexolite substrate. The HgI 2 crystals are grown and then

fabricated into the "back to back" units by EG&G, Inc., Santa Barbara.

The charge sensitive preamps are mounted on paper based printed

circuit boards and are held in place underneath the bottom detector by

a low mass stainless steel structure. A glass epoxy based printed

circuit board was avoided because of the high concentration of potassium 40 in natural glass (Knoll 1979, pg. 775). The preamps were based

on a low power design for SAS-C (Goeke 1975) and were modified and

hybridized by EG&G, Las Vegas. They were placed close to the detectors

to minimize input capacitance and noise pickup. These preamps were

chosen because of availability but future designs might consider

placing just the input field effect transistor (FET) in the well to

reduce passive mass. Small copper wires bring preamp power, high

voltage, and ground down the well to the detector assembly and the top

and bottom signals out of the well. The well is lined with electrically grounded aluminum (.25 mm) to provide electromagnetic shielding.

Outside the well and field of view of the detectors but inside the

pressure container are RC filters for the high voltage and preamp

power.

6.1.2 Bismuth Germanate Shield
.'.rl

69 T

very high photoelectric cross section and density. In terms of its

stopping power it can almost be considered "active lead". The crystal

is 7.5 cm in diameter and 12.5 am long with a machined 7.4 cm deep x

2.5 cm diameter axially symmetric well in which the HgI 2 detectors are

placed. The outside surface of the crystal is coated with reflective

paint and encased in aluminum to protect the crystal and prevent light

leaks. A Hammamatsu R1307 photomultiplier is mounted on the unpainted

bottom of the BGO crystal opposite the well (Fig. 5.2) and surrounded

by a u metal shield to protect it from magnetic field induced gain

changes.

At the time when the BGO was ordered from Harshaw Chemical Co. the

7.5 cm diameter was the largest size available. The 2.5 cm inner

d','',ameter of the well was determined by the preamp size and ease of

^,n,ztruction. In the energy range where BGO has the smallest attenual ion coefficient (-1 to 3 MeV), non-aperture gamma rays that intersect

the HgI 2 must traverse at least 2.5 cm of BGO (-1 attenuation length).

At lower and higher energies the attenuation is,much greater. The top

HgI 2 detector is placed 2.6 cm deep into the well measured from the top

of the BGO. The BGO then acts as the collimator with a flat topped

response and an effective circular field of view of -.12 steradians.

With such a large field of view the diffuse atmospheric and cosmic

X-ray background would dominate the background count rate. For the

diffracting concentrator concept to work, X-rays that can reach the

HgI 2 detector directly through the aperture must be stopped so that

only diffracted X-rays from a celestial X-ray source are detected. The

Io "direct" (undiffracted) field of view must be blocked using a graded

ifr•
76

t

shield of lead, tungsten, and tin. The lower atomic number materials

are included to stop the characteristic K X-rays of the higher. For

the test flight of the prototype detector this passive shield was

placed directly over the entrance window though in actual use with a

concentrator it would be further from the detector so as not to block

the diffracted X-rays. This is shown schematically in Fig. 6.2. The

transmission of 1 mm of Pb + 1 mm of `d + 2.5 mm of Sn is shown in Fig.

6.3.

6.2 Balloon Gondola and Associated Electronics

The detector prototype (BGO shield and HgI 2 detectors) is contained in a small (30 cm x 30 cm x 45 cm) aluminum framed package along

with its support electronics (Fig. 5.3)• Power for the electronics and

heaters is supplied by lithium batteries while the high voltage for the

HgI 2 crystals comes from a charged 0.2 microfarad capacitor. The

capacitor stays within 30% of its original voltage for -30 hours

because of the small leakage currents of the detectors, cables and

capacitors themselves. This detector biasing method also removes noise

contributions from low to high voltage converters. The detectors are

kept biased continuously before and during launch of the balloon to

avoid "cold start" problems (see Section 3.1). Low voltage power is

turned on at approximately 6 km altitude by parallel pressure switches

operating a relay. This conserves the batteries and protects against

electrical shorting at launch and landing.

Once power is turned on, data collection begins. The main part of

the electronics takes random, asynchronous voltage pulses from the

detector preamp, converts these pulses into synchronous, 8 bit digital

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Transmission of Graded Shield

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Figure 6.3 Transmission of photons vs. energy through a graded shield of

1mm of lead, 1 mm of tungsten and 2.5 mm of tin. The lower Z

metals act to stop the fluoresced K X-rays of the higher Z

metals.

80

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words and routes these words to an appropriate magnetic tape track.

The remainder of the electronics monitors th= condition of the gondola

(temperatures, voltages, currents, etc.) and records this housekeeping

data on a separate magnetic tape track. The gondola is self contained

during flight (i.e., no telemetry, no commands, no interface). However, it is essential that the cassette tape be recovered!

6.2.1 Event Logic Electronics

Pulse height spectra are recorded for both top and bottom HgI2

detectors. Output pulses from the BGO photomultiplier are fed into an

Amptek A-101 logarithmic amplifier/discriminator which produces a

digital pulse if the charge pulse amplitude is above a set threshold.

Only the BGO shield count rate is recorded; no pulse height information

is retained.

Small voltage pulses from both HgI 2 detector preamps are brought

to pulse shaping voltage amplifiers (time constant -15us). The last

stage of these amplifiers is a "peak stretcher" which retains the

maximum voltage attained by the pulse and holds it at the input of an

analog to digital converter until reset. If the pulse height triggers

the lower level discriminator, the pulse height will be converted 22

useconds later to an 8 bit digital word. For the top HgI 2 detector,

1

the event logic determines if an event occurred in the BGO crystal or

the bottom detector in the time interval from 13μS before the top event

to 9 uS after. If no "veto" event occurred then the 8 bit digital word

is shifted to the 'valid" buffer for tape track #1. If a "veto" event

occurred in either the BGO or the bottom crystal then the 8 bit word is

shifted into the 'veto" buffer for tape track #2. All events occurring

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in the bottom crystal are converted and shifted into the "bottom"

buffer for tape track #3• Once the shift into the buffer takes place,

the peak stretcher is reset and ready for a new pulse. Total time from

event to reset is approximately 50 microseconds. The event logic

diagram is shown in Fig. 6.4.

Because of the very low count rates in the HgI 2 detectors at float

(-1 ct/sec) the dead times due to the finite conversion time are

negligible. The BGO count rate of -1000 cts/sec plus the veto flag

duration of 22 microseconds gives a 2% probability that a valid event

will be vetoed by accident.

