source HESS J1745-290 in the Galactic Center
Dissertation
zur Erlangung des akademischen Grades doctor rerum naturalium
(Dr. rer. nat.) im Fach Physik
Spezialisierung: Experimentalphysik eingereicht an der
Mathematisch-Naturwissenschaftlichen Fakultät der Humboldt-Universität zu Berlin
von Philipp Wagner
Präsidentin der Humboldt-Universität zu Berlin:
Prof. Dr.-Ing. Dr. Sabine Kunst
Dekan der Mathematisch-Naturwissenschaftlichen Fakultät:
Prof. Dr. Elmar Kulke
Gutachter:
1. Prof. Dr. Thomas Lohse 2. Prof. Dr. Elisa Bernardini 3. Prof. Dr. Markus Böttcher
Tag der mündlichen Prüfung: 2.5.2017
This work presents a detailed study on the very-high-energy (VHE) γ-ray source HESS J1745-290 in direction of the Galactic Center using 10 years of data from the H.E.S.S. array of Cherenkov telescopes from 2004 to 2014 with the objective to search for variability of the γ-ray flux of this object. The question if HESS J1745-290 shows variability is of special interest, since the source is located at the same direction as the super-massive black hole Sgr A*. From the vicinity of this black hole variable radiation has been reported for different wavelength bands. The detection of a variability of the VHE γ-ray flux of HESS J1745-290 would favor the hypothesis of a connection of this object and Sgr A* which could not be confirmed so far.
The search for variability was performed for different timescales from minutes to years and indeed revealed evidence for variability in different statistical tests which will be discussed in detail together with systematic cross-checks. The study focuses on both variability with and without periodic character. While there is evidence for a long- term flux modulation with a period of 110 d at the 4.1σ significance level, the χ2 fit of a H.E.S.S. run-wise light curve from 2004–2014 shows variability at the 6.1σ level, which reduces to 4.5σ after adding a 10% systematic error to each flux measurement.
Also signs of variable behavior at a timescale of minutes were found at the 3.1σ level.
This tentative VHE short-term variability also shows quasi-periodic behavior as it was reported during infrared and X-ray observations of Sgr A*. Such a tentative long-term flux modulation with a period of 110 d has previously also been reported for the radio band.
Due to the similarity of time structure of the variability, which is reported for HESS J1745-290 in this thesis to observations of Sgr A* at other wavelength bands, the thesis will close with the discussion if these results can be considered to be first evidence for a link between HESS J1745-290 and Sgr A*.
Die folgende Arbeit beschäftigt sich mit der Quelle hochenergetischer Gammastrahlung HESS J1745-290, welche in Richtung des galaktischen Zentrums liegt und präsentiert die Analyse eines Datensatzes, der von den H.E.S.S. Teleskopen zwischen 2004 und 2014 aufgezeichnet wurde. Ziel der Untersuchung war es, eine zeitliche Variabilität des beobachteten Flusses festzustellen. Die Frage ob der Fluss von HESS J1745-290 variables Verhalten zeigt stellt sich, da sich die Quelle in der gleichen Richtung wie das supermassive schwarze Loch Sgr A* befindet, aus dessen Umgebung bereits variable Strahlung in verschiedenen Frequenzbereichen detektiert wurde. Die Beobachtung einer Variabilität des hochenergetischen Gammastrahlenflusses von HESS J1745-290 würde die Hypothese eines Zusammenhangs zwischen diesem Objekt und Sgr A* stützen, welche bis dato nicht bestätigt werden konnte.
Die Suche nach Variabilität wurde für verschiedene Zeitskalen von Minuten bis hin zu Jahren durchgeführt, wobei Hinweise für Variabilität in verschiedenen statistischen Tests gefunden wurden. Die Suche konzentriert sich auf Variabilität ohne periodischen Charakter sowie auf Periodizität.
Es wurden Hinweise auf eine Periode von 110 Tagen bei einem Signifikanzniveau von 4.1σ gefunden und auch der χ2 Fit einer H.E.S.S. Lichtkurve, die von 2004–2014 reicht, zeigt Variabilität bei einem Signifikanzniveau von 6.1σ, welches sich nach Anwendung eines systematischen Fehlers von 10% auf 4.5σreduziert. Auch Anzeichen für Variabilität auf einer Zeitskala von Minuten wurden gefunden. Diese Variabilität auf einer Zeitskala von Minuten zeigt quasi-periodischen Charakter, ähnlich derer, welche während Infrarot- und Röntgenbeobachtungen von Sgr A* festgestellt wurde.
Die Möglichkeit einer Verbindung zwischen HESS J1745-290 und Sgr A* und ins- besondere auch die Fragestellung, ob die hier präsentierten Ergebnisse als erster Hinweis auf solch einen Zusammenhang gewertet werden können, werden Thema der Diskussion am Ende der Arbeit sein.
Contents 7
1 Introduction 9
1.1 Astroparticle Physics: Studying Cosmic Accelerators . . . 10
1.2 The non-thermal Milky Way and the Galactic Center Region . . . 11
2 The High Energy Stereoscopic System 13 2.1 Air Showers and Cherenkov Light Cones . . . 14
2.1.1 The Cherenkov Effect . . . 14
2.1.2 Different Types of Air Showers . . . 15
2.2 The H.E.S.S. Experiment: Science and Detector . . . 17
2.2.1 The H.E.S.S. Site . . . 17
2.2.2 H.E.S.S. Phase I . . . 19
2.2.3 H.E.S.S. Phase II . . . 19
2.2.4 The different Subsystems of the H.E.S.S. Detector . . . 20
2.2.5 Calibration and Detector Simulations . . . 24
2.2.6 Detector Simulations . . . 25
2.3 Reconstruction and Particle Identification . . . 25
2.3.1 The Standard Hillas Reconstruction . . . 26
2.3.2 Particle Identification: Moving towards Pattern Recognition . . . 28
3 The Galactic Center Source HESS J1745-290 35 3.1 Variability of Sgr A*: multi-wavelength Results . . . 36
3.1.1 A Periodicity at about 110 Days . . . 36
3.1.2 Short and weak Flares during Infrared and X-Ray Observations . 38 3.1.3 The G2 Gas Cloud . . . 43
3.1.4 Giant X-Ray Flares . . . 44
3.1.5 Estimates of the Frequency of Galactic Center Flares . . . 44
3.2 Summary of H.E.S.S. I Results from 2006 . . . 45
3.3 Summary . . . 47
4 Search for Variability of HESS J1745-290 49 4.1 Methods: Variability and Periodicity Tests . . . 49
4.1.1 Light Curves: A Definition . . . 49
4.1.2 Periodicity Tests . . . 50
4.2 Data Analysis . . . 64
4.2.1 The Dataset, Cuts and Background Method . . . 65
4.2.2 Skymap of the Galactic Center Region . . . 67
4.2.3 The Spectrum of HESS J1745-290 with Loose Cuts . . . 69
4.2.4 Search for a long-term Variability . . . 73
4.2.5 Search for a Variability at a Timescale of Minutes . . . 98
4.2.6 Run-wise χ2 Test of the post-cut Event Rates . . . 117
5 Summary and Outlook 127 5.1 Summary . . . 127
5.2 Comparison to MWL Observations . . . 128
5.3 Cross-Checks and External Data . . . 131
5.4 Candidate Mechanisms explaining the observed Variability . . . 133
5.5 Conclusion and Outlook . . . 134
Bibliography 139
A Simulated Monte Carlo Event Displays for gamma-like Events 153
B Studies with respect to Broken Pixels 155
Introduction
The cosmic rays (CRs) were discovered by Victor H.E.S.S. in August 1912 during a balloon flight from Aussig to Pieskow which was preceded by several other measurements.
