Analyzing the High Energy Activity of Candidate Neutrino Emitter Blazars to Constrain their Observability through Deep Sea/Ice Cherenkov Telescopes

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Dipartimento di Fisica “E. Fermi”

Analyzing the High Energy Activity of

Candidate Neutrino Emitter Blazars to

Constrain their Observability through Deep

Sea/Ice Cherenkov Telescopes

ANKUR SHARMA

April 2020

Supervisor:

Dr. Antonio Marinelli

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Dissertation submitted to the Department of Physics of University of Pisa, Italy

for the degree of

Doctor of Philosophy (Ph.D.) in Physics

XXXII Cycle

Academic Year 2016-19

Jury Members:

Dr. Damien Dornic

Dr. Antonio Stamerra

Dr. Dario Grasso

Referees:

Prof. Soebur Razzaque

Dr. Carla Distefano

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A billion neutrinos go swimming in heavy water;

....one gets wet.

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Acknowledgements

This thesis is a result of the direct and indirect contributions of a number of people and in-stitutions, the contribution of all of whom has been monumental at different stages. I would therefore like to apologize in advance for any names that I omit or forget to mention, as it makes their contribution second to none.

At the very outset, I would like to thank my supervisor Dr. Antonio Marinelli. This journey and this thesis would not have been possible without your constant support. You have borne with my shortcomings day in and day out for three years straight and been my go-to person for all manner of problems. You were like a guardian and a friend whose presence made me feel much more comfortable finding my way in a rigid system with an alien language. Thanks a million for getting me through the labyrinth that the Italian bureaucracy is. It has been a great pleasure knowing you as a person and working under your guidance. I hope I achieved half of what you wanted me to when we began, and look forward to continuing our collaboration.

I would also like to recognize the support of all my colleagues and collaborators. Within the KM3NeT experiment, I would like to thank the members of the calibration, multi-messenger and high-energy astrophysics groups, especillay Dr. Agustin Sanchez Losa, Dr. Rosa Coniglione and Dr. Damien Dornic for their valuable insights. A special mention to Mr. Nafis R. K. Chowdhury for his constant flow of ideas on things to collaborate on and cooperation with the calibration work. Outside KM3NeT, I feel I owe sincere gratitude to Dr. Jose Rodrigo

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Sacahui Reyes for more than one reason. My apologies for constantly bothering you with the multitudes of light curves. Our skype calls were really interesting. I hope to continue our collaboration and meet you in person someday soon. My gratitude also to collaborator and fellow student Ms. Mabel Osorio for taking time to do the painstaking work even at a moment’s notice.

I am grateful to my reviewers for patiently going through my thesis and for their invaluable suggestions to bring out the best in my work. I would also like to acknowledge the support of Dr. Dario Grasso, my internal academic referee, for his outside suggestions and for the strong stands he took to ensure a smooth and timely completion of my work and dissertation. My colleagues in Pisa were also very kind and helpful throughout. Especially, our PhD represen-tative Francesco Di Renzo, who was happy to go through my long and sometimes resentful emails and provide solutions to all manner of problems. I wish you all the very best for your dissertaton and for your life.

My stay in Pisa was made much more enjoyable by my close group of friends, all of whom enriched my life in one way or another. Thanks a lot to Sabbir Ahmed, Shubhi Parolia, Su-vankar Roy Chowdhury, Anirban Ain, Rajesh Singh, Manoj Saini and a few others, for allowing me to couchsurf, for the late night walks, for the free dinners, for the long and competitive game nights and much more. I will cherish our memories for a long time.

My girlfriend was a constant pillar of support throughout this journey and I owe her a heartfelt gratitude for standing by me in the tough times and the good ones. Last but not the least, I would like to thank my parents and family to whom I am greatly indebted. I knew I had their unconditional support and that’s what kept my motivation high at all times.

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Introduction

The last few decades have seen great progress in our understanding of the high-energy Universe, with the advent of more powerful telescopes and great advances in software and simulations. A more recent development has been that of a multi-messenger approach towards tackling the not yet understood or observed sources and phenomena, combining information from various messengers, particularly that of electromagnetic (EM) waves, very high energy (VHE) neutri-nos, and very recently, also gravitational waves (GW). Yet, many facets of the high-energy Universe remain poorly understood. Of particular interest here is the acceleration of particles upto the extreme energies of 1020 _{eV. Cosmic ray protons have been observed with energies}

upto 3 × 1020 _{eV. But these cosmic accelerators haven’t been directly confirmed yet.}

Out of the three probes at our disposal to explore the high-energy Universe (γ-rays, cosmic rays and VHE neutrinos), photons with energy > O(102₋₁₀3₎_{GeV get absorbed by the}

Extra-galactic Background Light (EBL) and Cosmic Microwave Background (CMB), charged cosmic rays are subject to deflection by the Galactic and extra-galactic magnetic fields (EGMF) and also interact with the CMB at energies above ∼ 1020_{eV to produce neutrinos. This limits their}

horizon to ∼ 200 Mpc [1]. Neutrinos however, have a much longer mean free path than γ-ray photons and cosmic rays due to their charge neutral and weakly interacting nature and can be exploited to partially solve the problem of particle acceleration to VHE energies in cosmic sources.