6.2.2 Data Recording

Events arrive at the 8 bit wide x 16 :'t deep buffer asynchronously but are read out synchronously at 8 bytes per second (slow mode)

or at 1024 bytes/sac (fast mode). The fast mode is used on the ground

for calibration and testing purposes while the slow mode allows recording on tape for flight or long term background runs.

Separately for all three data tracks (valid, ve`o and bottom) the

8 bit pulse height is read out of the buffer and inserted into a serial

bit stream 515 bytes (=1 block) long. The bits are Manchester phase

encoded (Rooks 1979) and sent to the tape head amplifiers. The tape

recorder is a 4 track LTD 800, a long term data logger constructed by

Braemar Inc., and uses standard Phillips cassettes. The tape speed is

.08 inches/second and a 560 foot cassette tape (size 120) will run for

23.4 hours. The maximum bit density allowed is 800 bits/inch which

determines the 8 byte /second recording rate.

The tape is formatted (Fig. 6.5) into blocks of 515 bytes separL

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with a preamble (postamble) of 10101010. The next byte is a block

number that is incremented every block and is the same for all 4

tracks. For tracks 1, 2 and 3 the remainder of the 511 bytes are pulse

heights (mostly zero). Position on the tape corresponds to the event

arrival time. Events that arrive during the inter-record gap and post

and preamble are not lost but kept in the buffer until readout after

the block number.

Track 4 has the same tape format as tracks 1-3 but instead of

pulse heights the digital words correspond to various housekeeping

i data. Temperatures, currents, and voltages are all converted to a

voltage between 0 and 10 volts using linear transducers, amplifiers or

simple resistive dividers. These voltages are multip1x . 3d to another

analog to digital converter and the 8 bit digital word is again. inserted into a serial bit stream, encoded, and recorded. Also recorded is

the 16 bit clock/counter and the 16 bit BGO event counter that is reset

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every 2.91 seconds.

6.2.3 Ground Sup^^nrt end Data Analysis

Pulse height spectra from the top or bottom HgI 2 detectors are

accumulated by a Hewlett Packard HP85 desktop calculator v' q an interface that can read a track from the cassette tape or directly from the

gondola. A CS400A data reader built by Braemer Inc. reads and decodes

1 tape track at a time at 10 inches per second and sends this serial

bit stream to the interface which directs this data onto an IEEE 488

interface bus connected to the HP85• The HP85 can analyze up to 30

blocks at a time and will bin all events by their 8 bit pulse height

1
79

(256 bins). An arbitrary number of these spectra may be summed together. A 23 hour tape requires 20 minutes to reduce into a spectrum.

For calibration and testing, the interface can take data directly

from the gondola using the serial bit stream just before it is encoded.

This method can use either the slow or fast mode (8 bytes/sec or 1024

bytes/sec; manual switch selectable). The fast mode is used when the

^ount rate exceeds 8 cts/sec as it does using ordinary radioactive

sources. Overnight runs to measure laboratory background use the slow

mode with tape as does the actual balloon flight.

C -

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7.0 BACKGROUND SOURCES AND MONTE CARLO ANALYSIS

The background counting rate of an X-ray detector at high balloon

altitudes and in orbit ie due to high energy radiation at the detector.

The components of this radiation include: the prima ry cosmic rays and

their secondar'.esincludin ( g neutrons ) produced in nearby matter;

atmospheria gamma. rays from bremsstrahlung and decay of the secondaries; the diffuse cosmic R- and gamma ray background; aad induced

r radioactivity of the detector and its surroundings. The physics of

interaction of each type of radiation are quite different. Shielding

systems that work on one type of radiation sometimes fail to notice (or

even worse) enhance another type. Regardle3s of the type of radiation,

' if 40-80 koV of energy is deposited in the top HgI 2 detector it looks

like a signal R-ray unless vetoed by a simultaneous event in a shield.

The radiation environment at balloon altitudes (40 km or 3 gm/cm2

atmospheric column density) are summarized in Table 7.1. These fluxes

1

are not always what the main detector sees because the transmission of

the shield varies with energy and type of radiation. Each must be

(

(

treated separately using measured ambient fluxes to determine the

individual contribution to the X-ray background counting rate and then

summed together. Secondary radiations created in the detector are

treated along with the original radiation (i.e., electron bremsstrahlung is treated under the charged particles contribution while

{tt

9.

positron annihilation is treated under the gamma rays' contribution).

tg

Y A note on terminology: the detector is considered to be oriented

4. with the aperture pointing at the zenith so the primary HgI 2 detector

is referred to as the "top" detector while the HgI 2 detector below it

79

D

T

I

rte- 6
ee

Table 7.1

Integrated Radiation Flux at 3 gm/cm 2 and

Magnetic Latitude - 42 0 N (Palestine, Texas)

i

{ I

Reference

Ling 19'74

Radiation Type Flux

Gamma rays

photon energy >70 keV 7.6 photons/cm 2 sec

>200 keV 4.8 photons/cm 2 sec

Protons (E>3 GeV) •3 protons/cm 2 sec Meyer 1969

Electrons (E>1 MeV) 1.0 electron/cm' sec Verma 1969

(possibly factor 2-5 to high) Daniels 1974

Neutrons (E>.1 eV) 1.2 neutrons/cm 2 sec Armstrong

i 1973

i

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81 t

T

is the "bottom" detector. Events in the top detector are 'vetoed" if

there is a coincident event in the BGO or bottom detector or they are

"valid" if there is no coincident shield event.

7.1 Gamma Ray Induced Background

Gamma ray induced background in some ways is the most straightforward to understand yet the most difficult to calculate. The physics

of interaction with matter are well understood and the cross sections

can be expressed as analytical functions (McMaster, 1969). Yet the

i

complex geometry of the BGO shield combined with the high probability

of large angle scatters that are random in nature required the use of a

' Monte Carlo simulation. Before describing the details of the analysis e .

some general conclusions c= be stated. The gamma ray spectrum at

j float altitude has been modelled by Ling (Ling 1974) and is shown in i

•

Fig. 7.1 for 3 gm/cm 2altitude. The transmission of gamma rays through

2.5 cm of BGO is plotted on the same graph. There will be a certain

energy range where the transmission becomes non-negligible and the flux

of photons is still appreciable. This "shield leakage" must then be

folded with the detector response. A typical valid event would be a 2

MeV gamma ray penetrating through the BGO, Compton scattering in the

HgI 2 detector leaving a 50 keV Compton electron and then escaping

i

without being d^_.ected again, all imp-obable events. From Fig. 7.1 the

need for an aperture shield is evident to stop the direct flux of low

' energy (20-200 keV) X-rays.

