the anticoincidence system of the pamela satellite ... -...

The Anticoincidence System of the PAMELA Satellite Experiment

Design of the Data Acquisition System and Performance Studies

Johan Lundquist

AKADEMISK AVHANDLING

som med tillstånd av Kungliga Tekniska Högskolan framlägges till offentlig granskning för vinnande av teknologie doktorsexamen.

fredagen den 9 december 2005kl 14.15 i sal FB42,

AlbaNova Universitetscentrum, KTH Partikel- och Astropartikelfysik,

Roslagstullsbacken 21, Stockholm

Avhandlingen försvaras på engelska

Kungliga Tekniska Högskolan

Fysiska Institutionen

Stockholm 2005

Johan Lundquist: The Anticoincidence System of the PAMELA Satellite Experiment -Design of the Data Acquisition System and Performance Studies

AbstractPAMELA is a satellite-borne cosmic ray experiment. Its primary scientific objective is tostudy the antiproton and positron components of the cosmic radiation. This will be donewith unprecedented statistics over a wide energy range (~10MeV to ~100GeV). ThePAMELA experiment consists of a permanent magnetic spectrometer, an electromagneticcalorimeter, a Time-of-Fight system, a neutron detector and a shower tail catcher. An anti-coincidence (AC) system surrounds the spectrometer to detect particles which do not passcleanly through the acceptance of the spectrometer. PAMELA will be mounted on a RussianEarth-observation satellite, and the launch is scheduled for 2006. The anticoincidence sys-tem for PAMELA has been developed by KTH, and consists of plastic scintillator detectorswith photomultiplier tube read-out. Extensive testing has been performed during the devel-opment phase. Results are presented for environmental tests, tests with cosmic-rays and par-ticle beams.

The design of the digital part of the AC electronics has been realised on an FPGA (FieldProgrammable Gate Array) and a DSP (Digital Signal Processor). It records signals fromthe 16 AC photomultipliers and from various sensors for over-current and temperature. Italso provides functionality for setting the photomultiplier discrimination thresholds, systemtesting, issuing alarms and communication with the PAMELA main data acquisition sys-tem. The design philosophy and functionality needs to be reliable and suitable for use in aspace environment.

To evaluate the performance of the AC detectors, a test utilizing cosmic-rays has been per-formed. The primary aim of the test was to calibrate the individual channels to gain knowl-edge of suitable discriminator levels for flight. A secondary aim was to estimate the ACdetector efficiency. A lower limit of (99.89±0.04)% was obtained.An in-orbit simulation study was performed using protons to estimate trigger rates and in-vestigate the AC system performance in a second level trigger. The average orbital triggerrate was estimated to be (8.4±0.6)Hz, consisting of (2.0±0.2)Hz good triggers and(6.4±0.5)Hz background. Inclusion of the AC system in the trigger condition to reducebackground (for the purpose of data handling capacity) leads to losses of good triggers dueto backscattering from the calorimeter (90% loss for 300GeV electrons and 25% for100GeV protons). A method, using the calorimeter, for identifying backscattering eventswas investigated and found to reduce the loss of good events to below 1% (300GeV elec-trons) and 5% (100GeV protons), while maintaining a background reduction of 70%.

Descriptors: PAMELA, satellites, astroparticle physics, electronics, anticoincidence sys-tem, trigger system

Table of Contents

Introduction 5

Authors Contribution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5Outline of The Thesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5

Chapter 1. Cosmic-ray Physics 7

1.1 Cosmic-Ray Particles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71.1.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .71.1.2 Composition and Energy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81.1.3 Fermi Acceleration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .101.1.4 Cosmic-Ray Propagation. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .121.1.5 Local Effects . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .13

1.2 Search for Antimatter in the Cosmic Radiation . . . . . . . . . . . . . . . . . . . . 151.2.1 Heavy Elements. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .151.2.2 Antiprotons . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .161.2.3 Positrons . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .181.2.4 The Future . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

Chapter 2. The PAMELA Experiment 21

2.1 The Resurs–DK1 Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 222.2 The PAMELA Apparatus. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23

2.2.1 The Time-of-Flight System. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 242.2.2 The Magnetic Spectrometer . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .252.2.3 The Calorimeter. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 272.2.4 The S4 Scintillator. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .312.2.5 The Neutron Detector . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .312.2.6 The Trigger . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .312.2.7 The DAQ System . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

2.3 Current Status . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 34

Chapter 3. The Anticoincidence System 37

3.1 Purpose . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 373.2 System Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 483.3 The Scintillators . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 403.4 The Photo Multiplier Tubes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 403.5 Analog Front-End Electronics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 423.6 LED Monitoring . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 443.7 Environmental Tests. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44

3.7.1 Vibration Tests . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .453.7.2 Thermal Cycling Tests . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .46

2

Chapter 4. Anticoincidence Digital Electronics 49

4.1 General Specifications . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 494.2 System Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50

4.2.1 Interface Signals . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 514.2.2 FPGA Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 514.2.3 The DSP . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 52

4.3 FPGA: IO Block. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 544.3.1 DS-Link . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 554.3.2 The Data Receiving State Machine (DRSM) . . . . . . . . . . . . . . . . . . . . . 664.3.3 The Data Transmitting State Machine (DTSM) . . . . . . . . . . . . . . . . . . . 604.3.4 Control of the Data Flow. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .60

4.4 FPGA: Register Block . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 614.4.1 The Address Space . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .614.4.2 Operations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .624.4.3 Registers . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .63

4.5 DSP–FPGA Communication . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 644.6 FPGA: Physics Block. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

4.6.1 Standard Pulses . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .644.6.2 Shift-Registers and Trigger Delay. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .644.6.3 Counters & LED Flashing. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 674.6.4 DAC Programming . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .674.6.5 ADC Reading . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

4.7 Trigger-Busy Logic . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 684.8 Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69

4.8.1 Initialisation. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .694.8.2 Operation in Physics Mode . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 794.8.3 Calibration Procedure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .704.8.4 Status and Alarm . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .71

4.9 Radiation Issues . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 724.9.1 Long-Term Effects . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 734.9.2 Single Event Effects . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .73

4.10 Performance of the Digital Electronics . . . . . . . . . . . . . . . . . . . . . . . . . 74

Chapter 5. Performance of the AC Detectors 75

5.1 Experimental Set-up . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 755.1.1 The Drift Chamber (DC) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 755.1.2 Trigger . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .765.1.3 Charge Measurement. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .785.1.4 Data Acquisition . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .785.1.5 Track Reconstruction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .895.1.6 Data and Selections . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .81

5.2 MIP-Scale Calibration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 815.2.1 Pulse-charge Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 825.2.2 Charge to Pulse-height Conversion. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84

5.3 Efficiency Measurements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 845.3.1 Particle Scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .865.3.2 Event Classification. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .875.3.3 Inefficiency and Background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .90

Table of Contents 3

5.3.4 Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .905.4 Position Dependence . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 935.5 Detector Stability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 94

Chapter 6. Studies of a Second Level Trigger 97

6.1 Proton Flux . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 976.2 The First Level Trigger . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 996.3 The Simulation Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

6.3.1 Particles and Energies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .1016.3.2 Detector and Geometry . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

6.4 First Level Trigger Rates . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1036.5 The Second Level Trigger . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

6.5.1 AC Rejection and Backscattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .1056.5.2 Calorimeter Condition. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .107

6.6 Test-Beam Results and Simulation Validation . . . . . . . . . . . . . . . . . . . . 1096.7 Implementation. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

Chapter 7. Acknowledgments 111

Chapter 8. Bibliography 113

Introduction

Authors ContributionWhen I started at KTH I got involved in work surrounding the liquid argon calorimeter sys-tem for the ATLAS experiment at CERN. That work resulted in a licentiate thesis [45]which centres around tests performed to gain understanding of how the radiation environ-ment in the ATLAS detector will effect the components of the optical read-out link for theliquid argon calorimeter system.

After this I started working on the PAMELA experiment, where I assumed responsibilityfor the design of the digital part of the data acquisition for the anticoincidence system. SinceI had no prior knowledge of logic design, a period had to be dedicated to learning DSP andVHDL programming, in order to be able to build the system which now a part of the PAME-LA experiment, and is presented in this work (see Chapter4). During the this developmentvarious environmental tests of the anticoincidence detectors and electronics were performed(see end of Chapter3). In most of these tests I was responsible the data acquisition.