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The IceCube collaboration recently provided strong evidence for a neutrino signal of astro-physical origin [2], [3], [4]. It has also been shown that the Galactic contribution to this diffuse flux cannot exceed ∼ 10% of the all sky astrophysical excess [5], [6]. Most of this Galactic emission comes from the interaction of diffuse cosmic rays with dense molecular clouds, and the galactic point sources (like SuperNova Remnants, Microquasars etc.) contribute only a small fraction [7]. The SED of the muonic neutrino events from the Northern hemisphere is found to have a spectrum with a harder index (α ∼2.2) than the SED of the full sky HESE (High Energy Starting Events) events observed in 7 years [8]. It is also not yet clear which class of extra-galactic population contributes more to the remaining ∼90% of the observed astrophysical neutrino flux, with GRBs ([9]) having already been excluded as dominant con-tributors and limits existing on the contribution from Star Forming Galaxies (SFGs) ([10]).

Finding electromagnetic (EM) counterparts to the IceCube astrophysical neutrino signal can be a way to answer this question. IceCube has an alert program [11], where EHE and HESE neutrino events are reconstructed in real time, and those tagged with a high probability of being astrophysical in origin and with a small angular error, are announced for follow-up to the other astrophysical observatories around the world, covering the entire EM spectrum from radio to gamma. More than 30 such alerts have been sent out to date, although the first positive follow-up was received only very recently. On September 22, 2017, the IceCube col-laboration announced an alert corresponding to the detection of a muon neutrino (E ' 290 TeV) coming from the direction of the blazar TXS 0506+056. Follow-up observations by Fermi-LAT confirmed that the blazar showed increased gamma-ray activity (flaring activity) during the detection of the IceCube neutrino event. MAGIC telescope also observed the source in an increased flux state during a six day period following the detection of the Extreme High energy (EHE) event by IceCube. Further analysis with unblinded IceCube data revealed an excess of neutrinos from the direction of this blazar in a 3 month window in 2014-15 [12]. All these observations suggest a possibility of hadronic emission from the blazar TXS 0506+056, located at a redshift z = 0.3365.

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Earth. These jets can have bulk Lorentz factors of as high as 40 [13], and very high magnetic fields. Since speed is boosted in the direction of the jet due to relativistic beaming, we can expect to see particles accelerated to very high energies. Their rapid variability is also sugges-tive of emission from a compact causally connected region. Combined with high luminosity of these sources and the fact that their jets point towards us, makes them interesting candidates to look for signatures of VHE neutrinos.

In most cases, the SED of blazars are very well described by leptonic emission. The ab-sorption of the gamma-ray flux at high energies introduces an uncertainty that does not rule out the possibility of entirely lepton-dominated gamma-ray production, but also allows for the possibility of contribution to the SEDs from hadronic channels. Neutrino flux compared to the gamma-ray component is quite low and only the most optimistic models allow for a compara-ble flux between the two. The contribution of blazars to the TeV - PeV scale diffuse neutrino flux has been constrained [14]. Thus, it is possible that a few more years of observation by IceCube can result in the identification of hotspots in the sky which are ultimately resolved to find blazars. The addition of observatories like KM3NeT and Baikal-GVD will also help to increase the significance of multi-messenger observations in the near future. Thus, the case of multi-messenger studies of point-sources of neutrinos remains strong and blazars continue to be promising candidates.

Multi-messenger studies with neutrinos encounter a few difficulties. The first is the lack of statistics when looking at high energies (E > 100 TeV). Secondly, a lack of neutrino events correlated in time with other messengers, like the case of TXS 0506+056 where a neutrino event was observed in temporal coincidence with the flaring activity in γ-rays, makes it dif-ficult to identify the prospective sources of interest. Even if a potential correlation can be established between a neutrino event and EM data from a source in one wavelength-band, contemporaneous data on the source in other wavelengths is not easily available. There is also a lack of long-term monitoring of the sources in any one wavelength-band. It is thus of fundamental interest to find sources that can be correlated with high energy neutrino events, and an experiment that performs continuous monitoring of these sources, especially in the

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X-ray and γ-ray bands, since blazars emit most of their power in γ-rays.

This work focuses on multi-messenger analysis of blazars by searching for spatial correla-tions between the blazars from the Fermi-LAT catalogs (3FGL [15] and 3FHL [16]) and the IceCube neutrino events that are of astrophysical origin. The activity of the sources found in spatial coincidence with the neutrino events, is analyzed in γ-rays to determine their potential for very high-energy (VHE) neutrino emission. The Fermi Large Area Telescope scans the full sky every three hours and provides a continuous dataset in the energy range of 0.1 - 300 GeV. Long term data (∼ 10 years) from this telescope has been utilized for the γ-ray analysis. The

γ-ray data has also been used to estimate the expected neutrino emission from the selected

sources by applying a lepto-hedronic model from [17]. This model allows us to correlate the neutrinos at 10 TeV - 1 PeV to γ-rays at 1 GeV - 100 GeV, which falls in the Fermi energy range. The KM3NeT project is a new-generation km3 _{Cherenkov neutrino telescope, next in line}

to join the list of Mediterranean neutrino telescopes after ANTARES, NEMO and NESTOR. It will be divided into two components, KM3NeT/ARCA and KM3NeT/ORCA, each dedicated to neutrino astronomy and fundamental physics of the neutrino respectively. It has an im-proved design for the Digital Optical Module, where the segmented photocathode area allows for a more efficient rejection of background [18]. The proposed detector in its full configura-tion will be capable of studying astrophysical point-sources of neutrinos with a median angular resolution of upto 0.1◦ _[₁₉_{]. While it is expected to be more sensitive to the Galactic neutrino}

sources due to its location in the Northern hemisphere, its proposed sensitivity matches that of IceCube for the extra-galactic sources [20], and it will also provide a complimentary field of view to IceCube for sources in the Northern sky.