7.1.1 Monte Carlo Simulation

The Monte Carlo simulation of the gamma ray interactions with the

detector assembly (BGO, HgI2 detectors, passive aperture shield) used

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82

ORIGINAL PAGE 19

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Atmospheric gamma ray

ambient spectrum at

3 gm/cm2

102

103 104

Energy (keV)

Figure 7.1 Plot of ambient gamma-ray spectrum at 3 gm/cm 2 taken from Ling

(1974). Also plotted is the photon transmission of 1 inch of

bismuth germanate to show the effectiveness vs. energy of the

shield crystal where the minimum thickness is 2.5 cm. The

shield leakage will be a maximum between the energies of -•5

MeV to 3 MeV, where the atmospheric flux is still apnrecialble

V

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101

10-5

J 10-1

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61

as an input an assumed atmospheric photon flux density (function of

energy and direction and the geometry of the detectors). The output

consists of the resulting energy spectra of valid, and veto events for

the top detector, as well as the spectra in the bottom and BGO detect-

'p ors. Details of the technique and computer code can be found elsewhere

9 (D. Wood, MIT Senior thesis, 1982).

The simulation employed a ray tracing achnique and can be summarized as follows:

(1) A photon is generated from the input distribution with an

energy, direction, and starting point.

(2) With the attenuation coefficient of the material in which the

! photon is travelling (HgI 2 or BGO), the position of next

interaction is calculated based on the geometry of ti:e detector and direction of travel.

r (3) The type of interaction is determined based on the different

calculated cross sections (photoelectric, Compton scatter,

pair production). The energies and direction of any secondary

photons are generated if needed for the interaction and the

energy deposited is calculated.

(4) Secondary photons are traced from the interaction site starting with step (2) until the photon escapes entirely from the

detector or until its energy drops below a set discriminator

level.

(5) When the incident and all secondary photons are completely

^. traced the energy deposited in each detector is separately

P

accumulated and the appropriate energy bin in the bottom and

Maw

i
s

84 T

B00 spectra are incremented. A deposit in the top detector

will increment the valid spectrum unless a concurrent deposit

occurs above a discriminator level in the other two detectors.

In that case the veto spectrum would be incremented.

(6) All energies are cleared and the number of incident photons is

incremented and another photon is generated in Step (1).

Three types of photon interaction were used in the simuls tion:

photoelectric absorption, Compton scattering and electron-positron pair

production. The attenuation coefficients used for the different

detector materials are shown in Fig. 3. 1 and 3.3.

For the photoelectric interaction the simulation would generate a

K-alpha X-ray photon if the energy absorbed was greater than the

bind_nig energy of the K shell in the material absorbing the initial

photon. The K. alpha photons were assigned the K-alpha-one energies of

77.1 keV and 70.8 keV for bismuth and mercury, respectively. The K

esca=pe photons of iodine, germanium and oxygen were ignored because of

their low escape probabilities. The strength of these line features

are very important for determining the contribution of gamma rays to

the detector background. K fluorescence is a "signature" of photon

interaction and should scale as the gamma ray continuum. All other

radiations produce a continuum background in the 20-200 keV range which

cannot be separated from that produced by gamma rays. The line features supply a measurable test for the simulation's prediction.

The Compton scatter cross sections and angular distributions are

based on the Klein-Nishina formula (Knoll 1979, pg. 69)• In the energy

range where this cross section dominates (.5-2MO), the Monte Carlo
L

85

simulation is most needed to follow the mulitple scatterings. Fig. 7.2

is a two dimentional visual representation of the Monte Carlo pr-,gram

showing the various scatters and interactions.

1

Pair production, although relatively infrequent because it requires a photon with an energy greater than 1022 keV, generates the

positron annihilation line at 511 keV which could also be a measurable

prediction of the simulation. However, the dynamic range of the

electronics (10-250 keV) did not allow this measurement.

The gamma ray spectral distribution that was input to the simulation was adapted from Ling (Ling 1974) and can be expressed as

dN. 16E-1.4 photons/cm 2 sec keV

at 3.5 gm/cm 2 along with a line feature at 511 keV of strength .3

photons/cm 2 sec. The directional distribution was assumed to be

isotropic in the upper and lower half spaces. The ratio of upward to

downward flux at different energies was estimated from Ling's angular

distribution and was fitted by the function:

upward_flux 1.7 In (E/87 keV + 1.84) downward }lux

During an actual flight the balloon cannot be expected to arrive

exactly at a desired altitude used in the simulations. Fortunately the

variation with depth of the photon flux is slight for lower energies.

Descending from 3 gin/cm 2 to 5 gm/cm2 there is a 10% increase in flux at

300 keV while effectively no variation occurs at higher energies (>1

MaV) (Ling 1974).

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86

2-Dimensional Ray Trace

photoelectric ---0 pair production

absorption

Compton scatter •.... escape from crystal

Figure 7.2 Two dimensional plot of results from a Monte Carlo simulation

of gamma-ray interactions in the shielded HgI, detector

assembly. The simulation followed 50 photons generated on the

outside of the bismuth germanate shield until they were

absorbed or escaped from the shield. The photons were

distributed in energy according to the ambient gamma-ray

spectrum at 3 gm/ cm' (Fig. 7.1). Thereforc, most of the

`

interactions occur near the outside surface of the BGO shield.

I

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87

ORIGINAL PAGr E3

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m

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Energy (keV)

Figure 7.3 Predicted valid background spectrum due to gamma-rays for the

top mercuric iodide detector at 3 gm/cm' float altitude. The

small excess between 100 and 300 keV is due to the leakage

through the passive aperture shield. The atmospheric positron

annihilation line at 511 keV is also apparent.
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Gamma Ray Induced Backaround

-- -

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a ^

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^' I

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Figure 7.4 Predicted vetoed (top detector) and bottom detector spectra at

3 gm/cm' due to gamma - rays. The spectra are featureless except

for the Hg and Hi K X-rays and the positron annihilation line

along with its Hg K escape feature, 70 keV lower in energy.
L^•^.4 (D'

t

89

7.1.2 Results of Simulation

The results of the Monte Carlo simulation are shown as energy

spectra for the top (valid and veto) and bottom detectors in Figs. 7.3

and 7.4 and are taken from the thesis of D. Wood (Wood 1982). The

valid spectrum attained a broad general maximum at an energy of -200

keV and a local minimum between the energies of 25 to 140 keV. A

strong line feature exists at 511 keV due to the atmospheric line

feature. Aside from the 511 line,the valid spectrum exhibited no other

resolvable line features. However, the processes which produced line

features in the bottom and veto spectra were not expected to contribute

to the valid spectrum because the detection of K-a photon and positron

annihilation photons required concurrent interactions in either of tho

anticoincidence detectors that produce such photons. These events are

always vetoed if the energy deposited in the veto detectors exceeds

the discrimination level. The simulations verified this condition.

