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Production mechanisms

High energy gamma photons are produced, by charged cosmic rays, according to two main mechanisms: leptonic and hadronic[14][15].
• In the leptonic processes, gamma photons are emitted through one of three scenarios[12][16].
– Synchrotron radiation of relativistic electrons gyrating in a magnetic field.
– The inverse Compton scattering of low energy photons by relativistic electrons.
– Bremsstrahlung radiation emitted by accelerated electrons in the electrostatic field of nucleus.
• The hadronic process is based on the decay of neutral pions π0. These neutral pions are produced through nuclear interactions between high energy protons or heavier nuclei with ambient protons, nucleons or photons in dense regions of interstellar medium.
Other mechanisms of gamma photons production are the annihilation of dark matter particles (Weakly Interacting Massive Particles WIMP)[16][17] and the primordial black hole evaporation.

Detection mechanisms

It is possible to directly detect gamma rays in space but at high energies above 1010 eV , large detection area is required. This is because of the low flux estimated by some events per year per m2[16]. Besides, at these energies it is impossible to focus these photons of wavelengths shorter than interatomic distances. Hence, wide detection area is needed and this can not be afforded by the relatively small space-based detectors (e.g, FERMI-LAT ∽ 1m2 and AMS). For the very high energy gamma rays, the solution is then to indirectly measure them on the ground.
Considering the opacity of the atmosphere, when high energy gamma rays enter the atmosphere, they interact with its nuclei and produce cascade of secondary particles (electrons, positrons, photons) called extended air shower (EAS)3. The number of these secondary particles increases while their energy decreases. It depends on the nature and energy of the primary particle whether this air shower can reach the ground or die out earlier at high altitude.
Ground based detectors measure gamma-ray induced particle-cascades in the atmosphere by detecting those particles that reach the ground (air shower detectors), by detecting the fluorescence photons emitted from the shower during its development through the 3So do hadrons and leptons as well. In contrast to electromagnetic air shower, hadronic cascade is generally irregular and grows asymmetrically around the direction of the incident particle[18]. atmosphere4 (fluorescence detector in the Pierre Auger observatory) or by their Cherenkov light (Cherenkov telescopes). Fig.1.5 shows the indirect detection of V HE − γ − ray on ground according to these two scenarios.
Figure 1.5: Indirect γ − ray detection scenarios. Right: air shower particle detector (detects cascade particles that can reach the ground). Left: Cherenkov telescope (images Cherenkov light emitted by secondary gamma-ray induced particles that do not reach the ground). Here are some examples of gamma rays detectors in space. OSO-3 (third Orbitingg Solar Observatory), SAS-2 (the second Small Astronomy Satellite), COS-B, EGRET (Energetic Gamma-Ray Experiment Telescope), AGILE (Astro-rivelatore Gamma Immagini LEggero) and Fermi-LAT (Fermi Large Area telescope) are telescopes that were (are still for some of them) landed on satellites.
CANGAROO (Collaboration of Australia and Nippon (Japan) for a GAmma Ray Observatory in the Outback), HESS (High Energy Stereoscopic System), MAGIC (Major Atmospheric Gamma Imaging Cherenkov Telescopes), VERITAS (Very Energetic Radiation Imaging Telescope Array System), and HAWC (High-Altitude Water Cherenkov Observatory) are examples of on-ground very high energy gamma rays (30GeV − 100 TeV ).
4Atmospheric nitrogen molecules being excited by shower particles (electrons and positrons) emit isotropically fluorescence photons with wavelengths between 300−400nm[19][20]. The fluorescence light emitted is proportional to the energy deposited in the atmosphere by charged particles in the shower.
Light observed on the ground, after some corrections related to the varying attenuation of the atmosphere and to the signal contamination by Cherenkov light, can give a measurement of the energy of the primary particle[21].

Cherenkov light

When secondary charged particle in gamma induced cascade move in the atmosphere with a speed greater than the speed of light (v > c/n), it will asymmetrically polarize the medium in front of and behind it causing an electromagnetic shock wave. Consequently, it will emit light near the ultra-violet5 called Cherenkov light. Cherenkov light is emitted in a cone around the direction of the original particle. It makes on the ground a circle of about 250m diameter (for near vertical showers) called Cherenkov pool[22].
Many on-ground very high energy gamma ray telescopes were built to detect photons by the Cherenkov light. Some examples are HESS, MAGIC, and VERITAS. Fig.1.6 shows the energy range covered by some of these telescopes besides that aimed to be covered by the future CTA.
Figure 1.6: Integral sensitivity for CTA from MC simulations, together with the sensitivities in comparable conditions (50 h for IACTs, 1 year for Fermi-LAT and HAWC) for some gamma-ray observatories[3].