Sensitivity of the KM3NeT/ARCA component to the blazar TXS 0506+056 is also calcu-lated, and a portion is dedicated to the data analysis of two lines of KM3NeT/ARCA and the time calibration of both components of KM3NeT (KM3NeT/ARCA and KM3NeT/ORCA) using atmospheric muons. The thesis is layed out as follows:

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The key features of the 3 most important messengers: cosmic rays, gamma-rays and neutrinos are covered. Their spectrum on Earth, sources of origin and mechanism of production are discussed. The sources of high-energy neutrinos within and outside the galaxy are also described.

• Chapter 2 provides a theoretical background on AGNs. Blazars are described in detail, with the properties of their jets and their morphological structure. Their broadband SEDs and models describing them, the possibility of proton acceleration in their jets, discernable through the detection of high energy neutrinos on Earth is also considered. The major phenomenological motivations for this work, in context of the recent advances in the field of neutrino astrophysics are outlined.

• Chapter 3 gives a synopsis of the neutrino telescopes. A historical background is provided and the interaction mechanisms of neutrinos and their detection principles are described. An overview is provided on the current and upcoming neutrino observatories in the world. Important results and milestones in the last few years are discussed.

• In Chapter 4, the upcoming KM3NeT experiment is described in greater detail along with its two component detectors: KM3NeT/ARCA and KM3NeT/ORCA. Latest updates and physics results from both the detectors are presented. The second part of the chapter deals with the results of the data analysis of KM3NeT/ARCA, and time calibration studies performed for both KM3NeT/ARCA and KM3NeT/ORCA, carried out for this thesis.

• Chapter 5 deals with the analysis of a small sample of blazars selected for their spatial coincidence with neutrino events that are likely to be astrophysical in origin. The chapter details the criteria for selection of neutrino events and the blazar sources from different catalogs. The bulk of the chapter deals with the gamma-ray analysis of the sources within the sample to understand their potential for high energy neutrino emission. Their SEDs are also considered. The last part of the chapter looks at the expected neutrino emission from a select few blazars of the sample, obtained by applying lepto-hadronic models to the gamma-ray data. Through this excercise, constraints are obtained on the

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duration of flaring activity in gamma, for the sources to be observable by the neutrino observatories on Earth.

• Chapter 6 is dedicated to the sensitivity studies of the KM3NeT/ARCA telescope to one particular blazar from the sample: TXS 0506+056. This blazar is interesting due to its recent association with a very high energy neutrino event in its gamma-ray flaring state by IceCube. Monte Carlo (MC) simulations have been used to calculate the time-dependent sensitivity for this source during its gamma-ray flare using varying time bins and different models to estimate the expected neutrino flux. A study of the differential sensitivity of the KM3NeT/ARCA detector to the direction of this source is also presented.

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Multi-Messenger Astronomy

Multi-messenger astronomy proposes to derive information from more than one messenger (particle/wave) to increase our understanding of the various phenomena of the high-energy Universe. It promises to answer some of the long sought-after questions in astrophysics by combining resources and adding up the information available through each individual compo-nent to increase the observational significance or utilizing the information obtained through one component to perform further investigations with another component in the manner of a target of opportunity (ToO).

While the field itself is relatively new, most of the messengers used in multi-messenger astronomy have a long history. This chapter describes the individual components of multi-messenger astronomy in some detail, glossing over their origin, flux strength, detection on Earth, and their role within a larger context. Neutrinos are given special attention, and the possible sources of very high-energy neutrinos are discussed.

1.1 Cosmic Rays

Cosmic rays were first discovered more than a century ago in 1912 by Victor Hess in his balloon borne experiment. But even after a 100 years of observations and investigations, their origin, especially at the highest energies, is not fully understood.

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1.1.1 Spectrum

Cosmic rays are charged particles (mostly protons, but also e−_{, e}+ _{and atomic nuclei) that}

impinge isotropically on the Earth from various sources inside and outside the galaxy. Cosmic ray spectrum has been measured to a great accuracy and has revealed some fascinating infor-mation about their origin. As seen in Fig. 1.1, the spectrum extends from about 1 GeV to almost 1021_{eV, roughly 12 orders of magnitude. Below 1 GeV, the cosmic ray flux is driven by}

solar modulation, and above ∼ 5 × 1019 _{eV, a cutoff is expected due to the absorption of the}

cosmic ray protons on the CMB photons (GZK cutoff) [21,22]. The Greizen-Zatsepin-Kuzmin limit (or GZK limit) allows to define a cosmic ray horizon of about 50 Mega-parsecs (Mpc), from beyond which, cosmic rays above the cutoff energy should not be visible.