To evaluate the performance of the AC detectors, I planned and performed a cosmic-ray testat the University of Naples (see Chapter3). The primary aim of the test was to calibrate theindividual channels in the anticoincidence system on to obtain knowledge of suitable pulsediscrimination levels. The test allowed the efficiency of the detectors to be evaluated.

The latest work I have made in the PAMELA experiment concerns trigger rate simulationsof the experiment in space (see Chapter6). The aim has been to estimate and find a methodfor using the veto capability of the anticoincidence system on the significant background,without losing good events due to backscattered particles. This work has involved both sim-ulation and test-beam studies.

Publications related to the work in this thesis: [43,44,54,56,57,58,59]

Outline of The ThesisIn Chapter1 a brief introduction to the field of cosmic-ray physics is presented. Chapter2gives an overview of the PAMELA experiment and the different sub-detectors. Chapter3 isdedicated to the anticoincidence system, describing its purpose and design. A presentationof results from environmental studies is also made. Chapter4 gives a detailed description ofthe anticoincidence digital electronics, and in Chapter5 is presented results of anticoinci-dence performance tests. Chapter6 describes methods and results for trigger rate and back-scattering simulations of the PAMELA experiment in space.

Chapter 1

Cosmic-ray Physics

The term cosmic-rays, in everyday speech, often refers to the ionising particles that contin-uously penetrate our daily environment and interact with matter here on Earth, and whoseorigin can be traced back to sources outside the Earth’s atmosphere. This radiation is com-posed of secondary particles (mostly electrons and muons) produced in interactions betweenthe Earth’s atmosphere and the so-called primary cosmic-ray particles that impinges on it.These primary particles form an interstellar plasma that fills the Galaxy and is a part of whatis called the interstellar medium (ISM). What is referred to as primary cosmic-rays mightindeed be secondary particles in the sense that they can have been produced in interactionsbetween the cosmic-ray plasma and other constituent particles of the ISM. However, thecosmic-rays must have primary sources of origin. These are believed to be located mainlyinside the Galaxy. The study of cosmic-rays allows the mechanisms of production and prop-agation of particles within our Galaxy to be studied. The antiproton and positron compo-nents of the cosmic radiation also gives the possibility to explore physics beyond theStandard Model, which otherwise may only be possible by using detectors located at parti-cle accelerators.

1.1 Cosmic-Ray ParticlesEvery second around 1000 charged cosmic-ray particles per square meter hit the Earth’s at-mosphere. About 90% of these are protons, 9% are alpha particles and 1% are electrons, andthere are also some small abundances of heavier nuclei [28].

1.1.1 HistoryThe presence of these particles was discovered in 1912 by Victor Hess and his colleagueWerner Kohlhörster when they were conducting an experiment in which a hot air balloon,carrying an electroscope, was used to measure the altitude dependence of the natural back-ground radiation [33,37]. At the time this radiation was thought to originate from the Earth.However, in Hess’s experiment the electroscope discharged more rapidly as the balloon as-cended in the atmosphere which meant that the ionising radiation increased with altitude andtherefore must originate from outside the Earth’s atmosphere rather than from the ground.For some time it was believed that the radiation was electromagnetic in nature.

In 1929, an experiment of Walter Bothe and Kohlhörster based on coincidence measure-ments using two Geiger detectors allowed the charged nature of cosmic-rays to be deter-mined [20]. A few years later (1932) Carl D. Anderson used a cloud chamber anddiscovered the positron by measuring the curvature and energy-loss of cosmic-ray particles

8 Chapter 1 – Cosmic-ray Physics

in a magnetic field [4]. This was the first experimental evidence of the antimatter predictedby Dirac. That discovery was an important milestone in the development of particle physics.In the following years many new particles were discovered in cosmic-ray experiments, andcosmic-rays provided a source of particles for high energy physics research.

With the advent of particle accelerators, cosmic-rays lost their role as a particle source. Thestudy of the flux, composition, origin and propagation of cosmic-rays has gradually growninto an independent field of research. However, during the last decades the interplay be-tween cosmic-ray-physics and particle-physics has undergone somewhat of a revival. As-troparticle and cosmic-ray physics aims to answer fundamental physics questions connectedto phenomena like; neutrino oscillations, magnetic monopoles, dark matter and baryonasymmetry.

A wide range of experimental techniquesare being used to pursue this scientific pro-gram. At ground-level, large area detectorarray experiments are constructed to meas-ure air showers initiated by high energycosmic-rays. To study the primary cosmic-ray particles before they interact in theEarth’s atmosphere, high altitude balloonshave frequently been utilised. At the alti-tudes reached by these balloon-borne ex-periments, there is however still anoverburden of residual atmosphere whereparticle interactions can occur. The way toavoid this limitation is to use satellite-borneexperiments.

1.1.2 Composition and EnergyThe main component of the primary cos-mic-ray flux is hadronic in nature andseems to originate from within our Galaxy.The flux is approximately isotropic andranges between 2 and 4 cm−2s−1 at 1 astro-nomical unit. Galactic cosmic-rays (GCR)are mainly composed of hadronic particleswith the abundance: 90% H, 9% He and1% heavier nuclei. An electron componentof ~1% and traces of antiprotons (0.001%)and positrons (0.1%) are also present.

Hadronic ComponentThe hadronic component of the cosmic radiation has an abundance of different nuclei whichis almost the same as those in the solar system. This is illustrated in Fig.1–1 where the rel-ative abundance in the Solar system (dashed line) is compared to that of cosmic-rays. Both

Figure 1–1: Relative abundance of cosmic-rays (solid line) and of the elements present in the Solar system (dashed line). The numbers are normalised to the abundance of Si (=100).

CO

Ne

He

MgSi

S

A CaTi Cr

Fe

Ni

Mn

Co

V

KP

N

Cl

Sc

Na Al

F

B

Be

Li

106

105

104

103

102

10

1

10–1

10 –2

10 –3

10–4

10–5 2 4 6 8 10 12 14 16 18 20 22 24 26 28

NUCLEAR CHARGE NUMBER

RE

LA

TIV

E A

BU

ND

AN

CE

1.1 – Cosmic-Ray Particles 9

compositions show a clear even-odd effect (nuclei with an even atomic number are moreabundant than those with an odd one). This composition and the shape of the cosmic-ray en-ergy spectrum supports the assumption that most Galactic cosmic-rays derive from super-nova explosions (discussed later in this chapter).

Differences between Solar and GCR abundances are explained by spallation processeswhere carbon, oxygen and iron nuclei interact with the ISM as they travel through the Gal-axy. The two groups of elements Li, Be, B and Sc, Ti, V, Cr, Mn are almost absent in theSolar system while a significant amount of them is found in cosmic-rays. The abundancesobserved in cosmic radiation are interpreted as a result of the collisions of nuclei with theinterstellar medium and the subsequent fragmentation of the primary nucleus into lighter el-ements. The products of these spallation processes are Li, Be, B, in case of primary carbonand oxygen, while Sc, Ti, V, Cr, Mn are produced by primary iron nuclei.

Collisions of primary cosmic-ray nuclei with the ISM also cause production of long-livedradioactive isotopes, which can be used as cosmic-ray “clocks”. By comparing the observedabundances of these radioactive nuclei with the amount of stable secondary species in thecosmic radiation one can establish for how long cosmic-rays are confined in the Galaxy andinvestigate the density distribution of the gas encountered by the particles. It has been de-duced that an average GCR spends several million years wandering around our Galaxy be-fore reaching Earth. Because of their charge, most cosmic-rays propagating in the Galaxyhave an energy low enough to be deflected by the interstellar magnetic field, and once theyreach the Earth all directional information about their source has been lost.

The differential cosmic-ray flux is shown in Fig.1–2 as a function of energy. Above 10GeV/nucleon the energy spectrum of primary cosmic-rays can be described by a segmented pow-er law spectrum:

It spans over a large number of decades. For the highest energies (above 1019eV) the indexappears to be somewhat smaller. There are two breaks in the spectrum around 1015–1016eV(the “knee”) and 1018–1019eV (the “ankle”). Exactly what causes the knee is still unknown.Several different explanations have been proposed. Some connect the steepening of thespectrum to different acceleration mechanisms or to a change in the GCR propagation whichleads to a more rapid escape from the Galaxy, others explain the knee, in terms of differentproperties of the source itself [61]. However, neither of the models has become broadly ac-cepted.

The highest energy cosmic-rays measured to date have had more than 1020eV, equivalent tothe kinetic energy of a baseball travelling at approximately 100km/h.