Predicted integrated count rates of the continuum background and

line features due to the gamma ray flux are tabulated in Tabld 7.2.

Only the K fluorescence line features can be compared directly to the

flight results (Section 8) because the continuum background must be

summed with background produced by neutrons and charged particles.

7.2 Neutron Induced Background

A significant contributor of background events in the HgI 2 detector are neutrons generated by primary cosmic ray interactions with

oxygen and nitrogen nuclei in the atmosphere. The event rate in the

HgI 2 detectors is significant because the shielding schemes for suppressing gamma rays and charged particle3 are not effective at atten1.

r..v.c
0,

Line Features t+

Bismuth K Xray

Mercury K Xray

2.6 x 10 - cts/sec cm

1.0 x 10-2

90

ORIGNN1',L PAUL ' fs

OF POOR QUALITY

Table 7.2

Predicted Background Count Rates in BGO-Shielded

Top Hg2 2 Detector Due to Atmospheric Gamma Rays

at 3 gm/cm and Magnetic Latitude 42°N

t.

Continuum:

Energy(keV)

40 - 80

80 -120

120-160

160-200

200-240

240-280

Valid t Backround Rate (10 -5 cts/sec cm keV)

1.3 + .4

.8+.3

.3 + .1

1.3 + .2

1.2 + .3

1.4 + .4

t A "valid" event is an event not accompanied by

a coincident event in either veto shield der,ector.

tt Present only in vetoed spectrum

'i

^e

^j

uJ

L
R'".^ N P aTL -

- - -

CDI'

91

t

0

uating the neutron flux. The mean free path for diffusion and absorption of fast neutrons (En >1 MeV) in B00 is approximately 6 cm which is

comparable to the thickest part of the shield. The only advantage of

this situation is computational in that the effect of the neutron

background can be calculated analytically assuming a slab detector in

an isotropic flux of neutrons. The complexity enters in the types of

neutron interactions and their corresponding cross sections: elastic

scattering, inelastic scattering, and neutron capture with a prompt

gamma ray and delayed beta decay. Errors in these calculations will

result more from extrapolated neutron cross sections and fluxes rather

than the simplified geometry.

The upper atmosphere neutron flux has been measured and modelled

by a number of authors (Armstrong 1973; Marker 1973; Klumpar 1973) at

various atmospheric depths and magnetic latitudes. From thermal

energies to 20 MeV the flux at 3 . 5 gm/cm 2 , magnetic latitude a-42°N

(Palestine, Texas) can be expressed as F(En)-.1En-. 83 neutrons/cm 2 sec

MeV. The spectrum remains flat from 20 to 50 MeV and than falls off as

E-2 above 50 MeV. In the energy range 10 to 100 MeV, 80% of the flux

at 3.5 gm/cm 2 is upward moving neutrons. To calculate the event rates

in the IigI 2 crystal for the different types of neutron interactions,

the physics of the interaction as well as the tabulated cross sections

(Hughes 1964) determined the energy range of interest of the incident

neutrons. For example, only neutrons with energies greater than 1.3

MeV could elastically scatter off iodine and leave 40 keV in the

detector due to simple conservation of momemtum.

^- The maximum energy transferred in an elastic scatter of a neutron
92 I

and an atomic nucleus is proportional to 4A/(A+1)' where A is the

atomic weight (Knoll 1979, pg. 571). This fact gives HgI 2 an advantage

over most other X-ray detectors due to the high atomic weight of Hg and

I (200.5 and 127, respectively). The elastic event rate at a given

energy, E, will be due to neutrons with energies, En , greater than

E(A+1) 2 /4A. Using tabulated elastic scatter cross sections (extrapolated as a constant to higher energies) with the incident neutron f ax

the event rate was separately determined at 40 and 80 keV in cts/sec

cm keV for Hg and I 2 and a power law was fit to these points. The

results are shown in Table 7.3 and in Fig. 7.5.

Inelastic scattering from a nucleus by a neutron leaves the

nucleus in an excited state which will decay by emitting a gamma ray.

This is an important process because even though the cross sections are

saall, the detector sees the events in a small energy bandwidth. But

with a high resolution detector, these line features can be "resolved

out" (i.e., the signal in this small bandwidth is ignored). The line

features of interest are the 59 and 203 keV gamma rays of '27 I and the

160 keV line of 199Hg. Higher energy gammas have a higher probability

of escaping the crystal and a lower production probability due to small

cross sections and fewer neutrons above the required threshold energy.

Starting with the neutron flux above the threshold energy (-inelastic

gamma energy) and tabulated cross sections, the number of gammas

produced is calculated. This is multiplied by the probability that the

gamma, will escape the detector without interacting. The line strengths

(cts/cm2 sec) are given in Table 7.3 and plotted in Fig. 7.5 folded

with a Gaussian onergy response with 5 keV FWHM. The 59 keV line
I

I

\'1F I `J

93

ORIGINAL PAGE 18

10-4

OF POOR QUALITY

Calculated Neutron Induced Background

Elastic Scatter

59 keV

10-5

U 2M ker'

Q)

U

N

E

1281 Beta Decay

168 keV

s= 10-6

0

U

Inelaetic Scatter

Gam Rove

110 101

102

103

a

re

Energy (keV)

Figure 7.5 Predicted neutron induced valid background showing the three

types of interactions. The elastic scatter power law results

from the power law energy spectrum of the fast neutron flux.