The global spectrum of cosmic rays shows two kinks. The spectrum transitions from a solar wind dominated one upto about ∼ 5 GeV into one dominated by cosmic rays from galactic sources like SuperNovae (SNe). Upto an energy of ∼ 3 PeV, the spectrum can be defined by a power law, with E−2.7_{. At the break, known as the knee, the spectral shape changes}

to E−3_{. The knee is proposed to be the characteristic of the limit of galactic confinement}

([23], [24]). It is accompanied by a change in the composition, with an increase in the frac-tion of heavier nuclei [25]. The origin of the flux between ∼ 1015_{− 10}19 _{eV is not clearly}

understood, although it is believed to be dominated by exra-galactic sources. The second break, or the so-called ankle, is observed at energies of ∼ 1018_{eV, where the spectrum softens}

again to E−2.7_{. The ankle signals a flux of extra-galactic origin. The spectrum falls of at}

∼ 1011 _{GeV, which is expected due to the GZK cutoff, but the uncertainties in measurements}

due to the low flux at that energy do not allow a conclusive establishment of the cutoff. It is still possible that the spectrum extends beyond these energies but higher energy particles have not yet been observed due to their extremely low flux or limitations of the detection facil-ities. Cosmic rays above ∼ 1018 _{eV are referred to as ultra-high energy cosmic rays (UHECRs).}

An anisotropy has also been measured in the arrival directions of cosmic rays by IceCube [27] and other experiments, which shows a spatial dipole. The effect is energy dependent,

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1.1.2 Origin and Sources

Figure 1.1: The spectrum of cosmic rays from various experiments. Figure from [26].

but is seen from energies of a few TeVs, all the way upto a few EeVs. The reason behind the anisotropic distribution of flux is not yet understood.

1.1.2 Origin and Sources

As previously stated, the origin of a large part of the cosmic ray spectrum is unclear. It is also possible that the fluxes at different energies are contributed by different classes of sources. While cosmic rays detected on Earth upto a few GeV originate from the Sun, the energy range of 1010_{− 10}15 _{eV is generally assumed to be dominated by galactic cosmic rays, of which a}

major share is from the cosmic rays accelerated at the shocks of the SuperNova Remnants (SNRs) in the galaxy.

Below the knee, the contribution is still believed to be mostly galactic, however, the sources that can accelerate particles upto PeV energies within the galaxy are not clearly understood.

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Apart from galactic supernovae, possible sources include pulsar jets, microquasars, galactic center and the Central Molecular Zone (CMZ). Around the ankle region and beyond, the flux should be dominated by extra-galactic sources, with evidence suggesting AGNs and their jets as the plausible acceleration sites for these UHECRs [28]. Gamma-ray bursts (GRBs) can also possibly account for a fraction of the extra-galactic cosmic ray flux.

The power law spectrum of cosmic rays indicates a non-thermal process behind the ac-celeration of cosmic ray particles. The shock fronts of supernova remnants (SNRs) and the jets of AGN are believed to be the sites of acceleration of protons (and electrons and other nuclei) in a process known as first order Fermi acceleration. In this process, acceleration occurs at the boundary of two colliding plasmas, where the particles are trapped due to magnetic confinement and repeated elastic scatterings on both sides of the shock front lead to a gain in energy [29]. Alternatively, acceleration of particles can also be achieved through deflections in molecular clouds through the process of second order Fermi acceleration [30, 31]. The global spectrum of observed cosmic rays differs from the predicted value of E−2 _{at the source}

due to propagation effects like the modulation due to the irregular component of the galactic magnetic field, and interactions with the interstellar gas [32], and also due to energy losses through radiative processes.

Fermi acceleration is effective only until the time required for the energy gain through acceleration is less than the energy loss time-scale, or until the scattering length is less than the shock radius. The maximum energy gain by a particle under this process is subject to the confinement time within the shock. Since the confinement time depends directly on the dimensions of the source and the strength of the magnetic field, the simple criterion proposed by Hillas [33], which essentially conveys that acceleration can progress only until the gyro-radius does not exceed the size of the accelerator, can be used to determine the maximum energy attained by a charged particle within a shock:

Emax< ZeBR ≈ Zβ R kpc ! B 10−6_G × 1018_eV _(1.1)

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where B is the magnetic field, Ze is the particle charge, R is the radius of the accelerator, and β = v

c. Eqn. 1.1 shows that for larger accelerating sources, lower magnetic fields are required. The possible acceleration sites can be represented through a Hillas plot (Fig. 1.2) which is a graph between the size R of the source and the magnetic field B. It shows the various source classes that can be responsible for the acceleration of particles in various ranges of size-magnetic field strength. As seen from Fig. 1.2, GRBs, AGNs and neutron stars are all capable of producing a 1020 _{eV proton, while SNRs will not be able to accelerate protons}

beyond PeV energies.

Figure 1.2: The Hillas plot showing the various source classes capable of cosmic ray acceleration. Figure from [34].

The theory of Fermi acceleration is based on empirical evidence. Although, there is no conclusive proof as yet to establish this mechanism as the principal cause of acceleration in astrophysical sources. Neutrinos, which are unaffected by propogation effects, due to their charge neutrality and weakly interacting nature, can provide a useful insight into the proton

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1.1.3 Detection

spectrum at the source, since they are produced in the decays of charged pions, which them-selves are a result of the interactions between the accelerated protons and target photons (also neutrons and protons).

1.1.3 Detection

The flux of cosmic rays decreases rapidly with energy (dN dE ∝ E

−α_{, α = 2.7, 3). From ∼ 10}4

particles per sq. meter per second at 1 GeV, it falls to only a few particles per sq. meter per year at 1015 _{eV. At the highest energies (E > 10}18 _{eV), only 1 particle is expected per square}

meter per century. This vast difference in flux with energy implies that different techniques are required to detect cosmic rays of different energies.