ElectronsEnergetic cosmic-rays that collide with the ISM can interact and produce a large variety ofsecondary particles, some of which will decay. Electrons can be produced in these decays(e.g. from pions via the decay ). The interstellar secondary production mech-

dNdE------- E

α–∝ ; α2.7 E 1016eV<3.0 1016 E 1018eV<


anisms for electrons in processes like this one will yield an equal number of positrons. How-ever, the measured positron fraction , is found to be only a few percent in theexplored energy interval (below ~50GeV) [24]. This fact suggests that there is mainly a pri-mary production of electrons in the Galaxy. The subject of cosmic-ray positrons is discussedlater in this chapter.

Like the hadronic component, electrons interact in the interstellar gas, but in addition to thisthey suffer significant energy losses when they interact with the interstellar magnetic fieldand the microwave background photons. These additional interactions, which include emis-sion of synchrotron radiation and bremsstrahlung as well as inverse Compton scattering,causes the energy spectrum to be dependent on parameters such as: the average diffusiontime in the Galaxy, the average strength of the Galactic magnetic field and the energy den-sity of the background radiation. The energy losses due to inverse Compton scattering withthe microwave background radiation also limits the lifetime for electrons and thus the dis-tance they can travel. This implies that the sources must be intergalactic. Electrons (and pos-itrons) are thus the only cosmic-ray components for which extragalactic contributions canbe excluded.

1.1.3 Fermi AccelerationHow do cosmic-ray particles acquire their kinetic energy? Enrico Fermi developed a theoryfor stochastic acceleration of cosmic-rays [26]. When cosmic-ray particles “collide” with

(a) (b)

Figure 1–2: Energy spectra of cosmic-rays a) Total spectrum.[79] b) Individual spectrum for compo-nents H–C (from [25]).

H

He

C

Fe

10-1

10-1

10-4

10-4

10-3

10-2

100

102

10-7

10-7

10-6

10-5

10-8

10-9

10-10

10-13

10-16

10-19

10-22

10-25

10-28

102

109 10

310

210 10

410

610

11 105

1013 10

710

1510

1710

1910

21

Kinetic Energy (MeV/n)Kinetic Energy (eV)

Ankle

(1 particle per km2 and year)

Knee

(1 particle per m2 and year)

(1 particle per m2 and second)

Diffe

ren

tia

l F

lux (

m2 s

r s G

eV

)-1

Diffe

ren

tia

l F

lux (

m2 s

r s M

eV

/n)-

1

e+ e+ e−+( )⁄


moving clouds of plasma, they gain or lose an amount of energy proportional to the plasmavelocity and to the particle energy. Whether a particle gains or lose energy will depend onthe relative velocity between the particle and the plasma (v). Since there is no upper limit tothe final energy but only a lower one (E=0), there is however a net gain in the particle en-ergy after many encounters. As the energy gained at each encounter is proportional to (v/c)2,this mechanism is called second order Fermi acceleration. In many cases it leads to a powerlaw energy spectrum of cosmic-rays, but because of the quadratic dependence the processis quite slow and inefficient in comparison to what is needed to sustain the observed cosmic-ray spectrum.

A similar stochastic acceleration process is believed to take place in shock-fronts thatsweeps through the ISM with a velocity much faster than the speed of sound in the interstel-lar gas. The principle is very similar to the second order mechanism, but here an energy pro-portional to v/c is gained each time the particle crosses the shock front. This mechanism istherefore referred to as first order Fermi acceleration (illustrated in Fig.1–3). With succes-sive diffusive crossings of the front the calculated spectra reasonably well reproduces ameasured power law spectra with spectral index around −2. Taking into account that higherenergy particles escape the Galaxy more easily, this index could be modified to the observedvalue of 2.7.

First order Fermi acceleration is believed to be atwork in many different types of shocks: the termina-tion shock of the Solar and the Galactic winds, the ac-celeration shock near a super-massive black hole(which is believed to exist in the center of our Gal-axy) and the expanding envelopes of exploded super-novae. With the observed rate of supernovaexplosions in our Galaxy (approximately one every50–100 years), only a few percent of the total energyreleased in a typical explosion (~1056J) needs to betransferred to cosmic-rays to maintain the observedintensity.1

The supernova acceleration model works well belowthe knee, but it can not produce enough GCRs above1018eV [28]. Moreover, since the gyroradius in theGalactic magnetic field at ~1018eV is of the order ofthe size of the Galactic disk, a transition from Galac-tic to extragalactic cosmic-rays is expected to occuraround this energy. Experimental evidence for such atransition is an important question [11,31,23].

1. The transfer rate of energy from supernovae to cosmic-rays is balanced by the total power of the cos-mic rays leaving the Galaxy, which is estimated to be of the order of 1047W [12].

Figure 1–3: Schematic view of “first or-der” Fermi acceleration. A shock wavemoves through the intergalactic mediumwith velocity v which is much faster thanthe speed of sound in the interstellargas. It can be shown that, on average,the energy gained in each crossing (dE)is equal to (3/4)⋅E⋅(v/c) [28]

E

v

E+dE

random diffusionchock fro

nt


1.1.4 Cosmic-Ray PropagationOnce produced, cosmic-rays propagate through the ISM, experiencing several physicalprocesses: diffusion along magnetic field lines, scattering on magnetic field irregularities,spallation reactions with the interstellar gas and radioactive decays. Cosmic-ray propagationmodels take into account all these processes and attempt to explain the fluxes observed atEarth. Descriptions of two of the most frequently used models, the Leaky Box Model and theDiffusive Halo Model, are given below.

Both of these models rely on a general transport equation [28]:

where is the density of a given nuclear species of type i located at x and withenergy between E and . The first term in the right side describes diffusion, D beingthe diffusion coefficient. The second term is related to the energy gains or losses and thethird to convection with velocity u. The fourth is the source term and the fifth takes into ac-count possible losses of nuclei by collision or decay (at rate p). The last term represents cas-cade and nuclear fragmentation processes (v is particle velocity, ρ density, m particle massand σ is the reaction cross-section).

Leaky Box ModelThe simplest approach to the solution of the transport equation is the Leaky Box Model. Re-placing the diffusion term by and neglecting collisions and all the other energychanging processes as well as convection, the solution for a point source is:

in which is interpreted as the mean time spent by cosmic-rays in the Galaxy and con-sequently , is the mean amount of matter traversed by a particle of ve-locity ( being the density of the interstellar gas). can be found by fitting the observedratios of secondary to primary cosmic-ray nuclei (e.g. ). Several extensions to theLeaky Box Model have been proposed, which differ widely in the energy dependence of

and in the features of the containment box.

In the leaky box model the particles are evenly distributed in the Galaxy. Cosmic-rays dif-fuse freely in a confinement volume, which could be either the Galactic disk or the halo, andthey are reflected at the boundaries. They have constant probability per unit time of escapinginto intergalactic space at each encounter with the boundary, governed by the characteristicescape time from the confinement volume. Fig.1–4b illustrates the Leaky Box Model.

Diffusive Halo ModelA completely different approach to the propagation problem is to consider a non-constantdiffusion term in the transport equations [29,39]. In these diffusion models the density of

Ni∂t∂

-------- ∇ Di∇ Ni( )⋅ E∂∂ dEi

td--------Ni E( )– ∇ uNi E( )⋅–=

+ Qi E t,( ) piNivρm------

σi k, E E ′,( )dEd

-----------------------------Nk E′( ) Ed∫k i≥∑+–

Ni E x t, ,( ) EdE dE+

N τesc⁄–

N E t,( ) N0 t τesc⁄–( )exp=

τescλesc ρβcτesc= βc

ρ λescp p⁄

λesc


cosmic-rays is not homogeneous and the isotropy of the Leaky Box Model is lost. Amongthem the Diffusion Halo Model (DHM) is the most used. In this model it is assumed thatcosmic-ray sources are located within the thin Galactic disk and the escape into the halo andeventually into the intergalactic space is driven by diffusion. The velocity of cosmic-raysstreaming away from the halo depends on the diffusion coefficient and on the halo size. TheDiffusion Models deal with a more realistic physical scenario, while the Leaky Box Modelsare often preferred since for many purposes the results are equivalent but obtained throughsimpler calculations.