The gamma-rays are the result of nuclei being excited to

higher energy states by inelastic collisions with fast

neutrons. The 11 ". beta decay induced energy spectrum results

from the capture of 20-100 eV neutrons by "'I and then beta

decavina (half life - 25 minutes).
i

94

Table 7.3

Predicted Background Count Rates in BGO-Shielded

Top HgI 2 Det2ctor Due to Atmospheric Neutrons at

3 gm/cm and Magnetic Latitude 42°N

Continuum:

Valid Background Rate (10-

5

 cts/sec cm keV)

Energy (keV) Elastic Scatter 128I beta decay Total

40 - 80 1.00 .20 1.20

( 80 -120 .64 .20 .84

} 120-160 .46 .21 .67

160-200 .36 .21 .57

200-240 .30 .21 .51

240-280 .26 .20 .46

Line Features (inelastic scatter):

59 keV 1.40 x 10_4 cts /sec cm 

1 160 keV .07 x 10

203 keV .25 x 10-4

I

^I
95

rt

u

feature has been observed in all balloon-borne NaI and CsI detectors (-

9 Matteson 1977; Peterson 1975) and is unfortunately right in the middle

+

of the energy range of ircerest (40-80 keV).

Neutrons can be captured, (n, Y) reaction, by the 127I nuclei or

the seven stable isotopes of Hg. The prompt effects of the capture are

the emission of many gamma rays up to the bin,ing energy of the neutron

-8 MeV. The small probability that these photons will be absorbed in

the thin HgI 2 crystal combined with the high probability that the

associated photons will be absorbed in the veto crystal leads to an

insignificant valid contribution from these prompt events.

The decay of the daughter products of this same neutron capture

may produce delayed gamma, R and beta rays. 128 I has a half life of 25

minutes and quickly comes to equilibrium at balloon altitudes emitting

a beta ray with maximum energy of 2.2 MeV. The beta rays that leave

energy in the 40-80 keV range are the initial low energy betas and the

high energy betas that are created near the surface and deposit 40-80

keV before escaping the HgI 2 crystal. The deposited energy spectrum is

calculated assuming the beta emitters are distributed evenly throughout

the detector using a Monte Carlo technique. The resulting spectrum is

plotted in Fig. 7.5 and the contribution to the valid spectrum is given

in Table 7.3.

Neutron capture in the Hg isotopes usually leads to another stable

Hg isotope with prompt Y-rays. The unstable daughter products are

197 2033Hg' 205Hg with half lives of 65 hours, 46 days, and 5 minM; utes, respectively. Only 205 Hg, which emits a beta particle, has a

8o

short enough half-life to be noticed in a balloon flight, though all

_ _ _.•. _. ., .wry ..'___ .':.-:R ,/^
U .^

96

must be of concern for possible long duration space applications.

197H decays b electron capture Hg y y p giving a 77 keV gamma ray. Though the

thermal cross section is high for the 196 Hg (n,Y) 197 Hg reaction, the

natural abundance of 196 Hg is only 0.15% and therefore its contribution

is negligible. Based on the thermal cross section for Hg as a whole and

assuming 6% (natural abundance of 204 Hg)of the Hg undergoing neutron

capture end up as 205 Hg, the beta decay response spectrum is a factor

of -8 less than that produced by 1281.

Elastic scattering of neutrons is the main contributor to the

neutron induced valid background in the energy range 40-80 keV. Though

this rate is less than an equally photon efficient NaI detector due to

the higher atomic mass, it is an intrinsic background limit and only

large amounts of neutron shielding would decrease this background

source.

7.3 Charged Particle Induced Background

High energy charged particles wi17 also contribute to the background count rate of an astronomical X-ray detector. The source of the

particles are the primary cosmic rays, (mostly protczs and alphas) and

the secondaries they produce in the atmosphere and nearby matter

(mostly electrons and mesons). Due to the geomagnetic cutoff, only

protons greater than - 3.5 GeV reach the upper atmosphere above

Palestine, Texas (magnetic latitude a-42°N; cutoff rigidity 4.5 GV;

Peterson 1975)• Almost all of the charged particleo are highly relativistic (E>>M0 c 2 ) and therefore highly penetrating, resulting in geometric simplifications. Protons directed at the HgI 2 crystal will

reach it and deposit -400 !:eV ( minimum ionizing) as well as depositing

U

T
ri I (r

is M+t

97

r'+

20-50 MeV in the BGO on the way in or out. The total count rate in the

HgI 2 is therefore determined by the charged particle flux though veto

efficiency approaches 100%. The effects of charged particles that can

increase the valid background are decays of spailation products and

bremsstrahlung photons,

The flux of high energy primaries in near earth space is isotropic

and to first order the effective solid angle 0, for a detector close to

the earth, is 2n steradians. Of the primaries, 15% are alphas by

number and a few percent are CNO and heavier elements, the rest are

protons. Above 3.5 GeV the flux can be modelled as a power law (see

Meyer 1969 for review).

dN p

dN= 6.3xE

-1.6

 protons/cm 2 sec sr MeV

E>3500 MeV

This leads to an expected integrated count rate in the HgI 2 of .28

cts./cm2 sec due to primaries including the cosine a dependence of the

effective area (plus .03 cts./cm2 sec due to alpha particles).

The electron flux at high balloon altitudes (-3 gm/cm 2 ) is dominated by albedo electrons produced by primaries in the atmosphere below.

The downward moving electrons are ,just the re-entrant albedo electrons

and to a good approximation can be set equal to the upward moving flux

(Daniels 1974). Again a power law fit to this omnidirectional differential vpectrum above Texas is (Verna 1967):

d

.05E-1'3 electrons/cm2 sec sr MeV

10 McV<E<1000 MeV
,G

Using the range of electrons in BGO to determine the lower energy bound

(33 MeV from sides and 80 MeV from bottom) leads to an integrated count

rate in HgI 2 due to electrons of .32 cts/sec cm 2. The integrated count

rate is not very sensitive to the assumed low energy cutoff (i.e.

geometry) while the measured fluxes have disagreed by factors of 2-5

(Daniels 1974).

All charged particles reaching the HgI 2 detector must pass through

the BGO and be vetoed except those that pass through the passive

Pb-W-Sn shield o^*er the aperture. The high energy particles will

continue through the top detector and interact in the bottom HgI 2 and

probably the BGO as well. The only particles that will leave energy

between 40 and 80 keV in the top detector alone are electrons in a

small bandwidth with initial energy of -8 MeV. The estimated count

rate between 40 and 80 keV is .06 cts/hr, a negligible contribution.

Though charged particles dominate the total count rate of HgI 2 at

float, their valid prompt contribution is nil.

7.3 .1 • Spallation and Bremsstrahlunk

Though prompt charge particle interactions will be vetoed, the

spallation daughter products with their delayed decay can be a problem,

especially in orbit where the trapped protons and long flight duration

can lead to a significant background (Peterson 1975). The geomagnetic

cutoff and short duration (..24 hours) of balloon flights combine to

decrease this contribution to a small fraction of the valid background.