At the lower energies, balloon borne and satellite experiments like the AMS [35] are suffi-cient to carry out the task. They can detect the type of particle and its momentum. Above 1 TeV, the surface area of these class of experiments are no more sufficient to collect the low flux. Additionally, longer exposure times are required. This is achieved by ground-based experiments which can cover several km2 _{of area and allow for the long exposures required}

to measure the low fluxes. Pierre Auger Observatory in the recent past [36], and the current Telescope Array (TA) experiment [37] are prime examples of such experiments. The detection techniques employed by these experiments include the detection of the air showers produced by the high energy cosmic rays interacting in the atmosphere, and the measurement of the flu-orescence produced in the air by charged particles exciting nitrogen atoms in the atmosphere. The ground-based experiments cannot identify the type of the primary particle based on these techniques.

Cosmic rays lose directional information quickly since they are deflected by the galactic and inter-galactic magnetic fields. But for protons with E > 5 × 1018 _{eV, the gyro-radius}

in the galactic magnetic field of B ' 3µG comes out to be of the order of the dimensions of the galactic disk (' 15 kpc). These protons can therefore escape the galaxy travelling in straight lines. It should be noted that the above relation does not hold for heavier nuclei since the gyro-radius depends on the charge Z of the particle, and heavier nuclei will be deflected

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1.2. GAMMA-RAY PHOTONS

more than protons by the same magnetic field. Thus, proton astronomy can be feasible for the energy range of E > 1018.5−19 _{eV, where the deviation due to inter-galactic magnetic}

field (IGMF) might be small enough that protons can be traced back to their sources, atleast at low redshifts. Practically, there is still some way to go before detectors that can collect statistically significant flux at those energies are constructed. However, another constraint on the observation of UHECRs comes from the GZK cutoff, which limits the portion of the Universe that can be explored with cosmic rays. The mean free path of the cosmic ray particle before its interaction with CMB photons is restricted to a few hundred Mpc [1].

1.2 Gamma-ray photons

Gamma-rays (γ) refer to a form of electromagnetic radiation with energy higher than ∼ 100 keV. They are produced in many processes including radioactive decays here on Earth, and EM cascades in the atmosphere generated by cosmic rays. The Sun is another source of gamma radiation on Earth. But gamma-rays can also originate outside the Solar Sytem, in a variety of thermal and non-thermal processes. Their production mechanisms include brehmsstrahlung of relativistic electrons, e+_e−_{pair annihilation, nuclear decays, synchrotron radiation of relativistic}

electrons (and protons), inverse Compton scattering, and hadronic processes that produce π0

which decays into two gamma-ray photons. Dark matter annihilation is also speculated to produce gamma-rays.

1.2.1 Leptonic production

The Fermi acceleration process is applicable to all charged particles including electrons and positrons. But due to their low mass, electrons lose their energy very rapidly via synchrotron radiation in the magnetic field of the source. Additionally, the synchrotron photons, if dense enough, can get upscattered by these relativistic electrons, further constributing to their energy loss. Thus, the radiative losses of the electron population can contribute to the energy content of the gamma-rays, giving rise to a leptonically generated gamma-ray flux. The mechanism described above, known as the Synchrotron Self-Compton (SSC) mechanism, is the natural explanation for the spectral energy distributions (SEDs) of most of the AGNs. It is described

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1.2.1 Leptonic production in detail in section 2.2.2.

Assuming a power law spectrum for the accelerated parent electrons with index αe, the slope of the synchrotron radiation curve should also display a power law behavior [38]. This gives: νFν = Eγ2 dNγ dEγ ∝ E(αe+1)/2 γ (1.2)

The spectrum is also dependent on the magnetic field strength B as:

E_γ2dNγ dEγ

∝ B(−αe+1)/2 (1.3)

The synchrotron self-absorption process determines the left slope of the synchrotron curve, as the self-absorbed flux increases with decreasing energy of electrons. Thus, we get the familiar expression:

E_γ2dNγ dEγ

∝ E_γ5/2 (1.4)

The peak frequency of the distribution, known as the synchrotron peak frequency (νs peak), is the turning point between the synchrotron self-absorption and the synchrotron spectrum itself. Synchrotron photons can range in energy between ∼ 10−7 _{eV, reaching as much as 10}9

eV in the extreme cases. However, for the observed AGNs, the distribution of Synchrotron emission peaks mostly in the eV - keV range. More details on the synchrotron radiation pro-cess, the power emitted, the single particle spectrum and the photon distribution can be found in [29] and [31].

Additionally, the co-existence of a population of electrons and the synchrotron photons produced by these parent electrons, gives rise to another interaction probability: the inverse Compton (IC) scattering of the synchrotron photons on the relativistic electrons. The energy gain of photons depends upon the interaction cross-section, which varies from the Thomson to the Klein-Nishina regime. The isotropic radiation field of the synchrotron source appears fairly anisotropic to the ultra-relativistic electrons (γ 1). Thomson scattering of this highly

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1.2.2 Hadronic production

anisotropic radiation systematically reduces the electron kinetic energy and converts it into inverse-Compton radiation.

The Thomson cross-section is applicable in the limit of hνΓ mec2 (where hν is the photon’s energy, and Γ is it’s Lorentz factor), and yields the following photon distribution:

νFν = Eγ2

dNγ

dEγ

∝ E(αe+1)/2

γ (1.5)

In the Klein-Nishina regime (hνΓ mec2), this changes to:

νFν = Eγ2

dNγ

dEγ

∝ Eαe

γ ln(Eγ) (1.6)

The final spectrum obtained for IC scattering from a quasi-thermal population of electrons (“thermal Comptonization”) [29] has contributions from many orders of scattering of many particles, and the information of the particle distribution is lost in the final spectrum.