His important discovery, which was awarded with the Nobel Prize in physics in 1936, can be considered to be a mile stone on the way towards modern particle physics and also the starting point of a new discipline: The astroparticle physics which needed about 60 years after the discovery of cosmic radiation to emerge as independent discipline next to the classical accelerator based particle physics. Today the astroparticle physics studies the most violent phenomena in the universe and particles reaching energies that no accelerator on earth has reached so far. While the Large Hadron Collider (LHC) in Geneva reaches a center of mass energy of 13 TeV since 2015, modern γ-ray detectors measure particles at several 10 TeV. Even so-called PeVatrons, accelerators reaching energies in the petaelectronvolt (PeV) range are in sight now. While the majority of CRs reaching the atmosphere of the earth consists of protons and heavier nuclei, there are also highly energetic electrons (positrons) and photons.
There should be a correlation between regions of CR production and the direction of high energy photons which are emitted due to the interaction of CRs with magnetic fields and matter in their environment. An important hadronic production mechanism of high energy photons is via the production of secondary π0 mesons which decay into photon pairs afterwards. Photons with an origin in theπ0 decay usually have larger energies than those which are produced via synchrotron emission. Their minimum energy is half the mass of theπ0 in the rest frame which implies that photons which are produced via the π0 decay have energies of at least ∼70 MeV. In general photons with energies ≥1 MeV are called γ-rays. Next to these hadronic production mechanisms γ-rays can also be produced by leptons. While at MeV energies the synchrotron emission is the dominant process, for the γ-ray production in the GeV and TeV range Inverse Compton (IC)
scattering where γ-rays (which could be produced as synchrotron radiation before) gain energy from an interaction with relativistic electrons is the most important mechanism.
The photons produced by these different mechanisms cover a large energy range up to 100 TeV. Due to the fact that γ-rays are not bent by magnetic fields like electrons and protons they form an excellent tool to observe the universe at high and very high energies (VHE, E > 0.1 TeV). From the 70ies of the last century onwards systematic efforts were undertaken to study γ-rays and the γ-ray astronomy was born as a new discipline. While the γ-ray sky at energies of about 100 MeV is dominated by Galactic and extragalactic diffuse emission, more and more point like sources were discovered in the GeV and TeV range.
Since photons with energies larger than 10 eV do not pass the atmosphere of the earth there are only two methods to observe them: The first strategy is to go to space and observe with a satellite experiment. Prominent detectors of this type are the Energetic Gamma Ray Experiment Telescope (EGRET) [102] and its successor the Fermi Large Area Telescope (LAT) [36]. On the other hand, there is a ground-based approach using so-called Imaging Atmospheric Cherenkov Telescopes (IACTs) which use the earth’s atmosphere as calorimeter and detect cosmicγ-rays indirectly using the Cherenkov light of the showers which the initial particles create by their interaction with the atmosphere.
The concept of IACTs to observe cosmic γ-rays was introduced at the Fred Lawrence Whipple Observatory [78] in the late 1960s with a single 10 m diameter Cherenkov telescope and led to a successful series of discoveries which is continued by recent projects like the High Energy Stereoscopic System (H.E.S.S) [107] among others which will be used in Chapter 2 to illustrate the concept of IACTs in more detail. Although significant progress in understanding the origin of CRs was made during the last decades, it is still not fully understood. Concerning the highly energetic γ-rays, several categories of objects emitting such radiation could be identified some of which will be shortly explained in the next section. The last section of this introductory chapter is dedicated to the central region of the Milky Way from the VHE γ-ray perspective with focus on the central γ-ray source HESS J1745-290 which is the main subject of this thesis.
1.1 Astroparticle Physics: Studying Cosmic Accel- erators
One of the key science targets of astroparticle physics is studying cosmic accelerators and mechanisms producing the high energy particles arriving at earth. Although this mission is far from being finished at the time of this writing, there are several successful models
explaining parts of the phenomena. The most prominent acceleration mechanism is the so-called Fermi or diffusive shock acceleration [68] which charged particles experience when they are reflected at so-called magnetic mirrors. These magnetic mirrors can be formed by magnetic inhomogeneities which are related to shock-waves for example. By repeated downstream to upstream reflections the particle gains energy and the resulting spectral distribution created by this process can be approximated by a power-law. This process is also called first order Fermi acceleration. Shocks which are strong enough for a significant acceleration are produced in violent astrophysical processes like super- nova explosions which attain velocities of the order of 104km s−1 [69] or during jet- like outbursts close to super-massive black holes, mostly located within so-called active galactic nuclei (AGNs) which can be found in the center of many galaxies. Furthermore, charged particles can also undergo second order Fermi acceleration which takes place in case they interact with randomly moving magnetic mirrors (e.g. magnetized gas clouds).
Whenever such a mirror is moving towards the direction of movement of the reflected particle it will gain energy by the reflection. For theγ-ray astronomy this means that the observed HE and VHE emission is expected to be associated with supernova explosions and their leftovers and pulsar-wind nebula (PWNe) and their corresponding pulsars on the one hand and super-massive black holes on the other hand. Indeed many of these objects have been detected as γ-ray sources meanwhile and form an important component of our picture of the non-thermal universe and our galaxy.
1.2 The non-thermal Milky Way and the Galactic Center Region
The discoveries made during the last 20 years dramatically changed the non-thermal picture of our own galaxy, the Milky Way. While at the beginning of this century only a couple ofγ-ray sources like the famous Crab Nebula [48] were known, until 2015 more than 150 individual TeV sources could be identified, whereby H.E.S.S. and Fermi LAT made a significant contribution. Mostγ-ray sources within our galaxy could be identified as SNRs or PWNe but there are several γ-ray sources the nature of which has not yet been revealed.