The thin HgI 2 crystals will be most sensitive to the radiation of

daughters that decay by beta emission (plus and minus) and electron

Nor _ ®^
^carrrw nisa. ^-^ 0

L

99

I

capture Most of the high energy gammas will escape and/or interact in

the BGO.

The spa'. lation of iodine has been well studied due to the many

detectors containing iodine ( NaI, CsI) flown in orbit as well as

irradiation experiments (Dyer 1980). A quick order of magnitude

calculation can be done using geometric cross sections and assuming

(worst case) every proton interaction with an iodine or mercury nucleus

decays later by electron capture. These crude assumptions give background features at 30 - 34 keV (iodine daughters) or 76 - 80 keV ( mercury

daughters) with intensities of .8 cts/hr and .5 cts/hr, respectively.

To reach this equilibrium rate would take weeks in orbit (Dyer 1980).

Most of the electrozi capture decays also emit gamma rays ths. , t have a

high probability of interacting in the BGO shield and being vetoed.

Though not required for a short balloon flight, a detailed analysis of

the proton induced spallation, peripheral, fission and evaporation

cross sections of the seven stable isotopes of mercury should be

attempted to dc'

ormine the equilibrium counting rate for HgI 2 in orbit.

A cursory look at the semi- empirical cross section formulae ( Silberberg

1973) shows that the most probable product isotopes ( mercury and

tellerium) are either stable or have relatively long half lives (-days)

so that this possible background source cannot be measured in a

. 1 standard" balloon flight.

An indirect method. for an electron to deposit energy in the top

HgI 2 crystal without being vetoed is by bremsstrahlung in the passive

graded shield and then stopping or backscattering. To estimate this

effect, the active volume used is the area of tin "seen" by the HgI2

i
^

rr

i.^

^e

d

w

ik

c

p

t

L

^it

p!

4

I

100

times one attenuation length of the emitted photon. Photons created

above this volume are self-absorbed in the passive shield. Using the

relativistic bremsstrahlung cross section (Jackson 1975, pg. 713) and

electron flux in the field of view gives .06 cts/hr between 40-80 keV

due to electrons between 8 and 9 MeV. Higher energy electrons will

continue through to the active shields and veto their own bremsstrahlung radiation.

Because relativistic charged particles interact strongly with

matter they are easily detected and vetoed though not easily stopped.

Only through indirect means (spallation) can they deposit energy in the

top HgI 2 crystal without being vetoed. A typical balloon flight is not

long enough for these reactions to build up though longer duration

satellite flights must be concerned with this source of background.

r*y

t

I

L

r,%W-
c

101

t

8.0 RESULTS OF BACKGROUND MEASUREMENTS AT SEA LEVEL

AND AT P'LLOON ALTITUDES

8.1 Sea Level Background

The radiation envirionment on the ground consists mostly of high

energy gamma rays due to the decay of naturally occurring radioactive

isotopes such as potassium 40 and uranium and thorium daughter products. Primary cosmic rays and neutrons are shielded by the atmosphere

but secondary muons do reach ground level. The gamma ray intensity is

a strong function of the type of building where the measurements take

place. Wood frame buildings are usually less active than concrete

buildings due to the amount of radionuclides in sand and cement.

Long term background spectra were accumulated with the prototype

detector gondola in many locations and many configurations. Fig. 8.1

shows the normalized count rates for spectra obtained in: Building 37,

5th floor, MIT; a ground floor preparation room at the National

Scientific Balloon Facility in Palestine, Texas; and the iron cave in

the basement of Building 2 at MIT. Integrated count rates are compared

in Table 8.1 along with the BGO shield count rate. The iron cave

results provide an upper limit to the intrinsic background of the

prototype detector (background due to radioactive material used in the

construction of the detector assembly itself) though extrinsic muons

are more likely to be the cause of this background. The gamma ray

contribution can be scaled with the bismuth K line observed in the veto

a

spectrum.' The veto ratio (the number of vetoed counts divided by the

total counts) increased in the iron cave as the bismuth line decreased.

This would be expected if the iron cave background is dominated by

T

It

ry

T

J
m

X

N

E

U

U

m

m

m

U

11)

I®

v

W

H

z

0

U

0

z

C

s

U

M

II i 102

wGINAL PAGE fl

OF POOR QUALITY

Balloon Lab

MIT Bldg.37

L m L Ln

ENERGY (keV)

Figure 8.1 Sea level valid background spectra in top mercuric iodide

detector for various experiment sites.

Y

5

4

3

2

1

e

5

4

3

2

1

0

5

4

3

2

1

e

m

V
183

Table S.1

Measured Background Count Rate of BGO-Shielded

Top HgI 2 Detector for Various Surface Locations

Valid Count Rate (40-80keV) Veto Bismuthtt BGO

Location (ots/sec cm2 keV) Ratio K X-raj Count Rate

Balloon Lab 1.8x10-5 91% 14x1O-4 1231 cts/sec

MIT Bldg. 37

Balloon Base 1.0x10-5 B9% 5x10-4 791 cts/sec

Palestine, TxIron Cave .2x10-5 94% 1x10-4 527 cts/sec

MIT Bldg. 2

t Ratio of number of vetoed counts to total counts (veto + valid)

over the energy range 40-160 keV

tt Counts/sec cm observed in the vetoed spectrum

-- J
u

1!!

C^.

L

4

r

^i

k^.

k^ y

i

1'

185

ORIGINAL PAGE Is

Bi+Hg K X-ray OF POOR QUALITY

30

BGO Veto Off

Bottom Veto Off

20

10

0

Ln

m m m

m in 

m

mm

m

- .. N N

ENERGY (keV)

Figure 8.2 Valid background spectra in top mercuric iodide detector

shoving the effectiveness of the anticoincidence shields.

3

m

X 10

fV

E

U

U 0

m

N

4J 30

U

lf)

m

20

v

LLLI

F—

10

z

0

C

0

Co

z

C) 30

0-'

Y

U

m

20
4 i

106

nRtr.,jNv IC FdCr ra

OF P0013 QUALITY

1500

HgIZ Flight Detector Calibration Spectrum

M

1000

z

0U

500

0

M M Ln ^ 00 m

CHANNEL NUMBER

Figure S. Actual mercuric iodide flight detector calibration spectrum

using a "'Am radioactive source. The energy resolution at

300V (flight bias) was 5.5 keV FWHM at 60 keV.

r

i

!i

ALIT
{ I I

197

1^1

so the quantization error of one channel/keV was negligible. This

detector was chosen for its stability and size and operated at 300

volts to decrease leakage current "bursts" that mimic background

counts. Some energy resolution was sacrificed since the dominant

background contributors produce a continuum of counts vs. energy.