1.2.2 Hadronic production

Astrophysical sources like SNRs, microquasars, and AGNs have an ambient field of photons generated through several self-balancing processes. The presence of dust particles inside or surrounding these environments is also widely accepted. These ambient photons and dust particles can act as targets for the relativistic hadrons co-accelerated with the electrons by these sources. There are two possible reaction mechanisms that take place with varying probability depending upon the matter content of the source and the environment surrounding the source. The first one, known as the astrophysical beam dump process [39], occurs when two hadrons interact to produce mesons and more nucleons. The energy of the hadron beam is dumped into these new particles, which themselves decay and transfer their energy to the daughter particles, mostly γ and ν:

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1.2.2 Hadronic production

The second mechanism is the photo-hadronic interactions (also called photo-pion interac-tions), and occurs when the hadrons interact with the ambient photon field or the photons of the CMB or (EBL) to produce charged and neutral π, which again decay into γ and ν. The process is intermediated by a ∆+ _{resonance, and is also responsible for the GZK limit:}

p + γ → ∆+ → p + π0 (1.8)

→ n + π+ (1.9)

The π0 decays into 2 γ, and is responsible for the hadronic gamma-ray flux. The energy

threshold for this interaction, in the center of mass frame, is roughly twice the mass of the

π0, i.e. ' 300 MeV. On average 10% of the parent proton’s energy is transferred to the γ.

The decay of the charged π into neutrinos is covered in Section1.3.1.

The energy range of the γ-rays produced through this process can be estimated based on the energetics of the reaction [40]. In the rest frame of the π0, the energy of the 2 γ will be E_γ∗ = 67.5MeV (c = 1), half of the rest-mass of the π0_{. The distribution of the photons will}

also be isotropic, which in terms of cos θ∗ _gives: dN

d cos θ∗ =

1

2 (1.10)

where cos θ∗ _{is the angle that the photon makes with the direction of the π}0_{’s momentum,}

and the factor of 1

2 is obtained from normalization. The energy of the photons in the observer’s

frame is readily obtained by Lorentz transformation of the photon’s energy in the π0_rest-frame,

and using the relation E∗

γ = p ∗ γ for photons: Eγ = Eγ∗Γ + p ∗ γβΓ cos θ ∗ _(1.11) = mπ0 2 Γ (1 + β cos θ ∗ ) (1.12)

where β and Γ are the velocity and the Lorentz factor of the π0 _{respectively. A decaying} π0 can thus produce γ in the energy range 0 for a π0 at rest to Eπ0, for a π0 with Γ 1.

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1.2.2 Hadronic production dEγ = mπ0Γβ 2 d cos θ ∗ _(1.13) dN dEγ = dN d cos θ∗ d cos θ∗ dEγ (1.14) = 1 2 2 mπ0Γβ (1.15) = 1 mπ0Γβ (1.16)

The spectrum of the primary protons in an astrophysical environment that interact to pro-duce the π0_{s, dictates the spectrum of the γ-rays from the source. The spectral index for the} γ-rays produced in a source with proton spectral index αp = −2is also -2. The lower threshold of the SED is determined by the π0_{s at rest, while the upper bound is set by the maximum}

energies that the protons can reach in the astrophysical accelerator. The spectrum observed on Earth is also suppressed at higher energies by propagation effects which are discussed in the next section.

Evidence for hadronically produced γ-rays from astrophysical sources is scarce. The SED of the galactic SNR IC 443, in the Gemini constellation, was analyzed by the Fermi collaboration in 2013, and the data showed a better fit assuming a hadronic origin for γ-rays (see Fig. 1.3). Another SNR, W44, also showed a better fit to its SED with the so-called π0_{-bump, as}

com-pared to the leptonic processes [41]. To disentangle the origin of the γ-ray spectrum of sources like SNRs, microquasars and AGNs, a fit to the global SED that is in agreement with data at all energies is required. It can be obtained by modelling the SED based on assumptions about the parameters of the source (like size of the emitting region, source populations, distance of the source and magnetic field strength) and various processes involved in the production of the EM spectrum from the source, including brehmsstrahlung, Bethe-Heitler pair production, synchrotron radiation from the primary electrons and protons, and the cascades of π decay products, inverse Compton scattering of synchrotron photons, direct γ-rays from π0 _{decay and}

a few other minor contributions. This requires a comprehensive and intensive algorithm that simulates all these processes individually and takes into account the self-interactions as well. There are tools available online that can achieve this to some degree [42, 43, 44]. Of these,

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1.2.3 Extra-galactic Background Light and Gamma-ray propagation

[44] is the oldest one that can simulate these radiative processes, while [42] can be used to fit the SED data with a few input parameters.

Figure 1.3: The SED of galactic SNR IC 443, fit using Fermi data. The best-fit is obtained assuming a

π0_{-decay origin of the data. The grey region is the systematic uncertainty due to the diffuse flux estimation.}

Figure taken from [41].

1.2.3 Extra-galactic Background Light and Gamma-ray propagation

The extra-galactic background light refers to the integrated luminosity of the Universe in the UV, optical and IR and is the radiation accumulated over several epochs from galaxy forma-tion processes that has redshifted to longer wavelengths over time. It permeates the Universe uniformly [45] and forms the second most energetic diffuse background in the Universe after the CMB [46].