1.1.5 Local Effects

Solar ModulationThe plasma emitted from the Solar corona is called the Solar wind. It consists of particlesaccelerated by Solar flares or shock waves driven by coronal mass ejection. The solar windhas speeds of about ~350km/h. It reaches out beyond Pluto (see Fig.1–5a) and carries theSolar magnetic field. This magnetic field in turn deflects low energy particles of extra-Solarorigin and prevents them from reaching the Earth. The observed cosmic-ray flux at Earth istherefore anti-correlated to the Solar activity. The effects can be seen for cosmic-rays up toenergies of tens of GeV. In order to estimate the Galactic cosmic-ray flux in this region agood knowledge of the effects of the Solar activity is required. At energies below ~10MeVcosmic-rays spectrum is dominated by solar particles due to this effect.

The Sun’s activity shows a periodic behaviour with an 11 year time interval. The number ofsunspots is an observable parameter used to monitor the Solar activity. By measuring thecosmic-ray spectra over an extended period, information of the interstellar spectra and theeffects on particles of different sign of charge (since the Sun reverses magnetic polarity eve-

(a) (b)

Figure 1–4: Illustration of different models for cosmic-ray propagation in the Galaxy. a) the Leaky Box Model b) the Diffusive Halo Model

HALO

GALAXY GALAXY

HALO


ry 11 years) can be collected. The solar activity is anticorrelated to the neutron flux onground. The neutrons are used as a probe of low energy cosmic-ray interactions in the at-mosphere (see Fig.1–5b).

Geomagnetic Cut-OffThe radius by which a magnetic field of strength B deflects a charged particle is proportionalto and a factor that only depends on the incident angle between the particles momen-tum and the direction of the magnetic field. R is the particle’s rigidity defined by (pbeing the particle momentum and Ze the charge). Thus, particles with the same rigidity andincident angle follow the same trajectory in a magnetic field.

The geomagnetic field deflects particles approaching the Earth from outer space. For a givendirection in the sky there is a therefore a certain critical cut-off rigidity needed to overcomethe geomagnetic field in order to reach the top of the atmosphere. The full solution of theproblem is non trivial. However, the momentum needed for a vertically incident particle toreach the Earth’s atmosphere depends on the geomagnetic latitude , and can be estimat-ed by [40]:

(e is the electron charge). This equation takes into account that the cut-off is higher near theequator where the field lines are perpendicular to the motion of the particle, than at the poleswhere the magnetic field is parallel to the particle trajectory.

(a) (b)

Figure 1–5: a) The Solar wind forming the heliosphere will prevent some of the cosmic-rays from enter-ing the Solar system. b) The number of observed sunspots (bottom) and the measured atmospheric neu-tron flux (top). The anticorrelation between the number of sunspots and the neutron flux [71] is clearly seen.

95 96 97 98 99 00 01 02 03 04 05 06

3000

50

150

3200

3400

3600

PA

ME

LA

(Ja

n-0

6)

YearC

ount ra

te (

a.u

.)S

unspots

Plasma Sheet

Magnetopause

Magnetosheath

Solar Wind

Bow Shock

Polar

Cusp

R B⁄p Ze⁄

Lgeo

R14.9

e---------- L

4geocos≥ [GV]

1.2 – Search for Antimatter in the Cosmic Radiation 15

1.2 Search for Antimatter in the Cosmic Radiation

1.2.1 Heavy ElementsAccording to the Big-Bang theory, matter and antimatter should have been created in equalamounts in the first instants of the Universe. Thus, the evident asymmetry between matterand antimatter is a mystery. It has been suggested that this asymmetry is an effect of CP vi-olation in combination with out-of-equilibrium expansion of the universe and baryonnumber breaking mechanisms []. However, the asymmetry could also be spatial in nature,i.e. regions of antimatter could exist alongside regions of matter. According to theories, anyregion of the universe dominated by antimatter must be separated from matter dominatedregions on the scale of clusters or super-clusters of Galaxies. Otherwise intense gamma ra-diation would be expected from annihilation processes on the border of between the regions.According to present understanding the electromagnetic interaction of matter and antimattershould be identical (the photon is its own antiparticle), and thus antimatter can not be iden-tified by optical or radio astronomy observations. The detection of anti-nuclei heavier than

would however indicate presence of bulk antimatter in our Universe (e.g. productionthrough fusion in an “anti-star”), while finding could also indicate the existence of pri-mordial antimatter left over from the Big-Bang. No anti-nuclei have been found and theantimatter searches have only been able to set upper limits on the anti-helium–helium ratio(see Fig.1–6).

Figure 1–6: The ratio of anti-helium–helium in the cosmic radiation as function of rigidity (momentum/charge). So far no experiment has been able to detect any anti-helium nuclei and thus only been able to set upper limits.

HeHe

PAMELA EXPECTED


1.2.2 AntiprotonsAntiprotons have been observed in the cosmic-rays since 1979 by balloon-borne experi-ments [30,18]. Secondary antiproton production occurs when cosmic-ray nuclei undergo in-elastic collisions with the constituents of the ISM. In order to produce at least one antiprotonin a proton–proton reaction and at the same time conserve baryon number and charge, a min-imum of three additional baryons must be produced. An example is the reaction:

The production of antiprotons therefore requires high energies. From kinematic constraints,the production threshold is about 7GeV for the interaction above. One key feature, also dueto the kinematics of antiproton production, is that the antiproton spectrum produced by suchcollisions sharply decreases below 1GeV. This is often referred to as the “kinematic thresh-old” for antiproton production. The spectrum of antiprotons has a characteristic shape witha maximum at ~2GeV. This is the only strongly energy dependent production process incosmic-rays and this feature makes the antiprotons spectrum a unique probe of various phe-nomena.

Antiprotons can also be created at lower energies in proton–nuclei or nuclei–nuclei interac-tions, via so-called “sub-threshold” production [66,38,7]. Fig.1–7 shows the theoreticalcontributions to the interstellar antiproton source density spectrum, from various interac-tions between cosmic-ray particles and the ISM.

By measuring the spectral shape it is possible to measure the post-production accelerationor deceleration (e.g. Fermi acceleration) of the particles. Between 1 and 5 GeV the ra-tio should have a characteristic rise. The shape and position of the rise tells us whether ornot there is acceleration of the antiprotons in the ISM after they are produced and if they areproduced by collisions in the ISM.

Figure 1–7: The interstellar antiproton source density spectra resulting from the different interactions between cosmic-ray particles and target particles of the interstellar gas (from [63]).

p p+ p p p p+ + +→

p p⁄


Determination of the energy spectrum of antiprotons is essential for understanding the prop-agation of cosmic-rays as the interaction mean free path for antiprotons in interstellar spaceis much larger than the matter traversed by cosmic-rays in the Galaxy. This is especially im-portant when comparing propagation of light nuclei and antiprotons, since they have differ-ent production mechanisms even if they are both produced in collisions between cosmic-rays and the ISM. The major difference is that antiprotons are primarily produced by cos-mic-ray protons while light nuclei are produced by spallation processes and depends strong-ly on the heavier nuclei, which might have a different model of propagation.

“Exotic” primary processes for producing antiprotons; like the evaporation of mini blackholes [36,68,46], or the annihilation of dark matter particles in the Galactic halo [62,21,8],have been suggested. One popular candidate for the latter is the lightest supersymmetric par-ticle – the neutralino ( ) – which is stable due to R-parity conservation in some supersym-metric models. An accurate knowledge of the secondary antiprotons flux is crucial whentrying to isolate an exotic primary contribution.

Fig.1–8 shows a compilation of the current antiproton measurements together with sometheoretical models. The two solid curves are predictions of the interstellar secondary anti-proton flux based on the standard Leaky Box Model using different path lengths. Thedashed curve represents a calculation of the secondary antiproton flux derived from a diffu-sive model [8,9] based on proton data from the CAPRICE–94 experiment [17]. The dottedcurve is a prediction of the contribution from primary antiprotons produced by heavy neu-tralino (964GeV) annihilation [69], which would detected as a distortion to the secondaryspectrum.

Figure 1–8: The measured antiproton spectrum. adapted (from [13])

χ̃0


An experiment aimed at detecting antiprotons must be capable of charge identification andseparation of antiprotons from the high background flux of electrons that decreases fromabout 103 times the antiproton component at 1GeV/c to less than 102 above 10GeV/c.

1.2.3 PositronsWhile most of the observed electrons are believed to be of primary origin, the measured pos-itron component is consistent with a pure secondary production. The main argument for thisis (as mentioned earlier) that the measured positron fraction ( ), is found to beonly a few percent. Moreover, the spectral power law index for positrons below 15GeV isfound to be about −3.1 [10] and consistent with secondary production. The theoretical pos-itron energy spectrum can be estimated by starting from the pion and kaon production cross-sections and then applying propagation models which take into account the additional ener-gy losses [60,50,64].