8.3 Balloon Flights

The small prototype detector gondola was flown twice from

Palestine, Texas. The first time, flight 1282P, was a "piggyback"

flight in which the small gondola was strapped on the side of a large

(1400 kgs.) payload flown by Martin Isreal's group at Washington

University, St. Louis. The first flight was launched at 12:25 UT, May

8, 1982. It reached float altitude at 14:48 UT at an altitude of 38 km

(3.8 gm/cm 

2

vertical atmospheric column density). The flight lasted

35.5 hours and was successfully recovered in Mississippi. Analysis of

the housekeeping data showed that the batteries began to fail after 13

hours, probably due to a defective lithium cell since they were expected to last 35 hours.

Another problem observed in the first flight was multiple triggering on energetic events due to extreme saturation of the voltage

amplifier.. The extra events had a characteristic signature of preferentially occurring with a channel number of 4N-1. This effect was

duplicated in the laboratory by applying large pulses to the input of

the voltage amplifier as would have occurred during flight if a very

energetic particle interacted in the detector. This was a small effect

in the veto and bottom spectra but it dominated the valid spectrum,

increasing the count rats by a factor of 5• Two modifications were

1" J
made to the electronics. Diodes were placed at the input and on the

feedback loops of the voltage amplifier so that large pulses from the

preamps would only slightly saturate the last stage. In addition, a

time constant was changed in the event logic so that the analog to

digital converter would not accept another pulse for 7504sec after it

converted a pulse. Despite this lock-out period, the additional dead

time was negligible (.07% for a count rate of -1 ct/sec). The change

in the voltage amplifier required a new energy calibration (see above).

The second flight, 1296P, was a solo flight for the small gondola

using a small (62x10 3 m3 ) balloon. The launch was at 12:00 UT on June

' 21, 1982 and a float altitude of 40.2 km (2.8 gm/cm 2 vertical atmospheric column density) was reached at 14:01 UT. The flight was terminated at 21:00 UT and reached ground 43 minutes later -60 km Northwest

of Pecos, Texas. The payload struck a 130,000 volt power line, resulting in the parachute being ripped apart and the payload free-falling

the last 15 meters. Though the styrofoam burned on the outside, the

payload received only minor mechanical damage and the data tape was

recovered intact. The free fall also demonstrated the ruggedness of

BGO which operated flawlessly on its return to the laboratory.

8.4 Flight Results and Comparison to Predictions

The background spectra of flight 1282P for both HgI 2 detectors,

top and bottom, are shown in Fig. 8.4. Only the vetoed spectrum of the

top detector is presented because the valid spectrum was contaminated

i

with spurious events due to multiple triggering of the event logic by

1.

energetic paticles. The spectra from flight 1292P are presented in

Figs. 8.5 and 8.6. The electronic modifications made between flights

p
109

FLIGHT 1282P MAY 8, 1982

Top H 9 I 2 Detector

Vetoed Spectrum

t

(•.I

...(•r1 .. l:,., ^^ .:fin

20

M

^m

0

L.LJ

f^

20 OU

O

Z

O

C^

U

Y

U

m 10

ri<

^Y

K'

W^

N

10

U

Bi+H9 K X-ray x 1 sigma

U

0L i

m

m U-) m Ln

ENERGY (keV)

Figure 8.4 Vetoed and bottom spectra from flight 1282P. The veto rate of

the top detector was .15 counts/sec over the energy range 40

to 80 keV. As can be seen, the energy resolution of the bottom

detector (used as an anticoincidence shield) was substantially

broader than the top detector.

i,
v

^^ I

110

ORIGINAL PAGE 19

OF POOR QUALITY

20 FLIGHT 1292P JUNE 219 1982

Top H9 I 2 Detector

Vetoed Spectrum

Bi+H 9 K X-ray s 1 si gma

}a^

CV

U 10

U

a^

N

U

M

0

w

Fe

20 0

U

D

Z

O

O

L^

U

10

V

ryS^

C

P

f'.

i

0

® n m

CM

^-+ N N

ENERGY (keV)

Figure 8.5 Vetoed and bottom spectra from flight 1292P. Note the

decrease in count rate from the previous flight as well as

the resolved bismuth K-beta X-ray at 87 keV in the vetoed

spectrum.

JJ
N FLIGHT 1292P JUNE 21, 1982

CD

25

E

U

e

^"

30

III

ORIGINAL PAGE Irl

OF POOR QUALITY

U

N

y

20

U

Ln

I

L1J 15 Top HgI 2 Detector

Valid Spectrum

z

O 10

I

I -^-+++

CM

ENERGY (keV)

Figure 8.6 Valid spectrum from flight 1292p . The measured background

rate was 4.2x10 -6 counts/sec cm' keV (40-80 keV). There is a

possible feature at 59 keV (see text). The increase in count

rate below the 40-80 keV range of interest was due to detector

"bursts" (see text).

QZOO

Y 5

m

01

m

Q Cb

m

N
112

i I

were successful in suppressing the multiple triggering and the valid

spectrum (Fig 8.6) is free of the single channel peaks.

All these spectra were integrated over time starting at float and

ending at flight termination (or power failure for flight 1282P).

Arrival at float altitude was determined by when the count rate versus

i

time stabilized as shown in Fig. 8.7. The two maxima in the count rate

vs. time plot are due to the Pfotzer maximum, the altitude (_17 km.) at

which the flux of cosmic ray secondaries reaches its peak. The payload

passed through this region slowly on the way up and quickly on the way

l

II down.

The vetoed spectra for both flights are very similar in that they

are featureless except for the bismuth and possibly mercury K X-ray

line features. The K beta X-ray of bismuth was also resolved in the

second flight. The overall count rate decreased by about 25 to 35%

from the first flight to the second. This occurred in both detectors,

top and bottom. Changes this large cannot be explained by the differences in altitude between the two flights. As can be seen in Fig. 8.7,

the count rate was 30% higher -20 minutes before float which corresponds to a difference in altitude of 6 km. in contrast to the actual

difference of 2 km. in float altitude between the flights. The only

other differences between the two flights was the change in the event

logic electronics and the absence of the adjacent 1400 kg. payload in

YP

the second flight. Multiple triggering only accounted for a small part

n

of the higher count rate in the first flight based on the number of

events in the single channel peaks. A massive payload above the

detector as in flight 1282P, however, would produce an increase in the
4

U..

