Abuso Sexual

The optical signals are received and converted back to electrical ones again by fast GaAs PIN diodes on the receiver board. The electronic signals are split into two. One part of the signal is sent to the trigger branch and the other part is sent to the FADC system, see Fig. 4.2. There are discriminators with high gain, low gain amplifiers on top of the board. After the MUX-FADC was installed in February of 2007, the high and low gain splitter is not necessary anymore.

4.8.2

Trigger

The MAGIC trigger decision is defined by three different levels. In current operation, only the first two are used.

• Level-0: Very basic discriminating of each pixel. After extraction the signals and charge integration from FADCs, the total amount of charges is digitized and converted into FADC counts. For each pixel, a constant fraction discriminator (CFD) will select only the pulses which pass through the trigger level. The threshold can be easily adjusted by software for different NSB levels, or FOVs, or trigger algorithms (such as SUM Trigger). The discriminator level is decided in order to keep individual pixel rates between 100 and 400 kHz. For MAGIC, the level-0 trigger is restricted to the inner camera only. The total angular span is about 1 degree.

• Level-1: To search for next neighboring coincidence. We know that electromagnetic show- ers are more compact than hadronic ones. Moreover, the hadronic showers are more messy. Close-neighboring coincidence is used for discriminating the hadronic and EM showers.

Figure 4.7: Sketch of the trigger topology in the MAGIC camera.

The camera is divided into 19 overlapping cells and with 36 pixels each. Fig. 4.7 shows the level-1 trigger topology. Level-1 gives a trigger signal if for at least one of the cells a logical condition is fulfilled.

• Level-2: Higher level trigger depending on the topology of the trigger pattern.

An individual pixel rate monitoring system is used to prevent high trigger rates which might be caused by the bright stars in the pixels. The whole system is adjusted such that the data are typically taken at a rate of 250 Hz.

4.9

FADC/MUX-FADC

Once the trigger has arrived, the FADC system will continuously digitize and write 30 slices into the buffer. The dead time is about 2-3%. Under normal conditions, events are recorded at a rate of∼250 Hz. Nevertheless, the maximum tolerable rate had been about 600 Hz before the MUX-FADC was installed in February 2007. The night sky background level at the MAGIC site is about 130 MHz. MAGIC was using 8-bit, 300 MHz FADCs for each pixel. In order to mea- sure the Cherenkov pulse structure, the typical 2.5 ns FWHM signals are artificially stretched to 6 ns. The stretched pulse shape is unfortunately smears out the differences among pulse shapes induced byγ-ray, cosmic rays (mainly proton) and NSB photons. Furthermore, to provide a dy- namical range of∼1000, a two gain charge extraction is implemented. The signal is amplified by a factor of 10 and then digitized. If the high-gain signal exceeds 250 FADC counts, a GaAs switch is activated and the delayed (∼55 ns) low gain will be digitized. If the high-gain signal does not exceed that threshold, the recorded information is derived from the high-gain branch.

4.9 FADC/MUX-FADC 57

Figure 4.8: Performance comparison between old FADC and new MUX-FADC. The left plot shows the pulse sampling and the right plot depicts the pixel distribution of the integrated noise. Therefore the recorded signal is either full of 30 slices of high-gain or 15 high-gain slices fol- lowed by 15 low-gain slices. If the FADC does not record the low-gain signal, the 15 late slices of the high-gain signal should not contain any information from a signal and thus could be used as an estimation of the pedestals.

Figure 4.9: The different components of the calibration system used in the MAGIC telescope. After February 2007, a Fiber-Optic signal multiplex system which uses a 10-bit, 2GSamples/s FADCs to digitize 16 read-out channels consecutively has been installed [96]. The analog signals from each of the 16 pixels are delayed by using optical fibers and digitized sequentially with the single channel of Ultra fast FADCs. This multiplexed (MUX) FADCs read-out reduces the costs by about 85%compared to using FADCs for each readout channel. A comparison of the pulse sampling between old FADCs and the new MUX-FADCs is shown in Fig. 4.8. The total integrated time window is about 40 ns. If we dispose of the first and the last 5 ns switching noise, the effective time window is about 30 ns. The integrated noise is lower for MUX-FADCs than

for the old FADCs, see Fig. 4.8. A DAQ rate up to 100 MBytes/s corresponding to a trigger rate of 1 kHz has been achieved. The upgrade of the FADCs not only improves the rejection of the NSB photons, but also the γ/hadron separation. In Fig. 4.6, γ-ray, proton induced showers and muons have different distributions of the RMS value of the arrival times. The timing information is expected to improve the separation ofγ-ray events from background events. Including a new analysis techniques using timing information from the air showers, the sensitivity moved up by 40 %.

4.10

Calibration System

The whole electronic chain of the MAGIC Telescope is calibrated. The main purpose is to obtain the conversion factors and absolute signal timing information. An optical calibration system has been installed [97]. It is composed of several differently colored ultra-fast LEDs with three wavelengths, namely 370 nm, 460 nm and 520 nm. These LEDs can illuminate the camera with different color of light pulses at different intensities. (from 5 to 700 ph.e. for each inner pixel). The whole system is illustrated in Fig. 4.9. By using this calibration system, the whole dynamical range of the camera is accessible. The pulsed light is used to achieve a relative calibration of the data chain using the F-factor method (will be described in the next chapter). The PIN diode and three blinded pixels are used for absolute calibration of the camera. The calibration events are recorded in dedicated calibration runs and interleaved calibration events during the regular data taking procedures. As a result, the gain variation of the whole data chain (e.g. from PMTs or VCSELs) can be monitored during the data taking.