In Fig.1–9 the experimental situation for positrons is shown. The measured positron chargeratio is displayed together with some theoretical predictions. The dashed line is from a cal-culation based on an older standard Leaky Box Model assuming a purely secondary originof positrons [60]. The solid line is for secondary production using a diffusive model [50].The dotted line is a prediction of a primary contribution based on the annihilation of neu-tralinos with a mass of 336GeV [5]. Some data shows an excess of positrons above the fluxexpected by the simple Leaky Box Model which may indicate a rise at energies greater than15GeV [6] even though the spectrum should be steepening in this region. Thus, clearly the

Figure 1–9: The measured positron charge ratio (from [13])

e+ e+ e−+( )⁄

e+ e+ e−+( )⁄


positron spectrum is not understandable in terms of our conventional model for cosmic-rayproduction and propagation.

One possible source of this excess is the annihilation of dark matter particles in the Galactichalo. Other sources include magnetic pair production at pulsars and pair-production by gam-ma rays interacting with optical or ultra violet photons [52]. Observation of positrons overa very large energy range should yield new insights into Galactic processes.

A difficulty in detecting positrons lies, as for antiprotons, in the charge identification andthat the positrons have to be well separated from protons. The ratio between protons andpositrons lies between 103 at 1GeV/c to about 5⋅103 at 10GeV/c and a reliable positron iden-tification therefore requires a proton rejection factor greater than 104.

1.2.4 The FutureAll data in Fig.1–8 and 1–9 is consistent with a purely secondary origin for both antiprotonsand positrons. However, most of the data comes from high altitude balloon experimentswhich are limited in statistics due to data collecting times of 1–2 days. Therefore the steep-ness of the antiproton and positron spectra at higher energies has so far limited the measure-ments to below ~50GeV. The balloon experiments also have to take into account theinteractions occurring in the residual atmosphere (~5g/cm2) above the experiments by esti-mations derived from simulations, which increase the systematic errors on the data points.

The PAMELA experiment, which will be described in the next section, is a satellite-borneexperiment for cosmic-ray studies and in particular the study of antiproton and positronflux. It will remove the limitations mentioned above by taking data over a much longer timeperiod (~3 years), covering a wider energy range and eliminate the influence of the atmos-phere. Expectations after 3 years is shown in both Fig.1–8 and Fig.1–9 [13]. The PAMELAdata points include error bars (only statistical errors), and should be able to resolve both aprimary component of antiprotons or positrons in the cosmic-rays as well as be used to val-idate the different propagation models. In Fig.1–6 is also shown the PAMELA expectationvalue on anti-helium rejection.

−♦−

Chapter 2

The PAMELA Experiment

Since Hess’ day, many balloon-borne experiments have been performed, aimed at measur-ing the cosmic-ray energy spectra and composition. As mentioned in the previous chapter,at the altitudes reached by these experiments (~40km) there is an overburden of ~5g/cm2

from the residual atmosphere where particle interactions can occur. Satellite-borne experi-ments avoid this constraint. An other advantage is that a satellite allow for much longer pe-riods of data-taking, usually several years, resulting in a larger statistical sample. A naturaldisadvantage is that one has no access to the experiment after it has been launched into orbit.The experiment hardware must therefore be extremely reliable.

PAMELA (Payload for AntiMatter Exploration and Light-nuclei Astrophysics) is a particledetector system which will be placed on-board a Russian Resurs DK–1 satellite [80]. Thelaunch is scheduled for 2006. The sun-synchronous, 350–600 km altitude semi-polar orbitwill allow low-energy cosmic-rays to be measured while near the poles. The primary objec-tive of PAMELA is to measure the energy spectrum of antiprotons and positrons in the cos-mic radiation. At least 105 positrons and 104 antiprotons per year are expected. The statisticsfrom PAMELA data will exceed what is available today by several orders of magnitude andwill allow significant comparisons between the competing models of antimatter productionin the Galaxy. The main observational objectives of PAMELA can be summarized as fol-lows:

• flux from 80MeV to 190GeV ( flux from 50MeV to 1TeV)• flux from 50MeV to 270GeV ( flux from 50MeV to 800GeV)• light nuclei (Z

22 Chapter 2 – The PAMELA Experiment

2.1 The Resurs–DK1 SatelliteThe satellite that will carry the PAMELA experiment will be the Russian Earth-observingsatellite Resurs-DK 1 [80]. It is mainly designed to take high resolution images of theEarth’s surface in order to collect information about sea surface status, ice situations, mete-orological conditions in the polar regions and information for natural resource studies. Thesatellite will be oriented toward Earth in order to perform these observations. The PAMELAexperiment will be in a, pressurized container, mounted on a mechanical arm, on the side ofthe satellite. This container will be oriented so that PAMELA faces away from the Earthduring data taking. It will only be moved to a downward position during launch and orbitalmanoeuvres. Fig.2–1 shows a schematic drawing of the satellite.

The orbit is elliptical and semi-polar (i.e. will not pass directly over the geographic poles)with an altitude between 350 and 600 km and with an inclination of 70.4°. It will traverseboth the inner and outer radiation belts when approaching the polar regions, and also passthrough the South Atlantic Anomaly (SAA). In these places the particle flux will increasesignificantly due to low energy particles trapped in the Earth’s magnetic field.

The mass of the satellite is ~10 tonnes, while the total mass of the PAMELA container withthe experiment installed will be ~750kg. The average power consumption for the satellite isexpected to be ~2kW, and for PAMELA ~380W.

Figure 2–1: Schematic drawing of the Resurs DK-1 satellite. The shaded container on the side of the satellite houses PAMELA and will be in this position when measurements are made. During launch and orbital manoeuvres the container will be in the downward position.To the right is shown the Soyuz rock-et.

47m

2.2 – The PAMELA Apparatus 23

The launch-vehicle is a three-stage Soyuz–U rocket [87] (Fig.2–1b); which is the type usedfor manned missions. It has a high launch reliability. To date ~1700 Soyuz launches havebeen performed with a success rate of 96.8% [35]. Soyuz is used for manned and un-mannedmissions to the International Space Station and for commercial launches. The vehicles ded-icated to Resurs-DK1 is launched from Baikonur Cosmodrome in Kazakstan. The currentlaunch date is planned for 2006.

2.2 The PAMELA Apparatus

PAMELA is designed to measure energy of (and discriminate between) particles and anti-particles. Attaining the information necessary to do this for the wide energy range thatPAMELA intends to cover, requires an apparatus that makes use of a variety of particle de-tection techniques. The PAMELA detector is built in the form of a particle telescope. Thedifferent sub-detector systems include: a Time-of-Flight system, a magnetic spectrometerwith silicon tracking, an anticoincidence system, a silicon–tungsten electromagnetic calo-rimeter, a shower-tail catcher and a neutron detector. The positions of the different sub-de-tectors are shown in Fig.2–2.

The magnetic spectrometer allows for the sign of charge to be determined by means of de-flection in the magnetic field. It measures the absolute value of the rigidity, and since adown-going particle follows the same trajectory as a corresponding up-going antiparticle,

Figure 2–2: The PAMELA detector is built in the form of a particle telescope and consists of a Time-of-Flight system (ToF), a magnetic spectrometer with silicon tracking, an anticoincidence system (AC), a silicon–tungsten electromagnetic calorimeter, a shower-tail catcher (S4) and a neutron detector (adapt-ed from [74]).

AC (CAT)

Magnetic

Spectrometer

(behind CAS)

AC (CAS)

Neutron Detector

Calorimeter

ToF (S1)

ToF (S2)

S4

ToF (S3)

AC (CARD)

ANTIPROTON POSITRON

z

x

y

B

ToF (S1)

ToF (S2)

top AC (CAT)

side AC (CAS)

ToF (S3)

S4

Neutron

Detector

Calorimeter

Magnetic

Spectrometer

CARD


the Time-of-Flight is needed to decide the direction in order to determine the charge. A massmeasurement can be performed by the independent measurements of the rigidity and path-length, by the spectrometer, and the particle passage-time, by the Time-of-Flight system.The finite time resolution of the latter however puts an upper particle momentum limit to itsusefulness in mass determination. Above 1.5GeV/c separation between electron and protonmass is no longer feasible. However, the calorimeter is capable, beside the measurement ofthe energy release in its sensitive volume, to make a topological distinction between elec-tron and hadron induced signals. Thus for the purpose of separating positrons and antipro-tons from the main background of protons and electrons, the calorimeter will be used.Additional information useful in this separation is provided by the shower-tail catcher andthe neutron detector positioned below the calorimeter.