The direct detection of EBL is inhibited by the contribution from zodiacal light, which is orders of magnitude higher than EBL in luminosity. Measurements are available in the optical and near-infrared (IR) bands, but they are in contradiction with the data from VHE observations [47, 48]. An approach to understand EBL is thus the modelling of the EBL to

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1.2.3 Extra-galactic Background Light and Gamma-ray propagation

obtain its spectrum that can be verified through observations. These approaches fall under 2 major categories [46]:

1. Forward evolution, that assumes some initial conditions for the Universe and evolves them forward in time to obtain the luminosity and distribution at present epoch, taking into account some galaxy formation models. [49, 50] fall under this category.

2. Backward evolution, where the current galaxy population estimates are utlilized to ex-trapolate backwards in cosmic time. Examples include [51, 52].

Other empirical approaches utilize observational quantities like the star formation rate (SFR) or the energy density of the Universe to infer the galaxy evolution spread over a redshift range ([53, 54]) or the direct observations of galaxy evolution in the range of redshifts most affecting the EBL [46]. Of all these, the model from [51] is a comprehensive and empirical one, and has been used to account for the EBL absorption effect on γ-rays throughout this work.

The high-energy γ-ray photons from cosmic sources can scatter on the EBL photons through the Bethe-Heitler process [38] dampening the flux from these sources at high energies.

2γ → e++ e− (1.17)

The cross-section is dependent on the energy of the incoming photon, and the threshold energy for the interaction is determined by the mass of the electron. Moreover, the photon number density of the EBL also evolves with redshift. EBL aborption is thus a function of energy and redshift. The attenuation of the γ-ray flux observed on Earth due to EBL absorption can be described through an exponential multiplicative factor:

Fintr = e−τ (z,Eγ)Fobs (1.18)

where Fintr is the intrinsic flux from the source, Fobs is the flux from the source observed on Earth, and τ(z, Eγ)is the photon-photon opacity due to EBL as a function of the redshift z, and photon energy E . The τ factor is obtained from the models and determines the amount

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1.2.3 Extra-galactic Background Light and Gamma-ray propagation of absorbed flux.

Figure 1.4: Energy dependent cosmic horizons for γ-rays and UHECR protons. The scattering on CMB photons at 1020_{eV limits the cosmic ray horizon to 10 Mpc, while the EBL photons limit the observable sky in γ to a}

few Mpc above energies of a hundred TeV. Plot taken from [55].

The attenuation of the flux from astrophysical sources due to EBL absorption creates a horizon for γ-rays, similar to the GZK cutoff for cosmic rays (Fig. 1.4). For a source at

z ∼ 1, the observed flux becomes 0.1 times the intrinsic flux when we consider energies of

Eγ < 200 GeV. For Eγ ∼ 300 TeV, the energy where the interaction cross-section with CMB

photons becomes larger that the interaction with EBL, the horizon is as small as 1 Mpc. This implies that the fraction of the Universe that can be explored with high-energy γ-rays is also restricted. Beyond a few hundred Mpc, the pevatrons (sources emitting at PeV energies) and the UHECR emitters in the Universe can be accessible through a particle/messenger that does not get easily absorbed or deflected during propagation. This is where astrophysical neutrinos can play a distinctive role.

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1.2.4 Detection

While the observation of γ-rays started already in 1961 using Explorer XI [56], the field of γ-ray astronomy has only blossomed in the last 20-30 years when detectors began to be available to detect the GeV γ-rays, and the resolution became good enough to resolve individual sources. Today however, γ-rays upto TeV energy and above are regularly detected by telescopes around the world.

Presently, γ-ray astronomy is practised using two classes of detectors: space-based and ground based. The Earth’s atmosphere is not transparent to γ-rays, hence the direct detection of γ-rays is only possible through space-based telescopes. Since Explorer 11, other attempts were made to study the γ-ray universe from space in the last century, but the most remark-able among them was the EGRET instrument aboard the Compton Gamma Ray Observatory (CGRO). It was able to detect 270 individual sources between 100 MeV to 10 GeV and con-tributed to the most extensive γ-ray catalog at the time.

The Fermi Gamma-Ray Space Telescope, launched in 2008, hosts two detectors onboard, the Large Area Telescope (LAT) [57] and the Gamma-ray Burst Monitor (GBM) [58]. Fermi-LAT is a wide field of view (FOV), pair conversion telescope sensitive within the range of 100 MeV to ∼ 500 GeV. A pair conversion telescope converts the incident γ-ray photon into an

e+e− pair in the converter, which is tracked by the tracker to trace the direction and energy

of the primary photon, and the calorimeter absorbs the leptons to provide a calorific mea-surement of the energy. The convertor shield provides a veto against the background from charged cosmic rays. With a FOV of ∼ 2 steradian, the Fermi-LAT can scan the whole sky in close to 3 hours and provides an uninterrupted and comprehensive monitoring of the γ-ray sky. It also has a high angular resolution (< 1◦_{) for E > 100 GeV, and can be the ideal tool used}

to study the long term behavior of astrophysical sources in the GeV band. The 4FGL catalog from Fermi-LAT, with 8 years of exposure, contains over 5000 γ-ray sources, galactic and extra-galactic, as compared to the 3FGL catalog [15], which had slightly over 3000 sources (Fig. 1.5). Over 3000 out of these 5000 are AGN. The GBM instrument on the other hand,

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1.2.4 Detection

is complimentary to the LAT and is mainly designed to study transient events like GRBs. It has a sensitivity within the X-ray and γ-ray band (8 keV - 40 MeV), and has identified over 1400 individual GRBs, listed in the third Fermi GBM catalog [59].

Figure 1.5: The skymap of all sources within the Fermi-LAT 3FGL catalog from [15].