4.11

Observational Mode

During observation, the data-taking sequence is performed in two different modes. In gen- eral, an observation sequence consists of a pedestal run, a calibration run and a number of data runs. The pedestal runs contain 1000 randomly triggered events. These events contain the signal baseline and its fluctuations are given by NSBs as well as noise in the readout chain. The cal- ibration runs contain events triggered by the calibration system while synchronously uniformly illuminating the PMT camera by a flash LED or DC light source. These events are used to calcu- late an initial set of the calibration constants. Data runs contain events that triggered the telescope while tracking a particular sky position. Those runs contain interleaved calibration events at a rate of 50 Hz, which are used for a continuous re-calibration. Part of the events which do not contain any shower image are used to update the pedestal information. In addition to the data run, a number of subsystem monitoring files are produced by the telescope subsystems, for instance, the camera control system, the star field monitor and the central control system. These useful additional system information will be merged with the event data offline. At a certain point in the analysis (which will be described in the next chapter), they are then synchronized to each individual event.

4.11 Observational Mode 59

Figure 4.10: Illustration figure for different observational modes. In the ON-mode, the source is positioned in the center of the camera (red circle). For the WOBBLE observation, the source is placed at0.4◦ off center (marked as ”signal region” in the plot.) The background region in the WOBBLE-mode can be found at the opposite side of the signal region or even any other regions on the 0.4◦ circle. This provides a simultaneous measurement of the background. On average, every 20 minutes, the wobble positions are swapped. More background samples can be taken if the observation time is long enough. This increases the statistics of the background data. The plot is adapted from [10].

• ON/OFF mode : In this observation mode, ON and OFF observations are complemented by equal amounts of time. During the ON mode, the observed source position is tracked in such a way that the center of the camera is pointing to the source. While in the OFF mode observation, the telescope tracks a position in the sky where noγ-ray source is expected to be. In the ideal case, the OFF region of the sky should have the same background light and atmospheric conditions as the corresponding ON position. The swapping time is about 15-20 min between ON and OFF observations. The time should not be too long because the atmospheric conditions may change. Since the source is located in the center, thus this observation mode is more sensitive and suitable for measuring weak sources. The main problem is that the background and signal region may mismatch because they are taken at different times and the atmospheric conditions change.

• Wobble mode : In the Wobble observation mode, simultaneous recording of the ON and OFF data at different places in the camera is possible. The source position is located at a fixed distance of the camera center. The position which is at the opposite side of the source position is called ”anti-source position”. It can be used to determine the background. Sometimes, in order to get better statistics, two more background regions are defined at the same distance from the camera center, but ±90◦ displaced from the source position. In order to reduce the systematics because of the chosen source position in the camera (in case if the camera is not homogeneous), the source and the background positions are regularly swapped in every wobble observation. Since the source is located at offset of the camera center, the acceptance of the trigger area is reduced for higher energy showers. If the data was taken only with one wobble position, the background determination has additional uncertainties because the camera inhomogeneity may destroy the background and data sample. Thus the swapping of the two positions helps to reduce the systematic effects. Moreover, the inhomogeneity of the camera will be smeared out due to frequent swapping between different wobble positions.

4.12

MAGIC II

The MAGIC Collaboration has built a second telescope at 85m distance from MAGIC I on the the island of La Palma. With advanced photon detectors and readout electronics, MAGIC- II, the two telescope systems, will have a reduced analysis threshold. The overall sensitivity in stereoscopic/coincidence operation mode is expected to increase by a factor of 2-3. The expected sensitivity curve is shown in Fig. 4.11. In order to reduce the energy threshold, the overall light collection efficiency for the Cherenkov photons has to increase. Therefore the camera of the second telescope will be equipped with optimized Winston cones and photon detectors with the highest possible quantum efficiency (QE) up to33%. The testing characteristics of these photon sensors will be discussed in the following chapter.

Chapter 5

Photodetector Studies for the MAGIC II

Telescope

5.1

Introduction

The exploration of the high energy cosmic radiation is realized through increasingly complex experiments covering a wide energy range from underground detectors, ground-based detectors, satellites and space observatories. The photodetectors are a basic and crucial part for nearly all the experiments in astroparticle physics. Especially, observing high energy cosmic radia- tions, such as gamma rays, cosmic rays and neutrinos all use the techniques based on measuring Cherenkov or fluorescence light induced by particle showers in different materials. Efficient photodetectors are required and needed.

The sensitivity of the IACTs relies on the light detection efficiency of the photodetectors. Basically, the higher the detection efficiency of the photodetectors, the better the sensitivity and the lower the energy threshold of the telescope. The Cherenkov light intensity is weak, thus it is necessary to work with photodetectors in photoelectron counting mode. In addition, Cherenkov light pulses are short down to a few ns, and the response time of the photodetector is crucial. For all the above reasons, photomultipliers (PMTs) are so far the most reasonable, best suited and most stable devices. They have been used in high energy physics experiments for quite a long time and the technologies of production and development are mature.

For further lowering the energy threshold and improving the sensitivity of the system, the new telescope was designed and built. It is using higher sensitivity PMTs. Various measurements of different characteristics of these PMTs, such as a single photoelectron spectrum (SPE), afterpulse rate, timing characteristics, aging and photoelectron detections were performed. All the results from the measurements will be presented in this Chapter.