Originally PAMELA was designed with a Transition Radiation Detector (TRD) to help sep-arate electrons from hadrons. Due to production delays it is however not part of the PAME-LA instrument. In order to detect activity in the region originally intended for the TRD, ananticoincidence system (CARD) was instead placed in the space where the TRD was sup-posed to reside.

Below follows a more detailed descriptionof the various sub-detectors. The anticoin-cidence system is described in detail in thenext chapter.

2.2.1 The Time-of-Flight Sys-temThe Time-of-Flight (ToF) system [55]measures the velocity of the particlescrossing the apparatus and enables chargeidentification by dE/dx measurements inthe sensitive planes. The dynamic rangeallows identification of elements from Hto C.

In most of PAMELA’s different triggerconfigurations, the ToF system providesthe trigger electronics with signals forforming the main trigger signal that is dis-tributed to the various sub-detectors. ThePAMELA trigger is briefly discussed laterin this chapter.

The ToF system consists of three groupsof plastic scintillator counters (S1, S2 andS3) where each group is composed of two layers (see Fig.2–3). The first plane (S1) is placedon top of the instrument. Both S1-layers are 7mm thick. The second plane (S2) is locatedimmediately above the spectrometer. The S2-layers are 5mm thick. The last plane (S3) is

Figure 2–3: The ToF system.


placed between the spectrometer and the calorimeter, immediately below the magnet. BothS3-layers are 7mm thick.

The velocity of a particle can be calculated by combining ToF information with informationfrom the tracker about the particle’s trajectory length. The distance between S1 and S3 is78.3cm, which translated into time for 1.5GeV electrons equals to 2.6ns. With a timing res-olution (σ) of ~300ps the ToF system is able to discriminate electrons from protons up tothis energy at a level of 4σ. By identifying down-going particles system rejects up-goingparticles at a level better than 1 in 108. The dynamic range also enables charge identificationof elements from H to C by dE/dx measurements in the sensitive planes.

The ToF front-end electronic boards are divided in two sections where signals coming fromthe PMTs are converted either into time or charge. In the time section (used for timing meas-urement) a capacitor is linearly charged during the time interval that defines the passage ofthe particle. In the charge section (used for dE/dx measurement) a capacitor is charged withthe pulse itself. In both sections the capacitor is then linearly discharged for a time duringwhich a Time to Digital Converter (TDC) is gated.

2.2.2 The Magnetic SpectrometerThe magnetic spectrometer [2] is capable of determining the sign of charge and the rigidityof the traversing particles as well as performing dE/dx measurements in the sensitive planes.The spectrometer consists of a permanent magnet and a 6-plane silicon tracker positionedinside the magnetic cavity.

The Permanent MagnetThe permanent magnet is shown in Fig.2–4. It consists of five modules made of a Nd-B-Fealloy with a residual magnetization of ~1.3T, Each of these modules is composed of 16magnetized triangular blocks that are glued together to give the configuration shown in thefigure.

The cavity defines the PAMELA acceptance. It is 445mm tall with a cross-section of161×131mm2, which gives a geometrical factor of 20.5cm2sr. The magnetic field is orient-ed in the y direction and thus particles travelling along the z axis are deflected in the ±x di-rection depending on the sign of charge.

To enable precise rigidity measurements from the observed deflections, the magnetic field(B) has been carefully measured with a Hall-probe. In Fig.2–5a the By component measuredin the plane is shown as function of x and y. Fig.2–5b shows the By component meas-ured along the z axis. The mean magnetic field inside the cavity is 0.47T and has a maxi-mum value in the centre of 0.48T.

In order to minimize potential interference between the stray field outside the magnet cavityand the satellite instruments and navigation system, ferromagnetic shielding slabs have beenmounted around the magnet.

The Silicon TrackerThe silicon tracker is installed inside the magnet and allows the measurement of the momen-tum and the sign of charge of the crossing particle. The design goal for the rigidity is

z 0=


740GV/c. It is composed of 18 “ladders” of double-sided silicon strip detectors arranged on6 planes inside the permanent magnet. The detectors have a surface area of 70.0×53.3 mm2and a thickness of 300 µm; each side is divided in 1024 strips giving a total number of 36864electronics channels. In Fig.2–6a a tracker plane is shown.

Each silicon detector is segmented into micro-strips on both sides and consists of a high re-sistivity n-type silicon wafer with p+-type strips implanted on the junction side (x-view) andn+-type strips on the ohmic side (y-view). This is illustrated in Fig.2–6b. On the x-view sideof the silicon wafer the strip direction allows the position of the particle along the bendingdirection to be measured. The implantation pitch is 25µm and the read-out pitch is 50µm.In the y-view (non-bending magnetic field direction) the read out pitch is 67µm.

The efficiency for a single plane is of 90% for 1 minimum ionising particle (MIP). The sil-icon tracker is able to measure the absolute charge of particle up to Z=4. The dead-time in-troduced by the whole instrument is 1.1ms. The spatial resolution of the tracker is

Figure 2–4: The permanent magnet.

(a) (b)

Figure 2–5: a) By component measured in the plane as function of x and y b) By component meas-ured along the z axis

y

x

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(3.0±0.1)µm in the bending x view and (11.5±0.6)µm in y view. The momentum resolution(∆p/p) is an approximately linearly rising function from ~0.04 at 25GeV/c to ~0.25 at300GeV/c.

The front-end electronics (see Fig.2–6a) is based on a VLSI ASIC chip (VA1 [67]) whichcontains 128 charge sensitive preamplifiers connected to a shaper followed by a sample andhold circuit. It is divided into two sections (one per side), each containing 4 VA1 chips. Thesignals from the VA1s are then digitised in the ADC boards which are is connected to thefront-end by means of 50mm long 38–pin kapton cables. The digitised data are then sent viaserial links to the read-out boards based on a DSP where they are filtered by a compressionalgorithm.

2.2.3 The CalorimeterThe electromagnetic sampling calorimeter [16] is composed of 44 silicon sensor planes in-terleaved with 22 planes of 2.3mm thick tungsten absorber planes. As can be seen in Fig.2–7a, each absorber plane is sandwiched between two silicon planes. The full silicon sensorplane is composed of a square matrix of 3×3 detector wafers. These are large area devices(8×8cm2) and each of them is segmented into 32 strips with a pitch of 2.4mm and a thick-ness of 380µm. The strips are wire-bonded to the corresponding one of the other two detec-tors in the same row, thus forming 24cm long strips. To have a two-dimensional (x–y) read-out, the orientation of the strips of two consecutive silicon planes is rotated by 90°.

The calorimeter is modular with a basic module composed of two tungsten planes, eachsandwiched between two silicon detectors planes. The naming scheme adapted for the sen-sitive planes in a module is (from top to bottom) x-even, y-even, x-odd, y-odd. The 11 mod-ules slide independently into the common mechanical frame structure (see Fig.2–7b). Themodules are also electrically independent. The total thickness corresponds to 0.6λ (interac-

(a) (b)

Figure 2–6: a) Picture of a tracker plane.The plane is composed by three ladders inserted in an alumin-ium frame.The front-end electronics on the right is connected to the read-out boards by means of kapton cables b) Passage of a charged particle through the cross section of a silicon sensor. The ohmic and the junction sides are actually rotated by 90°.

X Side (Junction Side)

Y Side (Ohmic Side)

(Rotated 90 deg.)

double

metalization

p+ p+ p+

p+n+ n+

Al Al

25 nm

50 nm

SiO2 decoupling

blocking strip

SiO2

300 nm

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67 nmAl Al

n-type substrate

charged particle

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oard

fron

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tion lengths) and 16.3X0 (radiation lengths). The calorimeter has a mass of 110kg and a totalpower consumption of 75W.

The calorimeter has a constant energy resolution of 5.5% in the energy range of 20–200GeV.

Front-End ElectronicsThe core of the front-end electronics is a VLSI ASIC chip (CR1.4P [1]), designed for thePAMELA calorimeter. The main features of this chip are:

• a wide dynamic range of up to 1400MIP, where 1MIP ~5.1fC for a 380µm silicon detector

• large capacitance coupling capability (~180pF)• low noise and power consumption (6mW/channel).