Cherenkov Telescope Array (CTA)

CTA is a future gamma-ray observatory that aims to achieve full-sky coverage with considerable improvements over the existing instruments (Fig.1.6). Two sites in the two hemispheres are programmed for this purpose. The northern site will be situated in La Palma in Spain while the southern site will be in Paranal in Chile6.
Each site will be equipped by an array of telescopes of different sizes: Large, medium, and small size telescopes. That is how a wider energy coverage is possible between some tens of GeV (20GeV ) up to some hundreds of TeV . Fig.1.7 illustrates the energy rang of each telescope of the CTA.
The northern site will be dedicated to extragalactic physics (will not require coverage of the highest energies). While the southern site, will be optimized for galactic and extragalactic physics (hence it needs to be sensitive over the full energy dynamic range)[2].
Analyzing signals from an array of telescopes provides a stereoscopic view at large energies. This improves the angular and the energy resolutions (improve the quality of the data) besides allowing a better background rejection (better gamma-hadron separation)[2].
Thus CTA project will deliver better sensitivity thanks to the wider effective area covered by the array (an order of magnitude improvement with respect to the current minimum energy 30GeV detected by HESS), wider energy coverage (span about 4 decades between 20GeV and 300 TeV ), better energy and angular resolution, besides, a wider field of view (6 − 8 degrees).
The technical improvements that are needed to realize such performance are being studied in the institutes member of the CTA collaboration all over the world. This includes, among many others, the choice of the photodetecters. Two main candidates are photomultiplier tubes (PMT) and the silicon photomultipliers (SiPM). SiPM-based camera will equip the dual mirror Schwarzschild-couder (SCT) MSTs and most of the SSTs. The use of SiPM in the camera of the LST is under study.
Large Size Telescope (LST): It can also be called the low energy instrument in whose context lays the question of this thesis. It will be discussed in the next section.
Medium Size Telescope (MST): It is a 12m diameter telescope of one mirror Davis-Cotton design. Another design is also under study. That is Schwarzschild-couder dual mirror design (SCT) of smaller diameter that allows for a wider FoV and whose camera is based on SiPM detectors[23].
MST array covers an energy range from 100GeV up to 10 TeV . These energies are well covered by other current instruments. The particularity of the CTA is the use of larger number of MSTs which will improve the sensitivity due to the increased covered area and the better quality of shower construction. An additional advantage of extended MSTgrid (operating in two telescope trigger condition) is the possibility of lower triggering threshold since there are always telescopes sufficiently close to the shower core[22].
Small Size Telescope (SST): This is the smallest in size telescope (4−6m diameter) with a field of view of around 10 degrees. It covers the upper part of energy in the range covered by the CTA (10 TeV − 300 TeV ). Since at these high energies light yield is large, large number of this telescope will be used to cover a wide area of the southern CTA site[22]. The northern CTA site will not host any SST because it will focus on the extragalactic sources, hence it will be dedicated to lower energy range.
Several technical solutions are proposed for the SSTs: a single-mirror Davis-Cotton telescope (SST-1M) and two telescope designs with dual-mirror Schwarzschild-Couder (SST-2M). All the three designs will use SiPM-based cameras[24].
Figure 1.7: Dynamic range of CTA telescopes. Thin lines with small symbols illustrate the limited impact of a reduced dynamic range of the readout electronics (clipped at 1000 photoelectrons). The dashed black line with diamonds, shows the sensitivity if there was no electron background[3].


Large Size Telescope (LST)

It will be dedicated to the lower range of energy (down to few tens of GeV ) covered by the CTA. The advantage of detecting in this energy domain by a ground-based instrument is to provide a wider effective collection area and field of view, than that in the space-based telescopes like Fermi-LAT whose effective area is < 1m2.
A group of 4 large size telescopes will be installed in each CTA field. These telescopes will be placed in a compact configuration with inter-telescope distances of about 100m.
Each telescope is a 23m diameter and provides a field of view of 4.5 degrees.
The main goal of this LST grid is to achieve sensitivity in the energy domain from few tens of GeV (10 − 20GeV ) up to ∼ 100 − 300GeV 7. Even if event rate is high, the uncertainty in the systematic background seems to limit sensitivity improvement (i.e, higher detection rate but lower amount of detectable light). LST allows for a higher collection efficiency and multiple LSTs allows better background suppression [2].
Detecting photons in Cherenkov light with as low as possible energy (down to 10GeV ) requires efficient photosensors. The key feature for a better sensitivity is the photo detection efficiency (PDE: Eq.A.4) which is strongly related to the quantum efficiency (QE: Eq.A.3) of the photodetector devices in the camera[25]. For the first generation camera of the LST, photomultiplier tubes (PMTs) are planned to be adopted[26][27]. Advantages and disadvantages of such detectors will be discussed in the next chapter (Chapter.2).
An emerging technology in photo detection is represented by the solid state silicon photomultiplier (SiPM). These new detectors promise to provide a rival performance in respect to that of the PMT, particularly in their quantum efficiency. SiPM structure and performance will also be presented in the next chapter(Chapter.2).