An especially interesting point-like source of this kind is located right in the direction of the center of our galaxy where also the prominent black hole Sagittarius A* (Sgr A*) can be found. Sgr A* was discovered as a strong radio source in 1974 at the National Radio Astronomy Observatory (NRAO) [51] and its mass could be determined to be (4.31±0.06|stat ±0.36|R0)×106 solar masses (MSun) [91]. It is located at a distance
of about 8.5 kpc. The discovery of this VHE γ-ray emission from the direction of the Galactic Center (GC) by CANGAROO-II [27] and H.E.S.S. [51] in 2004 triggered many speculations about its origin. The corresponding source is called HESS J1745-290 [10]
in the following since due to the complexity of the GC region there are different sources and processes which could contribute to the observed signal. Despite the positional co- incidence of HESS J1745-290 with Sgr A* within the angular resolution of the H.E.S.S.
detector it could not be proven so far that the GeV and TeV emission from the GC is connected to this object. While the supernova-remnant (SNR) Sgr A East [92] could be ruled out to coincide with HESS J1745-290, the nearby pulsar wind nebula (PWN) G359.95-0.04 [109] is still a candidate to cause the VHE γ−ray emission from the direc- tion of the GC. Also exotic explanations like dark matter could play a role, although no evidence for that has been found so far. Furthermore, in 2013 the Nuclear Spectroscopic Telescope Array (NuSTAR) [1] discovered a new magnetar only 2.8′′ away of Sgr A*
which pulsates with a period of 3.76 s and showed an increased activity in 2013 [105].
The pulsed flux was observed at keV energies. A possible approach to understand the γ-ray emission from the GC is to search for a variability in time. While the diffuse emission, PWNe or a hypothetical Dark Matter (DM) halo produce a constant flux, the environments of black-holes are usually highly variable at all energies. This could be shown for many AGNs so far.
For a better understanding of HESS J1745-290 it is essential to investigate if this source shows a constant γ-ray flux like a typical PWN or SNR or indeed shows signs of variability which would support the hypothesis of a link between Sgr A* and HESS J1745-290. The main objective of this thesis is to perform such a study with the full H.E.S.S. dataset available by the time of writing and shed some light on this interesting question. After a short description of the H.E.S.S. experiment in Chapter 2 there will be an introduction to the data analysis techniques and statistical tests which are used in Chapter 3 and 4, followed by a discussion and interpretation of the results presented there.
The High Energy Stereoscopic System
The H.E.S.S. experiment, which is located in the Khomas Highland in Namibia, is an array of 4+1 IACTs and has been inaugurated in 2003. During its first phase (H.E.S.S. I) it consisted of four Cherenkov telescopes of the same type with a diameter of 12 meters.
In 2012, the array was extended by another, larger telescope starting the second phase of the H.E.S.S. project (H.E.S.S. II). The H.E.S.S. array can be operated in different setups either combining all 5 telescopes (stereo or hybrid mode) or using the large telescope (CT5) in stand-alone mode while operating the four small telescopes (CT1 - CT4) in a different sub-array. The full H.E.S.S. array with its five telescopes is shown in Fig. 2.0.1.
The following sections will present an overview on the principles of Cherenkov astronomy together with an introduction to the available methods of event reconstruction and data analysis.
In order to be able to detect cosmicγ-rays at the earth’s surface, the use of Cherenkov telescopes is necessary since high-energyγ-rays cannot penetrate the earth’s atmosphere without interaction and can only be detected indirectly from the ground level via the Cherenkov light cones, the secondary shower particles produce on their way through the atmosphere. In order to define the requirements for such a detector, one needs a detailed understanding of these Cherenkov air-showers. The design of the H.E.S.S. detector has been optimized with respect to being as sensitive as possible and also providing a relatively large Field of View (FoV) with a diameter of 5°. The angular resolution of the detector is better than 0.1° and its energy resolution is about 20%. Like that, it is possible to detect aγ-ray source with 1% of the flux of the Crab Nebula (which can be considered to be a standard-candle inγ-ray astronomy) at the 5σ level in only about 25 hours (with the H.E.S.S. I sub-array at low zenith angles≤20°). An important challenge
Figure 2.0.1: The H.E.S.S. Array with its 5 telescopes. The large telescope CT5 is surrounded by the four telescopes from H.E.S.S. I. Image: The H.E.S.S. Collaboration.
in Cherenkov astronomy is to distinguish the showers which are induced by primary γ- rays from a huge number of background showers, which are triggered by hadrons and also by electrons. Furthermore, light cones by muons contribute to the background.
The differences of showers with a γ-ray as primary particle and those having a different origin will be discussed in this chapter. The background suppression especially gets very important when CT5 is operated in mono-mode at energies below 100 GeV where the background contribution increases dramatically.
2.1 Air Showers and Cherenkov Light Cones
The following paragraphs give an overview of the characteristics of the different shower types which are relevant for Cherenkov astronomy after a short definition of the Cherenkov effect.
2.1.1 The Cherenkov Effect
The Cherenkov effect was discovered by the Russian physicist Pawel Alexejewitsch Tscherenkow and refers to the electromagnetic radiation which is emitted whenever a charged particle moves in a dielectric medium with a velocity which exceeds the speed of light in that medium. A prominent example of this radiation is the bluish light which is emitted by a nuclear reactor operated under water. The cone-like light emission is
created by a polarization of the surrounding medium in combination with a constructive interference of the electromagnetic waves which are created during that process. This constructive interference only takes place when the velocity of the particle is greater than the speed of light in the medium the particle is crossing, otherwise the elementary waves created during the interaction of the particle with the surrounding medium in- terfere destructively and do not form a coherent wave. The Cherenkov light is emitted in a cone with a characteristic opening angle θ with cos(θ) = nβ1 where β is v/c. Only particles with β > 1/n where n is the refractive index of the medium emit Cherenkov radiation and low mass particles dominate the emission. Considering the fact that the refractive index of the atmosphere varies as a function of the altitude, a shower on its way to the ground experiences different Cherenkov angles between 0.5° and 1.4°. The wavelength of the Cherenkov light of an air-shower varies between 250 nm (blue) and 700 nm (red).
2.1.2 Different Types of Air Showers
High-energy particles interacting with the atmosphere always induce cascades of sec- ondary particles which form an air-shower. It can either be triggered by electrons positrons or γ-rays with only electromagnetic interactions being present or by hadronic particles like protons or heavy nuclei, where the strong interaction is dominant and induces more heterogeneous cascades. During the development of such a shower the number of particles increases with each interaction or radiation length until the shower maximum is reached. At the same time, the energy per particle decreases, which results in the shower fading away after reaching its maximum.
2.1.2.1 Electromagnetic showers
For Cherenkov astronomy the electromagnetic showers are the most interesting class of showers since they allow the observer to detect VHE γ-rays. For this class of showers a combination of pair production and bremsstrahlung are the most important energy loss mechanisms. Ionization of molecules in the air only plays a role when the average particle energy reaches the critical energy of about 81 MeV in air where those losses are comparable to those by bremsstrahlung. At this point the shower development stops due to the electrons and positrons not being able to produce secondary photons any more. The first interaction for primary particles with energies > 10 GeV, which is shown schematically in Fig. 2.1.1, typically takes place at a height of several 10 km above the sea level. The depth of the shower maximum is approximately proportional
Figure 2.1.1: The typical start of an electromagnetic shower. The initial γ-ray triggers a cascade of secondary particles via pair production.
to the logarithm of the energy of the primary particle (a fact which also can be used for the energy reconstruction of the initial particle). While the pair production and bremsstrahlung processes are strongly forward-directed in the laboratory frame due to the large boost of the primary particle, the lateral shower development is dominated by elastic collisions of the shower particles with air molecules. An important quantity describing the transverse extension of the shower is the Moliere radius, which is the radius of a cylinder containing on average 90% of the shower particles. A typical Moliere radius of an electromagnetic shower at sea level is about 60 m to 90 m, which means that these showers are rather narrow and give a good guideline for the dimension of a telescope array together with the Cherenkov angle of the shower. Although showers induced by γ-rays and electrons are dominated by the same processes, there is a slight difference between these two shower types: For primary γ-rays the first interaction is a pair-production process while electrons (positrons) undergo bremsstrahlung first. Since the mean free path length for pair production at high energies is larger than the radiation length for bremsstrahlung, the height where the first interaction takes place slightly differs for the two types. Therefore, also the height of the shower maximum for electron- andγ-induced showers differs by about 5%. This is the only difference between these two classes of showers, which makes it difficult in practice to distinguish between them.