One CR1.4P chip has 16 channels, each with a charge sensitive preamplifier, a shaping am-plifier/filter, a track-and-hold circuit and an output multiplexer. The characteristics for onechannel are; a peaking time of 2ms, a linear dynamic range between 0.4 and 1200MIPs, asensitivity of 5mV/MIP and a counting rate of 30kHz. In order to monitor the chip perform-ance a calibration circuit has been integrated on the device which can apply the required sig-nal to the selected channel by means of a 2pF injection capacitor.

Six 16-channel CR1.4P-chips are needed to acquire the 96 signals from one plane (seeFig.2–8). They are multiplexed onto a single 16-bit Analog to Digital Converter (ADC).The data coming from the 44 ADCs (one ADC per silicon plane) on the front-end boardsare then collected and analysed by the read-out boards before being transmitted to the mainDAQ system. The read-out boards consists of four independent sections, x-even, y-even, x-odd and y-odd.

Each read-out board consists of: a micro-controller dedicated to the slow control of the cal-orimeter, a DSP which controls the acquisition and elaborates the data stream, and a FPGAto parallelize the ADC data.

(a) (b)

Figure 2–7: a) schematic view of a calorimeter module. The module is composed by two tungsten slabs, each sandwiched between four planes of silicon detector. b) assembly of the calorimeter modules. The 11 modules are inserted in the main mechanical frame. They are fully independent both from mechanical and electronic point of view.

Si plane (x-view)

Tungsten Absorber

Si plane (y-view)

x 11 modules

1 m

odule


Particle IdentificationAs mentioned in Chapter 1 PAMELA needs to identify from a background of p that in-creases from 103 times the component at 1GeV/c to about 5⋅103 at 10GeV/c and froma background of that decreases from about 103 times the component at 1GeV/c to lessthan 102 above 10GeV/c. Thus, PAMELA has to separate (p) from ( ) at a level of105-106. The tracking system provides reliable information on the sign of charge and rigidityover a wide range of momenta from about 50MeV/c up to several hundreds GeV/c and theToF system can with high accuracy select down-going particles. Most of the remainingidentification is provided by the calorimeter.

Fig.2–9 shows a display for two events collected at a test-beam. The left event is a 50GeV/c electron and the right is a 100GeV/c proton. Electromagnetic and hadronic showers differin their spatial development and energy distribution in a way that can be distinguished bythe calorimeter. For example, for electromagnetic showers the shower maximum is con-tained inside the calorimeter for the energy range of interest. Another is that the lateral dis-tribution of the hadronic shower is much wider. Also, the development of theelectromagnetic cascade is strongly related to the energy of the primary electron and theelectromagnetic shower will with high probability start to develop in the first two to threeplanes of the calorimeter. For electrons there is a linear relationship between the primaryparticle energy and the deposited energy (up to ~1TeV for the PAMELA calorimeter).

Test-beam data and simulations have been used to estimate the rejection capabilities of thecalorimeter. Fig.2–10a shows the number of strips hit in the calorimeter as a function of to-tal detected energy for 200GeV/c test-beam electrons and protons. The separation can clear-ly be seen. Fig.2–10b shows the efficiency (top) and contamination (bottom) of theelectron/positron selection as a function of particle momentum. The full circles representtest-beam data and the open squares simulated data. During the analysis, one aim was tokeep the electron efficiency at about 90%. As can be seen in the figure, the agreement be-

Figure 2–8: A schematic view of the calorimeter front-end electronics.

CLOCK

DATA

CONV

CR1.4P

CAL1

CR1.4P

CAL2

CR1.4P

CAL3

ADC

16 bit

CIRCUITSHOUSE KEEPING

(currents, voltages,

-

+

CR1.4P

CAL6

CR1.4P

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CAL4

temperature)

MULTIPLEXER

ANALOG

BOARD

CONTROL

FPGA

SIGNALS

FROM DAQ

CONTROL

e+

e+ pe

−pe

−p e+


tween simulation and experimental data is acceptable. These results refer to a partiallyequipped calorimeter, hence they can be considered as lower limits of the PAMELA calo-rimeter performances.

Figure 2–9: Two test-beam events. The left display shows a 50GeV/c electron and the right is a 100GeV/c proton. The shower profile in the calorimeter is different.

(a) (b)

Figure 2–10: a) The total number of strips hit versus the total detected energy in the calorimeter for 200GeV/c e− and p data taken at the SPS test-beam. The good separation between the two particle spe-cies is clearly seen. b) Top: The efficiency for identifying electrons using the electron/positron calorim-eter selection as a function of particle momentum. The full circles and open squares represents test-beam and simulated data respectively. Bottom: The proton contamination of the electron/positron calorimeter selection as a function of particle momentum. The full circles and open squares represent test-beam and simulated data respectively, the full star is for 3GeV/c simulated protons, the arrows indicate the 68% confidence level (i.e. no events survived the selection) and the dashed arrow is for simulated protons at 10GeV/c.

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2.2.4 The S4 ScintillatorThe Bottom Scintillator (S4) is placed beneath the calorimeter and above the neutron detec-tor. The dimensions are 482×482mm2 and the thickness is 10mm. It is read out by six PMTssituated along the two opposite sides of the counter. The main purpose of this additionalscintillator is to act as a shower-tail catcher that will improve lepton/hadron discriminationby measuring the shower containment in the calorimeter. In addition, S4 provides the triggerfor the neutron detector (described below) in case of high energy interactions. When the sig-nal detected corresponds to more than 10MIP, it is sent to the main trigger and, if the coin-cidence occurs, the neutron detector is triggered to acquire data. The S4 detector can also beused together with the calorimeter self-trigger to detect very high energy (>100GeV) elec-trons coming from outside the experiment's main acceptance.

2.2.5 The Neutron DetectorThe neutron detector (ND) is located below the S4 scintillator and its purpose is to comple-ment the electron-proton discrimination capabilities of the calorimeter. It consists of 36 3Hecounters enclosed in a polyethylene moderator enveloped in a thin cadmium layer. Thecounters are stacked in two planes of 18 counters each, oriented along the y-axis of the in-strument. The size of the ND is 600x550x150mm3.

The ND can help in electron-proton discrimination by exploiting the much larger neutronproduction due to hadronic showers than to electromagnetic showers in the calorimeter. Theevaporated neutron yield in a hadronic shower is 10–20 times larger than that from electro-magnetic cascades. As the 3He counters are sensitive to thermalized neutrons, evaporatedneutrons produced in the calorimeter are first thermalized, then detected in the neutron de-tector. The detection efficiency, taking into account the thermalization efficiency, is about10%. A combined analysis using the ND and calorimeter will allow primary electrons to beidentified at energies up to a fewTeV.

2.2.6 The Trigger

Trigger ConditionsThe dimensions of the ToF layers have been imposed by the projected trajectories of parti-cles passing “cleanly” through the tracking detector in such a way that the solid angle seenby a particle hitting all the three ToF planes is equal to the tracker geometrical acceptance(20.5cm2sr). The trigger system provides a fast signal, formed when coincident signalsabove a certain threshold occurs on the ToF outputs. Several configurations of the PAME-LA trigger are foreseen:

Main Trigger. It is determined by coincident signals in S1, S2 and S3. There are two sub-configurations (Subscripts 1 and 2 refers to the first and second layer of a ToF plane):

• (S11 or S12) and (S21 or S22) and (S31 or S32)• (S11 and S12) and (S21 and S22) and (S31 and S32)

Where the signal in each layer is given by the logical “or” of all the strips. The first config-uration is more redundant, while the second one has a higher efficiency in background re-jection.


Radiation Belts Trigger. During passage over the radiation belts and the SAA, S1 is ex-pected to saturate. It is therefore excluded from the trigger algorithm in this trigger config-uration

Calorimeter Self Trigger. The calorimeter can be operated in self-trigger mode, which al-lows the stand-alone detection of e± with an increased acceptance of 600cm2sr up to an en-ergy of ~2TeV. Since these events are quite rare, it is necessary to have a larger geometricalacceptance. The calorimeter triggers itself if there is an energy release corresponding to atleast 150MIP in any of the four sections.

S4 Trigger. When the bottom scintillator S4 detects a signal above the 10MIP threshold, ittriggers the neutron detector.

Trigger RateThe estimated trigger rate in the main configuration is shown in Fig.2–11. The contributionfrom the radiation belts are not included in this figure. During one orbit the trigger rate rang-es between a fraction of a Hz at the equator and a few kHz in the SAA. By “true” triggers ismeant triggers where a single particle has passed cleanly through the acceptance.