First generation camera of the LST

LST reflector is a parabolic dish of 23m diameter that consists of 198 segmented mirrors.
Each mirror is a hexagonal-shape of the size 1.5m flat to flat which makes a unit of ∽ 2m2. The total area of the LST reflector is then ∽ 400m2[27].
The camera is situated in the focal plane of the mirror and is conceptually divided into three parts: focal plane instrumentation, cluster electronics and the global camera elements. These three parts will be put inside a sealed structure with temperature control[26]. One can understand the motivation for such a structure since the camera of the first generation LST is planned to use PMTs. A motorized shutter will be used to protect the PMTs and to allow for a day-light operation with high voltage switched on.
Electronics inside the camera dissipate power (∼ 5 kW) which produce a considerable amount of heat that needs to be evacuated. In order to stabilize temperature inside the camera body, two scenarios are being studied: Air flow cooling system and water cooling system based on cold plates[26].
The first generation of the LST camera will be based on Hamamatsu R11920-100 PMT photodetectors (Table.1.1 lists the main characteristics of these PMTs). 1855 PMTs will be organized in 265 clusters of 7 modules each. Fig.1.8 shows a group of three clusters, while Fig.1.9 lists the different components of one cluster.

Table of contents :

1 Introduction to high energy gamma-rays astronomy 
1.1 Introduction
1.2 Cosmic ray
1.2.1 Energy spectrum
1.3 Gamma rays
1.3.1 Production mechanisms
1.3.2 Detection mechanisms
1.3.3 Cherenkov light
1.4 Cherenkov Telescope Array (CTA)
1.5 Large Size Telescope (LST)
1.5.1 First generation camera of the LST
1.5.2 Second generation LST camera
1.6 Conclusion
2 SiPM based camera 
2.1 Introduction
2.2 Photo detection principle
2.3 Photomultiplier tube (PMT)
2.4 Semiconductor photodetectors
2.4.1 PN junction
2.4.2 PIN junction
2.4.3 Avalanche Photodiode (APD)
2.4.4 Geiger-Mode Avalanche Photodiode (G-APD)
2.5 Silicon Photo-Multiplier (SiPM)
2.5.1 SiPM industrial variety
2.5.2 Multi-pixel Silicon photo multiplier (SiPM matrix)
2.5.3 SiPM in astrophysics instruments
2.6 4 × 4 SiPM matrix characterization
2.6.1 Test bench
2.6.2 Temperature Measurement
2.6.3 Pedestals
2.6.4 Gain variation with over voltage
2.6.5 Optimum operating point
2.6.6 Gain dispersion among the 16 pixels
2.6.7 Noise in SiPM
2.7 SiPM Modeling
2.8 Conclusion
3 ALPS Design 
3.1 Introduction
3.2 State of the art in readout-electronics
3.3 Readout chip, block diagram development
3.4 The preamplifier
3.4.1 Interface with the detector – Input stage Transimpedance amplifier Current amplifier (current conveyor)
3.4.2 Gain stage
3.4.3 Current mirror choice Simple current mirror Cascode current mirror
3.5 Proposed preamplifier and simulation results
3.5.1 Symmetric currents (transimpedance gain)
3.5.2 Symmetric loads (current gain)
3.5.3 Combination of transimpedance and current adjustment
3.5.4 Response speed
3.5.5 Dynamic range, Linearity
3.5.6 Signal to noise ratio
3.5.7 Power consumption
3.6 Slow control
3.6.1 8-bit DAC for per-pixel over-voltage control
3.6.2 Preamplifier gain correction
3.6.3 Noisy channel elimination function
3.7 Analog sum function
3.8 Trigger
3.9 Overall chip and layout
3.10 Single-ended to differential converter
3.11 Conclusion
4 ALPS experimental results and measurements 
4.1 Introduction
4.2 Test board design
4.2.1 Labview slow control interface
4.3 Test bench of Electrical tests
4.4 Preamplifier tests
4.4.1 Linearity and dynamic range
4.4.2 Signal shape
4.5 Slow control tests
4.5.1 Residual pedestal
4.5.2 Gain control
4.5.3 Gain equalization
4.5.4 Noisy channel suppression
4.5.5 Bypassing gain control
4.6 Sum function
4.6.1 The shape of output signal
4.6.2 Sum as a function of the controlled gain
4.6.3 Sum of different signals
4.6.4 Sum as a function of the input level
4.7 Test bench for ALPS test with SiPM
4.8 Signal shape
4.9 Conclusion
A Terminology of photodetector devices
B Electrical tests (continuation)
B.1 Residual pedestal
B.2 Gain control
B.3 Gain equalization
B.4 Sum as a function of the input level
C ALPS tests with SiPM (Specification)
C.1 Light emitting diode LED
C.2 Optical filters
C.3 Preamplifier test
C.4 Sum Function


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