2.1.2.2 Hadronic showers
Hadronic showers show a completely different behavior than the electromagnetic showers due to the fact that weak and strong interaction processes are involved. This leads to hadrons and leptons as secondary particles. During the first interaction of cosmic ray hadrons with nuclei in the atmosphere, secondary pions (and also protons and kaons)
are produced which then induce the secondary cascades. Since strong inelastic scatter- ing takes place, the final state mesons usually have sizable transverse momenta which results in a large lateral extension of hadronic showers in comparison to the narrow electromagnetic showers. While the neutral pions, which are produced in the initial re- action decay into two photons, induce an electromagnetic sub-shower, the charged pions (π+, π−) have muons and neutrinos as final states after their decay. While the neutri- nos obviously escape without an interaction, the muons often reach the ground level without further interactions emitting a ring-like cone of Cherenkov light. In the same way, the hadronic showers give a heterogeneous picture in the camera and show larger fluctuations across the telescope array: While a particular telescope sees a sub-shower which has been induced by a γ-ray from the pion decay a second telescope may see a clear muon ring. Also the shower width of hadronic showers is larger. That way they can be identified relatively easily by combining the information from several telescopes (in mono mode this is more complicated of course). Also the integrated pixel inten- sity measured in a specific telescope fluctuates stronger between the different telescopes for hadronic showers. The intensity maximum of hadronic showers is located deeper in the atmosphere than for γ-induced showers due to a larger free mean path lengths of hadrons.
2.2 The H.E.S.S. Experiment: Science and Detector
The next sections will describe the H.E.S.S. experiment in more detail and give an overview of past and ongoing science projects at H.E.S.S. Relevant sub-systems and data taking procedures will be introduced together with references providing further information.
2.2.1 The H.E.S.S. Site
H.E.S.S. has been built in the southern hemisphere in order to be able to observe the Galactic plane and especially the Galactic Center region at high elevations. The telescope array is located in Namibia, about two hours south-west of the capital Windhoek at the exact position with the coordinates 23°16’18” S, 16°30’00” E at an altitude of 1800 m above the sea level. The site was chosen due to its excellent climatic conditions for astronomical observations.
Figure 2.2.1: The inner Galactic plane as seen in the H.E.S.S. Galactic Plane Survey, where the color scale describes the statistical significance for an excess within a 0.22°
radius at each position. Figure from Ref. [53].
2.2.2 H.E.S.S. Phase I
The H.E.S.S. experiment was operated as four-telescope array from 2004 until 2012 and is considered to be the most successful ground-basedγ-ray observatory, achieving many high impact results like the H.E.S.S. Galactic Plane Survey (HGPS), which revolution- ized our image of the Galactic plane in the TeV range revealing an unexpectedly large and diverse population of over 60 TeV γ-ray sources. An excerpt of the result of the HGPS is shown in Fig. 2.2.1. Before this scan, less than 10 sources where known at TeV energies. Further highlights of Phase I were the stringent upper limits on the velocity- weighted dark matter (DM) annihilation cross-section for the Galactic Center halo, first evidence for a diffuse large scale emission in the Galactic plane or the first detection of a pulsar wind nebula in the Large Magellanic Cloud (LMC). Furthermore, shell-type SNRs could be identified asγ-ray sources. Some extragalactic highlights are the discov- ery of VHE γ-ray emission from the star burst galaxy NGC 253 or an upper limit on the extragalactic background light at optical/near-infrared wavelengths which, implies that the intergalactic space is more transparent toγ-rays than expected. Analyzing the full dataset of the Galactic Center region, H.E.S.S. recently found hints for Pevatrons (acceleration mechanisms reaching the PeV range) by showing that the diffuse emission close to the Galactic Center does not have a cut-off. The search for pulsed emission fromγ-ray pulsars in data from the four telescope array did not lead to any conclusive results, which put high expectations on Phase II [5, 11, 52,60, 89, 100].
2.2.3 H.E.S.S. Phase II
In 2012 the new telescope in the center of the array was inaugurated. With its large mirror area (diameter 28 m) it is the largest optical telescope in the world at the time of this writing. The new telescope was built to lower the energy threshold of the array by almost an order of magnitude from 100 GeV to about 10 GeV, which gives an excellent overlap with the Fermi LAT for this energy range, which is a satellite experiment. The low threshold was motivated by different physics objectives:
• Pulsed emission: Fermi LAT observed many γ-ray pulsars in this energy range the most promising of which being the Vela pulsar. First H.E.S.S. II results from 2015 confirmed the pulsations of the Vela pulsar seen by Fermi LAT at energies
>20 GeV [34].
• Gamma Ray Bursts (GRBs): So far no ground-based γ-ray telescope ever managed to detect a GRB. Due to its large effective collection area and its fast drive system CT5 is an ideal telescope to search for such GRBs.
• Dark Matter: The search for an annihilation signature of Weakly Interacting Massive Particles (WIMPs) or other exotic dark matter particles still is a hot topic. In 2012 there has been a promising signature in Fermi LAT data at 130 GeV close to the GC, which would have been an ideal target for H.E.S.S. II. However, this signal could be ruled out by the Fermi Collaboration before H.E.S.S. II could collect sufficient statistics to make a statement on this topic.
• The Galactic Center region: Due to its good angular resolution CT5 would be the ideal telescope to study the Galactic Center region and find unresolved sources in that region.
The list above shows that there is a diverse spectrum of science targets for H.E.S.S. II some of which still are waiting to be exploited.
2.2.4 The different Subsystems of the H.E.S.S. Detector
The different subsystems of the H.E.S.S. detector will only be mentioned very briefly here since a good overview can be found in the references quoted at the end of each paragraph. An exception here is the Data Acquisition System (DAQ), since the author has been working as an expert for this system from 2012 to 2015. The focus here will be especially on the steps which were necessary to prepare the system for the integration of CT5.
2.2.4.1 The Cameras of CT1 - CT4
The telescopes CT1 to CT4 are all equipped with cameras of the same type. Each camera in total contains 960 photo-multipliers (PMTs), whereby each PMT covers a solid angle of 0.16° in diameter. This gives a total field of view of 5° when the sub-array CT1 - CT4 is operated in a stand-alone mode. The PMTs are grouped in so-called drawers which contain 16 pixels each and contain the electronics which is needed for triggering and data-acquisition along with the power supplies for the PMTs. The PMTs have very short reaction times of a few nano-seconds and are sensitive in a wavelength-range of 300–700 nm which perfectly fits the Cherenkov light-spectrum created by an air-shower [50, 108].