I

113

ORIGINAL PAGE R!

OF POOR QUALITY

FLIGHT 1292P

INTEGRAL COUNT RATE VS. TIME

Bottom Detector (E>28keV)

f

3000

0

W M

0 N 0

N

M m v®i m

•-^ ^ N i ^ m v <' V'1

BOCK NUMBER

10 Blocks = 11 minutes

LnY

0 2000

J

M

m

W

a.

Ln 

1000

z00U

Figure 8.7 Count rate vs. time of the bottom detector during flight 1292P

showing both ascending and descending passages through the

Pfotzer maximum and the steady count rate at float altitude.

V
o

114 Y

local radiation flux due to interactions with the primary cosmic rays.

z

If the massive payload is responsible for the higher count rate in the

first flight it would be equivalent to an atmospheric slab of 9 gm1cm2

vertical column density.

As expected, the total number of valid background counts accumulated during the second flight was small (31 counts, 40 to 80 keV) and

the spectrum is featureless. There are a few extra counts (6 cts.

f total) close to 59 keV. For a flight of this duration, 2.6 counts

would be expected at this energy from inelastic neutron scattering from

the 127 I nucleus. Unfortunately, the tungsten K-alpha X-ray energy is

also 59 keV so this barely significant feature might also be fluorescence of the passive graded shield or a combination of both. Based

I on their timing signature (many simultaneous low energy events), the

excess number of counts at very low energies are due to leakage current

"bursts".

Table 8.2 contains the float background count rate predictions of

i

Section 7 as well as the measured values. The valid count rate agrees

surprisingly well. The gamma ray prediction of the Monte Carlo simulation (Wood 1982) can be directly tested against the strengths of the K

alpha and beta line features of mercury and bismuth in the vetoed

spectrum. During flight 1292P, the measured value was 25 p smaller than

4

predicted, which might indicate that the gamma ray contribution was

overestimated. This is also indicated in the energies above 160 keV.

If the gamma ray prediction is scaled down to match the measured

results for the line features and the higher energies then the discrepancy at lower energies would be enhanced. This discrepancy

F^
115

c+}

probably indicates that the neutron contributions at lower energies

were underestimated, which could easily be the case, considering the

uncertainties in the neutron fluxes and cross sections.

E.

VMg./
ft

118

ORIGINAL PAGE 18

OF POOR QUALITY

Table 8.2

Comparison of Measured and Predicted Count Rates at Float

in the Top HgI 2 Detector for Flight 1292P

Energy (keV)

40 - 80

80 -120

120-160

j 160-200

}(

200-240

Valid Count Rate (10-

5

cts/sec cm keV)

Predicted Measured

2.9 4.2 + .7

1.6 2.6 + .6

1.0 1.9 + .5

1.8 1.5 + .5

1.7 1.0 .4

IBismuth and Mercury K X-ray line featurett

Predicted Measured

3.6 x 10 -2 cts/sec cm 2.7 x 10-2 cts/sec cm2

t Sum of predicted background due to gamma rays and neutrons

(Tables 7.1 and 7.2 respectively). The gamma ray contribution

was determined by a Monte Carlo simulation while the neutron

contribution was calculated analytically.

tt Present only in vetoed spectrum

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117

9.0 CONCLUSIONS

Previous to the research work described in this thesis, HgI2

showed great theoretical potential for use in hard X-ray astronomy.

The combination of high energy resolution at room temperature plus a

high photon quantum efficiency seemed ideal for many balloon and

satellite detector applications. Yet, experience in the field operation of HgI Z was minimal and background count rate information at high

altitudes did not exist. Development of HgI Z specifically for use as

an astronomical X-ray detector here at MIT over the last 3 years has

changed this situation. Procedures for handling and packaging detectors have been deve) ned that produce stable and dependable HgI Z detectors for use in the field. HgI Z detector assemblies have been constructed and flown at balloon altitudes and the detectors operated flawlessly un all occasions.

A major limitation to the application of HgI Z in X-ray astronomy

is the inability to fabricate individual detectors larger than -1 cm2.

A scheme to overcome this limitation involves the use of X-ray concentrators to increase the effective area of a shielded, low-background

HgIZ detector. The concentrator consists of a paraboloid shell lined

with lithium fluoride crystals that Bragg diffract on-axis incoming

X-rays to the focus of the paraboloid, where a small area detector is

located. To establish the expected sensitivity of this concentrator

/detector design, the background counting rate of the shielded HgIZ

detector at balloon altitudes had to be determined. A balloon-borne

gondola containing a prototype detector assembly was designed and

constructed to measure this background rate. The prototype detector
1:8

consisted of thin HgI 2 detectors surrounded by a large bismuth germanate scintillator operated in anticoincidence. This gondola was flown

twice in the spring of 1982 from Palestine, Texas.

The second flight of this prototype instrument established a

background counting, rate for a bismuth germanate-shielded mercuric

iodide detector of 4.2±0.7 x 10 -5 counts/sec cm keV over the energy

range of 40-80 keV. This measurement was within 50% of the predicted

value. The prediction was based on a Monte Carlo simulation of the

detector assembly in the radiation environment at float altitudes

(-40km). The Monte Carlo ,zimulation can now be used t investigate the

detector design paramete:rs (geometry, discriminator levels, etc.) to

attempt to lower the background rate even further.

Based on a background rate of 4.2 x 10 -5 ets/sec cm 2 keV, calculations indicate that with 4 concentrators and 4 detectors systems a

celestial X-ray source with 1/100 the intensity of the Crab Nebula

i

could be detected in 1.5 hours. This excellent sensitivity is achieved

by simultaneously reducing the background by B'M shielding a HgI2

detector, and increasing the effective area using hard X-ray concentrators.

The prescription for reducing the background rate by a factor of

-10 over previous balloon-borne hard X-ray detectors can be summarized

very simply: use very high quantum efficiency detectors shielded by

dense, high Z active scintillators. Because they are well suited to

this experimental approach, HgI 2 and BCO are both likely to become

major detector materials in X-ray astronomy. Furthermore, their ease

of handling and improved performance at moderately reduced temperatures

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enhance their desirability for balloon and satellite use.

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p.; 8.