The foreseen average trigger rate is ~10Hz (see Chapter6), and the average event size ~6kBof compressed data. This results in a daily data volume of ~6GB. In case the amount of datashould exceed either the storage dedicated to PAMELA on the Resurs spacecraft or the dailylimit of 20GB of data down-linked to ground1, an on-line event selection can be appliedthrough a second level trigger, to be activated by an up-link command from ground. Thesecond level trigger is the subject of a study presented in Chapter 6.

Figure 2–11: The PAMELA trigger rate evaluated over an orbital period. The dotted curve shows an estimation of the total trigger rate while the solid curve shows the rate of only “true” triggers (i.e. trig-gers caused by a single charged particle cleanly traversing the acceptance of the experiment). The large peak to the right is due to the SAA [76].

1. This is not a technical limitation of the experiment, that can continuously acquire data with a trigger rate of ~60Hz.

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2.2.7 The DAQ System

PAMELA’s DAQ system is illustrated in Fig.2–12. The PSCU (PAMELA Storage and Con-trol Unit) includes a CPU based on a ERC–32 architecture (a SPARC v7 implementation)running the RTEMS real time operating system. It is custom-designed and space qualified.

The PSCU communicates with the satellite via a standard 1553B data bus, and handles allslow control, interaction with the satellite, data acquisition, storage and down-link. ThePAMELA InterFace board (PIF) transfers data from the detectors to mass memory. Two re-dundant mass memory modules of 2GB each support latch-up detection. When a latch-upoccurs, the operation is transparently switched to the safe module. A multipurpose Teleme-try and Control Board (TMTC) contains several interfaces to interact with the subsystemsof the experiment.

A so-called Intermediate DAQ System (IDAQ) has been included to perform acquisitionfrom the sub-detectors to the PSCU at a rate of 2MB/s. When receiving a trigger the PSCU(via the IDAQ), activates the procedure to read out the data from the sub-detectors sequen-tially. The IDAQ system is also responsible for main the second level trigger decision basedon data from the AC system and the calorimeter. The data is stored in the PSCU mass mem-ory. A few times per day, data are transferred to the satellite on-board memory (via a 10MB/s bus) where it is stored before being down-linked to Earth.

During each orbit, the satellite will send data to two down-link stations in Russia where ata transmission rate of 150Mbit/s. Two daily down-link windows of about 10 minutes eachwill be dedicated to the transfer of PAMELA data with a maximum limit of 20GB/day.

Figure 2–12: Scheme of the PAMELA data acquisition system. The interfaces (IF) between the IDAQ and the PSCU handle the data acquisition and transfer the commands (CMD) to the IDAQ. The commu-nication between PAMELA and the spacecraft is handled by the 1553B bus and by adapter modules for telemetry (TAM) and for data transfer (VRL) [76].

Neutron Detector

Tracker

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2.3 Current StatusThe PAMELA apparatus is now (november 2005) located at the TsSKB [87] factory in Sa-mara, Russia. Prior to the delivery, it was assembled at the INFN1 laboratories in Rome, It-aly, where the system was tested with ground muons for a period of several months.

The muon charge ratio (presented in Fig.2–13 a) was measured. Non interacting particles inthe calorimeter with charge one in the ToF scintillators were selected as muons. Low-energyprotons were rejected based on the ionization losses in the calorimeter. Their momenta weredetermined by the tracking system. In the figure a good agreement can be seen between thePAMELA data and other experimental results [32,51,15]

Fig.2–13b shows the velocity of the particles (in units of speed of light, β) measured by theToF system, as a function of their rigidity. Most events are relativistic muons. A small pro-ton component is visible at low momenta. The solid line indicates the theoretical β for pro-tons. As seen, the measured ToF resolution of 300ps will allow positrons to be separatedfrom protons and antiprotons from electrons up to about 1GeV/c.

Fig.2–14 shows two event displays. The first is a 2.8GeV/c negatively charged particle(top). The non-interacting pattern in the calorimeter indicates that the particle with highprobability is a µ−. The second is a 6.6GeV/c particle with an hadronic interaction in thecalorimeter, consistent with a proton signal. The solid lines in the event display indicate thetracks reconstructed by the fitting procedure [70] of the tracking system.

Fig.2–15 shows a photo of PAMELA taken in Russia prior to integration with the satellite.

1. Istituto Nazionale di Fisica Nucleare

Figure 2–13: a) The muon charge ratio at ground measured by PAMELA compared with the 2002 global fit of experimental data and more recent experimental results. The dashed lines indicate one standard deviation of the fit b) The particle velocity measured by the ToF system as a function of rigidity. The solid line is the theoretical β for protons.(both figures from [13])

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2.3 – Current Status 35

Figure 2–14: (top) The event display of a 2.8GeV/c µ− from ground data in the PAMELA appara-tus.(bottom) The event display of a 6.8GeV/c. hadron from ground data in the PAMELA apparatus.

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Figure 2–15: PAMELA just before integration with the satellite.

Chapter 3

The Anticoincidence System

The PAMELA experiment is equipped with an anticoincidence (AC) shield to help vetobackground caused by out-of-acceptance events. The primary sub-detector for measuringthe energy spectra of charged particles is the spectrometer. The AC system surrounds thespectrometer magnet and is designed to help reject particles which do not pass cleanlythrough the tracking system but nevertheless generates a trigger signal. The AC system alsoconstitutes a part of PAMELA’s second level trigger. In this chapter the AC system is de-scribed, and results of qualification tests for vibration and temperature tolerance are present-ed.

The CARD system is a smaller AC shield thatwas added to the experiment to cover the spacebetween the S1 and S2 ToF planes, which wasadditionally intended for the transition radiationdetector (see Chapter2). It consists of four smallscintillator planes similar to those in the mainlateral AC system. The electronics used to readout the CARD detectors is identical to that ofthe main system. This thesis concerns only thedetectors of the main AC system. The workdone on the construction and verification of theelectronics is however valid for both systems.

3.1 PurposeAs described in more detail in the last chapter,the main first level trigger condition in PAME-LA is the coincident of energy deposition in thethree ToF planes S1, S2 and S3. The majority ofthese triggers (~75%) are however expected tobe false, which usually means that they havebeen caused by interactions in the mechanicalstructure of the experiment itself or in the satel-lite. The former scenario is illustrated in Fig.3–1, where two different ways in which a par-ticle can interact in the magnet to produce a trigger signal, is shown. It is expected, and hasbeen shown, that this kind of false triggers are likely to be associated with activity in the ACdetectors [42].

Figure 3–1: Two ways in which a false trigger signal can occur. In both cases there is a high probability of activity in the AC detectors.Only detectors above the calorimeter are shown in the figure.

incoming

particle

incoming

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CAT

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38 Chapter 3 – The Anticoincidence System

These false background events can be substantially reduced during the off-line analysis byusing information from all the various sub-detectors, with the AC system providing addi-tional redundancy. The reduction can also be implemented on-line in the form of a secondlevel trigger (Chapter6 is dedicated to this subject).

3.2 System OverviewThe system consists of five anticounters - four lateral scintillators (CAS) and one top scin-tillator (CAT). It is placed around the spectrometer, covering the sides and the top of themagnet in the geometrical configuration shown in Fig.3–2. Each anticounter consists of aplastic scintillator sheet with photo multiplier tube (PMT) read-out.

For both CAS and CAT the PMTs are coupled to the scintillators with 7mm thick silicone“cookies”. The scintillators are then mounted into aluminium containers designed to protectthem from mechanical stresses during launch and light-leakage. These boxes also act asPMT holders to ensure a good scintillator–PMT coupling and provide means to attach thecounters to the PAMELA main structure. A high voltage divider is located directly behindeach PMT. The assembly of a CAS plane is shown in Fig.3–3 and of the CAT plane inFig.3–4.

As seen in the Fig.3–3, each CAS detector is read out by two identical PMTs. This is in or-der to decrease the chance of a single point failure. Also for redundancy, but mostly to coverthe full sensitive area (the shaded area in the figure), the CAT detector is read out by eightPMTs, two for each quadrant. The system has 16 channels in total.

A miniature low intensity 640nm light emitting diode (LED) is glued directly onto the lowerpart of the side anticounter scintillator and two LEDs are glued to opposite sides of the topanticounter scintillator. The LEDs allow the verification of functionality and stability of theanticounters during flight. In order to

the anticoincidence system of the pamela satellite ... -...

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