2.2.4.2 The CT5 Camera
The CT5 camera mimics the design of the H.E.S.S I cameras with the difference of providing considerably more PMTs in order to achieve a better angular resolution and
a lower energy threshold. The total number of PMTs used for the CT5 camera is 2048. CT5 is the first telescope in the world where the camera can be automatically removed from the telescope and put into a shelter in case of bad weather conditions or if maintenance is needed. In total the process of loading or un-loading the camera takes about 45 minutes [56].
2.2.4.3 The Trigger
The H.E.S.S. trigger system consists of two levels. The first level is installed directly at the individual telescopes. An individual telescope triggers when a threshold of 5.3 photo- electrons is exceeded for at least 3 PMTs within a sector of 64 PMTs in a time window of 1.5 ns. The Central Trigger forming the second level is combining the signals from the different telescopes and requires that at least two telescopes have triggered within a time window of 80 ns. In case an event is accepted by the Central Trigger, an event number is assigned together with a GPS time-stamp and distributed to all telescopes, which were involved in the trigger. Then the data is read out from the analogue ring buffers in the telescopes where it is temporarily stored and sent to the DAQ computing cluster where it is ultimately written to disk [90].
2.2.4.4 Data Acquisition, Run Control and Monitoring
In contrast to other Cherenkov telescope systems like VERITAS [81] and MAGIC [24], H.E.S.S. uses a combined system for Data Acquisition, Run Control and Monitoring which is referred to as DAQ in the following. It was designed to meet all science re- quirements of the H.E.S.S. array: An important factor here is the dead-time of the array which may not be increased artificially by the DAQ. Therefore, the system has to be able to deal with a data rate of at least 50 MB/s where the four H.E.S.S. I telescopes contribute 900 Hz with an event size of 4.5 kB and CT5 up to 3.4 kHz with an event size of 10 kB. In order to have some tolerance also with respect to possible further H.E.S.S.
upgrades at least 80 MB/s are required for the whole system. Furthermore, sufficient storage capacities in Namibia are needed in order to store the data until it has been safely transferred to the European data centers in Lyon (France) and Heidelberg (Ger- many). For the transfer to Europe the data are written on tapes, which are shipped via airplane to the two European data centers each of them receiving its own copy.
From the hardware perspective the H.E.S.S. DAQ during H.E.S.S. I consisted of 5 servers and 10 worker nodes. These are linked via optical 1 Gigabit/s connections to the telescopes via two main network switches. Furthermore, there is an internal, independent network for the communication between the different servers and nodes. A sketch of the
Figure 2.2.2: Schema of the network layout and storage units of the H.E.S.S. DAQ system. The telescopes are linked to the worker nodes and data servers in the control room building (upper part of the figure) via two main switches. All connections work at the gigabit/s level. The green lines illustrate a physically separated networks for the internal communication between data servers and worker nodes, while the black lines indicate the connections from/to the telescopes. The figure was taken from Ref. [3].
system layout can be found in Fig. 2.2.2. The task of the worker nodes is to receive the data from the telescopes and to write it to the actual storage space via a network file system (NFS). The nodes do not provide any facilities for long-term storage in contrast to the servers. Initially, only one of the servers was dedicated to storage providing 12 TB for data, while the remaining four servers were used for different tasks like hosting the H.E.S.S. database or user home directories and virtual machines which are needed by the cameras. Like that the initial DAQ system had 12 TB for storage while the rest of the available disk space was dedicated to other tasks. During H.E.S.S. I this was sufficient, however, to meet all requirements of a full hybrid system the H.E.S.S. DAQ system needed to be upgraded in 2013:
Considering the rate of the full system, under good instrumental and weather condi- tions one would expect a data rate of about 10 TB/months during the winter time when the nights are long. Therefore, 12 TB clearly were not sufficient to guarantee to store the data from at least three months on site which is also a design requirement. Furthermore, the data rates were limited by the fact that all the data had to be written via a single network interface to the same hardware RAID (RAID stands for Redundant Array of Inexpensive Disks, self-explaining) by both network transmission and the I/O rate of the RAID controller. In order to increase the bandwidth and storage capacity for H.E.S.S. II (and further upgrades), two more data servers were integrated into the system together with an upgrade of the disks forming the hardware RAID: For the two new servers 3 TB disks were used instead of 1 TB disks which were used before (the difference in price is only marginal). Furthermore, a third machine (the previous data server) was also up- graded with 3 TB disks. Like that there are three servers with a capacity of 36 TB each which could be mounted to a single virtual storage unit using GlusterFS [87], which is a network file system allowing parallel I/O operations in contrast to NFS, which is a serial filesystem. This allows a significantly higher bandwidth since the workload is distributed between the three physical servers in terms of disk usage and network bandwidth. In total the new storage units provide a volume of 108 TB for data only. Although this number may appear large it is justified for multiple reasons:
• Redundancy: The current size of the system allows one of the data servers to fail without the storage capacity of the system dropping below an acceptable size.
• Readiness for the H.E.S.S. I upgrade: The current setup will also provide sufficient bandwidth and storage capacity after an upgrade of the H.E.S.S. I cameras end of 2016.
• Delays in tape writing: Due to occasional delays in the tape writing procedure
temporarily up to 60 TB of disk space were used (about 5 months) since before deleting the data in Namibia it has to be verified that it has correctly arrived in Europe.
In a second step an additional GlusterFS partition was created combining two of the 12 TB servers as backup space and also for the system log-messages. It was not possible to create a merged volume of all five servers since the version of GlusterFS, which is used in Namibia, only supports equally sized backends for a particular volume.
A further important feature of the H.E.S.S. DAQ system is its ability to respond fully automatically to target of opportunity (ToO) alerts like GRBs. Information from the Gamma Ray Burst Coordinates Network (GCN) can be fed directly into the H.E.S.S.
DAQ system and a dedicated process, the GCNAlerter, reacts to the input automatically in case such an event occurred at coordinates which are observable. The observation starts automatically and no human input is needed, which significantly reduces the reaction time to the system for GRB observations to a range from only 30 seconds to 1 minute. More information about the H.E.S.S. DAQ can be found in Ref [3].
2.2.5 Calibration and Detector Simulations
The calibration and detector simulation are essential for dealing with the data recorded by the telescopes and producing solid analysis results. In the following a very short overview on the calibration procedure and the available simulation packages is given.
2.2.5.1 Calibration
Since the quality of the data taken with H.E.S.S. depends on both environmental and instrumental conditions, the data need to be calibrated before an actual analysis can be performed. For that purpose calibration runs are taken at regular intervals (see Ref.
[9]). During the calibration the ADC values of the PMTs are converted to intensity maps in units of photo electrons (pe). For the ADC values there are two channels available:
The low gain channel is used for strong signals which range up to 1600 pe while the high gain channel covers signals up to 200 pe. The calibrated intensities Ii can be obtained via the formula
Ii
pe = (Ai−Pi)
γi ×F (2.2.1)
where Ai are the ADC values for the channel i, which can either be low or high gain, Pi is the pedestal value for a PMT,γi the electronic amplification factor of the channel and F the flat-fielding coefficient. The pedestal is an additional signal which is caused
by electronic noise without light input. Furthermore, the night sky background (NSB) can cause noise at usually 1 pe level. Both effects have to be subtracted from the actual ADC value. The time dependence of the pedestal of each PMT can be extrapolated from the average value of the first two minutes of an observation run. The electronic amplification factor is measured every two days with a dedicated calibration run, the SinglePE run. The flat-fielding coefficients correct for inhomogeneities of different PMTs and are obtained by exposing the camera to a uniform illumination. At H.E.S.S. there are currently two different calibration chains in use, one being maintained by the MPIK in Heidelberg and the second one independently maintained by the French institutes of the collaboration.
2.2.6 Detector Simulations
Another crucial requirement for an actual data analysis are the detector simulations, which are referred to as Monte Carlo (MC) simulations in the following. Like for the calibration there are two independent MC simulation packages available in H.E.S.S.
The first, mainly maintained by MPIK is based on the CORSIKA package [49] for the shower simulations, while the French based counterpart is using CASCADE [2]. Both of these tools can be used to simulate showers induced by γ-rays and also other particles like protons. Models for the different interaction types are implemented. A full set of simulations covers a wide range of zenith angles, source offsets from the center of the camera and azimuth angles. For the simulation of the instrument response to these showers there are again two packages: For the Heidelberg chain simtelarray [23] is used, while in the French MC chain the response of the detector is simulated with SMASH [57]. Hereby the behavior of the full H.E.S.S. array is mimicked by simulating the two- level trigger system and taking into account properties of the mirrors (e.g. reflectivity) or the quantum efficiencies of the PMTs and also the detector electronics. The output from these simulation chains can be processed by the analysis software as if it was real data.
2.3 Reconstruction and Particle Identification
At H.E.S.S. there is a large variety of tools available which facilitate the analysis of the calibrated data. For briefness, the following section only focuses on those methods which are relevant for the actual data analysis in Chapter 4 together with the presentation of promising ideas which emerged during attempts of improving the current data analysis
Figure 2.3.1: The definition of the Hillas parameters (Figure from Ref. [31]).
methods.
2.3.1 The Standard Hillas Reconstruction
The classical method of shower reconstruction and particle identification in γ-ray as- tronomy is the Hillas method, which is well explained in Ref. [31]. It is a geometric approach simplifying the recorded image to an ellipse in the camera plane, which can be described by few parameters, which are discussed in the following. On these parameters box cuts can be applied.
2.3.1.1 Hillas Parameters
The so-called Hillas ellipses (shown in Fig. 2.3.1) are characterized by the following parameters:
• Length L and width W of the ellipse.
• The so-called size which is equivalent to the total image amplitude.
• The nominal distancedwhich is defined as the distance of the center of the camera and the center of gravity of the image.
• The azimuthal angle φ of the image main axis and the orientation angle α. Since the stereoscopic reconstruction approach was introduced by HEGRA and followed up by H.E.S.S. and later also MAGIC, the direction of the initial particle can be simply
Config. W¯smin W¯smax Lmins¯ Lmaxs¯ θ2 pe d
STD -2.0 2.0 -2.0 2.0 0.0125 70 2.0
Loose -2.0 2.0 -2.0 2.0 0.04 40 2.0
Table 2.1: The H.E.S.S. I STD cuts versus the Loose cuts.
reconstructed by the intersection of the main axis of the ellipses observed by the dif- ferent telescopes, which are projected into a common plane. The reconstructed energy is proportional to the signal amplitude measured in all telescopes. For the background suppression also a simple concept is used: For the so-called scaled cuts technique the width and length of the ellipses are compared to their expectation value and variance obtained from simulations, as a function of impact distance and charge. The scaled width and scaled length can be defined as:
Ws=W−< W >
σW , Ls=L−< L >
σL , (2.3.1)
For more telescopes involved these parameters can be easily combined to obtain the mean scaled parameters
W¯s =
PWs
√ntel, L¯s=
PSL
√ntel (2.3.2)
It could be shown from simulations that these mean scaled parameters are uncorrelated for γ-like events. Therefore, simple box cuts on these parameters can be applied as most simple discriminator betweenγ-like and hadron-like events since the distributions for hadron-like events show a tail towards large values of the ¯Ws and ¯Ls which is not present for γ-like events.
2.3.1.2 Hillas Cuts during H.E.S.S. Phase I:
The standardized Hillas cuts defined for H.E.S.S. Phase I are summarized in Ref. [8].
From the cuts which are summarized here mainly the Loose cut configuration is relevant for the analysis presented in this thesis. The Loose cuts were chosen since they provide the lowest energy threshold and largest integration radius of the official cuts for point like sources. In Table2.1 the Loose cuts are compared to the STD cuts.
2.3.1.3 Hillas Cuts for H.E.S.S. Phase II:
In order to be able to analyze H.E.S.S. II data, the author adapted the Hillas analysis which is implemented in the Heidelberg software to be operative for H.E.S.S. II data
Lmins¯ Lmaxs¯ W¯smin W¯smax θ2 pe CT 1−4 pe CT5 d
-2.0 1.4 -2.2 1.0 0.0132 30 50 2.0
Table 2.2: The H.E.S.S. II Hillas low energy cuts for the stereo mode.
as well. To reach this goal the lookup table scheme for the scaled parameters had to be adapted by introducing a new telescope type for CT5. This was necessary in order to correctly account for the larger mirror area of CT5 in comparison to CT1 - CT4.
This modified Hillas analysis was the first analysis chain to be operative for H.E.S.S.
II data within the HD software and has been used for the web summary in Namibia, which automatically analyzes the data from the last night on site as a first quality check. Furthermore, it was used for the training of other software chains like Impact [84]. Meanwhile, Impact is considered to be the standard method to analyze H.E.S.S.
II data within the HD software, however the available configurations use an amplitude cut of 70 photo electrons for CT1 - CT4, which leads to an energy threshold similar to that of the H.E.S.S. I STD cuts when CT5 is used in combination with CT1 - CT4. In order to have a cut configuration with a comparable energy threshold to the H.E.S.S.
I Loose cuts, a custom Hillas low energy cut configuration was defined for CT1 - CT5 which will be used in the following. The exact definition of these cuts can be found in Table 2.2.
2.3.2 Particle Identification: Moving towards Pattern Recog- nition
During the last years methods based on multivariate statistics have been used for various reconstruction and analysis tasks in different fields of physics ranging from energy recon- struction and particle identification at LHC experiments [114] to γ/hadron separation at H.E.S.S. [80]. This development was fostered by the TMVA software package, which has been developed at CERN [54]. The main advantage of these multivariate analysis (MVA) methods is that they are able to model non-linear correlations between the input variables.
The following section summarizes the studies which have been performed by the author with the intent of improving the current efficiency of γ/hadron separation at H.E.S.S. This initiated the promising approach of using state of the art pattern recogni- tion algorithms for particle identification tasks in ground based γ-ray astronomy. Only the basic ideas leading to this idea, which were considered to be worth to be put on record, are sketched here with the intention of not discussing them in too much de-
tail. A first implementation of the approach of reducing the γ/hadron separation to a pure pattern recognition problem has been realized in a master thesis where they are described in all detail (see Ref. [103]).
The idea of applying pattern recognition technologies for γ/hadron separation tasks at H.E.S.S. emerged from studying the distribution of pixel clusters in H.E.S.S. event displays, especially for low-energetic events with the intention of improving the separa- tion power of the existing chains. In Fig. 2.3.2a a typical simulated γ-like 4 telescope event with a reconstructed energy of 200 GeV is shown and in Fig. 2.3.2b the event dis- play of a typical hadronic shower with a reconstructed energy of 300 GeV can be found (also a 4 telescope event). For these event displays the information of all telescopes was projected into the same camera plane since this approach turned out to be very helpful when trying to distinguish between γ-like events and hadron-like events by eye. Self- experiments done by the author showed that doing aγ/hadron separation by hand for a region of interest using this representation one reaches an accuracy which is comparable to chains like TMVA. The quality factor or Q-Factor is the standard tool to quantify the performance of a set of cuts which is defined as:
Q= εγ εBG
whereεi=ni Ni.
The quantity εi defines the cut efficiency where ni is the number of events which pass the cuts. Ni stands for the number of pre-cut events. In case the analysis chain has a large separation power, the ratio ni/Ni is close to 1 for γ-like events, since most of them should be correctly identified as signal and small for the background, since the background should be rejected efficiently. The quality factor reached when doing the γ/hadron separation by hand is about 1.3 at zenith angles of 20° versus 1.2 for the simple Hillas chain.
The study on manual γ/hadron separation showed that the event displays carry suf- ficient information to successfully complete the task without additional information. An interesting question is upon which information a human event selector looking through thousands of event displays is relying in order to be able to quickly categorize an event.
An important role for the decision plays the distribution of pixel clusters in the camera together with the distance between these clusters. The γ-like events ideally only show one cluster per camera, which is describing the Hillas ellipse, while for hadron like events multiple clusters per camera are expected. A second important criterion for the human
event selector is the implicit information of the direction of the shower which is encoded in the event display when projecting all ellipses in a common camera plane. As soon as a human identifies ellipses which are pointing towards an intersection point of their axis he will classify the event asγ-like event. It should be easily possible to train a good pattern recognition algorithm on this behavior. During intents to quantify the distribution of clusters throughout the camera two quantitative discrimination variables were defined:
1. The first was the ratio of the number of pixel clusters observed in an event divided by the number of telescopes which participated in recording the event. This ratio is expected to be close to 1 for γ-like events.
2. Another variable is the maximum distance between two pixel clusters in a telescope.
Attempts were made to introduce these new variables into the existing framework of TMVA chains at H.E.S.S. Although these new discrimination variables lead to a 10%
improvement in sensitivity during MC simulations, the improvement could not be sta- bly reproduced when applying the classifier to real data. This is likely related to the robustness of different types of classifiers in case there are differences of quality between training and application data. So far so-called boosted decision trees (BDTs) are the only MVA method which has been used for particle identification tasks at H.E.S.S. [80, 84], however, in terms of performance they are comparable with artificial neural networks.
Therefore, the question arises which of the two methods suits better for this particular use case application. The main advantage of BDTs is that their decision can be retraced in principle, while NNs behave as black boxes which do not offer this transparency. For their traceability BTDs are the preferred choice in the analysis of financial markets and in some countries even required by law in that field. However, for the use in counting experiments with large statistics this advantage vanishes in practice, since it is not pos- sible to trace the decisions. On the other hand, in Ref. [98] the authors conclude that one of the advantages of neural networks over other MVA methods (including BDTs) is that they could be shown to be more robust against deviations in data quality between training and test samples. Taking this as given, neural networks seem to be a good choice for astroparticle physics: The quality of observation data is subject due to dif- ferent external systematic effects like light pollutions, changing atmospheric conditions or broken pixels and therefore a considerable deviation between MC data and real data conditions can be expected.
These studies with additional discrimination variables were briefly mentioned here for completeness, although in the end they were considered to be not sufficiently robust with respect to varying systematic conditions like broken pixels or missing drawers:
(a) A 4 telescopeγ-like event at a reconstructed energy of 200 GeV.
(b) A 4 telescope hadron-like event at a reconstructed energy of 200 GeV.
Figure 2.3.2: Typical H.E.S.S. event displays where the information from all telescopes is projected into a common camera plane.
Broken pixels can affect the number of clusters in a camera and therefore affect the sensitivity of a method, which relies on such information in a negative way (e.g. a Hillas ellipse is split into two by a string of broken pixels). However, taking into account the overall distribution of pixels showing an intensity over the camera could be helpful for the discrimination task. For example, a human event selector would not get tricked by a string of missing pixels in case the event display still contains implicit direction information. Therefore, the approach of counting clusters was discarded in favor of the completely new idea of applying pattern recognition software for particle identification tasks to the whole event. Thereby the information from all pixels of an event is taken into account by projecting the intensities of all pixels into a common plane instead of reducing this information to certain discrimination variables.
From all available technologies so-called Convolutional Neural Networks (CNNs) were considered to be the most suitable approach to process large numbers of H.E.S.S. event displays. Meanwhile, applications based on this new type of neural networks have sur- passed all previous methods of computer vision and current software implementations considered to be advanced enough to be applied to complex problems in research. Like Multi Layer Perceptrons (MLPs), these CNNs are trained via back propagation [98] but unlike classical NNs they consist of groups of neurons which are arranged in a way that they focus on overlapping sub-groups of the input data. By this overlap in combina- tion with different layer types the networks reach a certain level of invariance regarding shifting, scaling and distortion of the shapes in the images which makes them superior to classical MLPs. This robustness could be of advantage in situations where the ob- servation data differs from the simulation quality e.g. due to broken pixels. One of the most advanced software frameworks supporting CNNs is called Caffe and was created by Yangqing Jia [58]. It consists of C++ and CUDA libraries with Python and Matlab wrappers. Therefore, Caffe supports modern GPUs guaranteeing fast image processing times (as fast as 2.5 ms per second at a state of the art GPU).
So far, this type of pattern recognition technology has not been used in high energy physics due to the short time since the software packages providing methods for training and application of these CNNs are available and their still very experimental state.
Since the performance of these networks is known to scale with their complexity, there still is huge potential for further improvements. An example where the technology of CNNs performed better than the human mind was the solution of so-called CAPTCHAs which are commonly used across the world-wide web as reverse Turing test to prevent all kind of automatized requests [43]. The NN implementation for solving these puzzles reached an accuracy of 99.8%, whereas humans only reached an accuracy of 94% on
average. The authors attribute their success to the level of complexity used in the networks and hypothesize that especially the depth of the network architecture was crucial to reach such a high accuracy. Their network implementation used 11-layers representing more than 5.0×107 parameters. This example shows the huge potential of this technology. The first application of this technology at H.E.S.S. in Ref. [103]
already has reached a separation power comparable to that of a typical Hillas chain with a limited training sample in a limited hardware environment. Since the performance of the setup is expected to rise with complexity there should also still be a potential of improvement for the γ/hadron separation at H.E.S.S. or for future observatories of ground based γ-ray astronomy.
The Galactic Center Source HESS J1745-290
It is well known that the vicinity of the super-massive black hole Sgr A* in the center of the Milky Way is emitting radiation at many wavelengths from sub-millimeter to hard X-ray and across this huge frequency-range Sgr A* is also known as a highly variable source. For the energy band from 200 MeV to 100 GeV, which is covered by Fermi LAT so far no evidence for a variability for the GC has been published, which may be due to the limited angular resolution of that instrument at low energies where one obtains sufficient statistics for a search for variability in combination with the complexity of the Galactic Center region. The angular resolution of Fermi LAT at 200 MeV is about 10°.
Since the VHEγ-ray source HESS J1745-290 is located in the same direction as Sgr A*, there is the hypothesis of a connection between these two objects, which could not be proven so far. The detection of a variability of HESS J1745-290 in H.E.S.S. data would be the first evidence for a connection between Sgr A* and HESS J1745-290.
Of course, variability alone would not prove such a link, since there is always the possibility that it is caused by an unknown source close to Sgr A*. For example, in 2013 the magnetar SGR J1745-2900 [61, 62] was discovered at an angular distance of only 2.8” of Sgr A*, which is often also called the GC magnetar. However, so far no magnetar has been detected at GeV or TeV energies. In case the search for a variability of HESS J1745-290 in fact finds evidence for a variable γ-ray flux, a comparison of the timescale of the effect to the behavior of Sgr A* at other wavelength bands can help to distinguish between Sgr A* and the GC magnetar. In the end, a possible variability would not exclude that the observed signal from the direction of HESS J1745-290 is a superposition of a contribution of Sgr A* with the DC flux of the PWN and the diffuse emission or the GC magnetar but it would still be sufficient to claim the detection of
VHEγ-ray emission from Sgr A* in case it shows a time structure or even a periodicity, which is similar to that which was observed for other wavelengths.
The following section summarizes what was found out so far about the variability of Sgr A* at different wavelengths and also addresses some of the numerous open issues and questions about the properties of this interesting object. Furthermore, recent results about the GC magnetar are briefly summarized, followed by the discussion of a first published H.E.S.S. search for a variability of HESS J1745-290 from the year 2006.
3.1 Variability of Sgr A*: An Overview of multi- wavelength Results
During the last ten years, several multi-wavelength campaigns have collected substantial evidence that the black hole Sgr A* shows variability at timescales of years, hours and also even at a timescale of minutes. While there are indications for a periodicity with a period in the order of 110 days in radio data, various infrared (IR) experiments and X-ray observatories found flares with durations from minutes to hours.
3.1.1 A Periodicity at about 110 Days
In 2005 evidence for a long-term modulation of the radio signal from the direction of Sgr A* with a period of about 100–120 days was found in data from the Very Large Array (VLA) in New Mexico. The period could be observed across a large range of frequencies.
The amplitude of this reported periodicity increases with decreasing wavelength. The effect is not yet fully confirmed, but there are various hypotheses which let it appear consistent in the context of expected properties of the black hole Sgr A* [71].
For the analysis nearly 20 years of monitoring data from the VLA were taken into account at wavelengths of 1.3 cm, 2.0 cm, 3.6 cm, 6.0 cm and 20 cm. The analysis shows a clear peak in the measured power-spectrum which ranges from 100 –120 days.
The authors claim that they have shown with the help of Monte Carlo studies that the probability of obtaining this signal by random processes is less than 5%, which means that their result is not very significant and the effect is rather weak. An independent confirmation of the effect would be helpful.
There are various scenarios which could explain such an effect like a companion on a close orbit or the precession of the black hole’s accretion disk. Since in terms of period this 110 days periodicity is 4 orders of magnitude larger than the known inner disk movement of Sgr A* the latter is unlikely to be related to that tentative long-term
Figure 3.1.1: Here the geometry of Sgr A* is shown including an inner non-thermal halo and the accretion disk. The precession of the accretion disk might cause a periodic variation of the flux, which can be observed from earth. In this model d is the distance to the GC, Rh denotes the radius of the inner non-thermal halo of Sgr A* and Rd the radius of the accretion disk. Sis the spin vector of the black hole, while Lis the angular momentum vector of the accretion disk. The angleγ between L and S causes a periodical eclipse of the emission from the non-thermal halo towards an observer on earth. The figure was taken from Ref. [71].
variability. Also the presence of a companion is not the best candidate for an explanation for this phenomenon following a simple reasoning: An object orbiting around Sgr A* with such a period would be found on an orbit with a radius of about 60 AU. The VLA would easily have been able to resolve such an object. In case of an object rapidly circulating around Sgr A* the black hole would also show a characteristic proper motion which could not be observed so far. Liu & Melia [65] proposed a mechanism in 2001 which could account for the observed modulation without the assumption of a companion: In the Kerr Metric each gravitational acceleration acquires a dependence on the poloidal angle relative to the black hole’s spin axis, which means that matter orbiting around a black hole below and above the equator is subject to a force towards the black hole’s equator.
That way the angular momentum vector may start to precess around the black holes rotation axis. As shown in Fig. 3.1.1, the spinning angular momentum vector of the disk results in the disk periodically, shielding parts of the emission from the non-thermal halo of relativistic particles in the vicinity of Sgr A* from an observer on earth. Thereby highly relativistic particles on our line of sight, which are created around the black hole close to the event horizon, are absorbed. This can lead to the observed periodic signal modulation. Therefore, the tentative radio-variability of Sgr A* could be caused by a simple combination of a periodical disk precession and the presence of a non-thermal halo close to the horizon of Sgr A*.
A disk precession at a timescale larger than 100 days can also be observed for other black holes and black hole candidates like the Galactic Center micro quasar 1E1740.7-2942 [96]. Theoretically, such a precession can also affect the inner dynamics of different sections of the accretion disk which may connect it to the observed short-term variability which is described in the following sub-sections.
In case Sgr A* shows such a non-thermal halo of relativistic particles, also a possible VHE γ-ray emission from the vicinity of this object could be affected by the disk mod- ulation which makes it appear interesting to search for this tentative 110 days period also in the γ-ray regime.
3.1.2 Short and weak Flares during Infrared and X-Ray Ob- servations
During the last 10 years a large number of weak and short flares have been reported for Sgr A* and their frequency has been estimated to reach up to several flares per day. For the infrared (IR) band the Adaptive Optics Imager NACO 9, which is part of the Very Large Telescope (VLT) in Chile [42], among others, repeatedly reported flares. Further-