The surface areas of space-based telescopes are not sufficient to detect any meaningful flux of γ-rays above few hundreds of GeV. For this purpose, ground-based telescopes are used, which implement several indirect detection methods. While the γ-rays do not themselves penetrate the atmosphere, like charged cosmic rays, they interact with the nuclei in the atmoshphere to generate cascades of secondaries (EAS), which can reach the ground. This principle is exploited in two manners. The Cherenkov light produced by the whole EAS spreads out as it reaches the ground, forming a cone with radius 100-120 m on the ground. This light can be detected by an array of telescopes spread over a large surface area on the ground where each records the Cherenkov light from the shower in parts and reconstructs the whole shower event to deduce its properties (see Fig. 1.6). These are known as Imaging Atmospheric Cherenkov Telescopes (IACTs). The current generation of IACTs like MAGIC [60] (preceeded by HEGRA

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1.2.4 Detection

[61]), HESS [62] and VERITAS [63] cover the northern and southern sky, and have detected over 200 galactic and extragalactic sources with an energy sensitivity between ∼ 10 GeV -10 TeV [64]. The upcoming CTA experiment (Cherenkov Telescope Array) [65] will raise this threshold by an order of magnitude to 100 TeV.

Figure 1.6: A schematic of the detection techniques of Imaging Cherenkov (IACT) telescopes (left) and Water Cherenkov (WCD) telescopes (right) from [66].

Another possibility to reconstruct the CR or γ-ray primary is through the detection of the secondary particles of the EAS. Since the shower can be much more spread out, this requires detector arrays spread over very large areas (O(km2₎_{). Since these detectors require to be}

closer to the shower maximum, they are usually deployed at higher altitudes to increase ef-ficiency. The secondaries are detected either by using scintillators, that detect scintillation light through PMTs, or by Water Cherenkov Detectors (WCDs). WCDs use arrays of water tanks lined with PMTs that detect the Cherenkov light produced by the shower secondaries traversing through water (Fig. 1.6). In both the techniques, the final reconstruction is based on the timing and charge information relayed by the individual PMTs. WCDs are providing very useful results for TeV sources especially within the galaxy, like the HAWC experiment in Mexico, which recently released the > 100 TeV spectrum for the galactic γ-ray sources [67].

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1.3. ASTROPHYSICAL NEUTRINOS

to 10 TeV range, WCDs can cover the 101₋₁₀3 _{TeV band. IACTs work with a superior angular}

and energy resolution over the WCDs and are ideal for morphological and spectral studies. But they require good optical conditions and can only work on dark nights, while WCDs have a very high duty cycle and can work almost continuously without interruption, perfect for carrying out surveys. IACTs have a pointing capability (nowadays quite fast) and a narrow FOV, which can be exploited for observation of transients and Targets of Opportunity (ToO), while WCDs can be better used for studying emission from extended sources. In summary, the ground-based telescopes along with the space-based ones provide a wide array of instruments in the hands of astronomers to explore and understand the γ-ray sky, which, in the coming years is slated to provide many essential clues to our understanding of important astrophysical phenomena.

1.3 Astrophysical Neutrinos

Figure 1.7: The spectrum of neutrinos from various sources [68]. The Solar, Reactor, SuperNova, Atmospheric and Terrestrial neutrino fluxes have been measured while others represent expected values.

Neutrinos are demonstrably an essential probe to explore the farther reaches of the Universe, especially at energies above TeV. But neutrinos are produced in a variety of astrophysical

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1.3.1 Production

sources and vary in energy from a few micro-eV (10−6 _{eV) for cosmological neutrinos to}

PeV and above for neutrinos from extra-galactic sources like AGN (see Fig. 1.7). Solar neutrinos and neutrinos from SuperNova bursts have lower energies (few MeV - GeV) and require dedicated detection facilities, although indirect detection is sometimes possible with the neutrino telescopes. Source like SNRs are suspected to produce neutrinos upto PeV energies. Beyond that, extragalactic sources like AGNs, GRBs etc. are believed to be the prime candidates.

1.3.1 Production

The production of neutrinos at high energies proceeds through either one of the two processes discussed in Section 1.2.2. Charged πs are produced during the astrophysical beam dump process and/or the photo-pion interactions. These decay into muons and electrons, emitting neutrinos in the process [69]:

π+→µ+_{+ ν} µ (1.19) µ+→e+_{+ ν} µ+ νe (1.20) and π− →µ−+ νµ (1.21) µ− →e−+ νµ+ ¯νe (1.22)

A contribution from the decay of K and neutrons is also expected, which becomes in-creasingly significant with rising hadronic losses in the source. Distinction between ν and ¯ν is currently not possible, however the different flavors can be distinguished in the detectors, as is explained in Chapter 3. The flavor ratio at the source, as seen from eqn.s 1.20, 1.22 is

νe : νµ : ντ = 1 : 2 : 0.

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1.3.1 Production

equally distributed between the 4 leptons produced after the pion and muon decays [70]. Their spectral index is also correlated to the spectrum of the parent protons, but for p-γ interactions, it is dominated by the spectrum of the target photons [71].

A clear relation is evident between the cosmic rays, γ-rays and neutrinos. Astrophysical sources that are capable of producing high-energy cosmic rays by accelerating protons, are also capable of producing γ-rays and neutrinos since the processes involved in their production are similar. Together, they account for the energy output of the source. Furthermore, the spectral indices of all three are related as: αCR ∼ αγ ∼ αν, which has a theoretical expec-tation of -2 if the Fermi acceleration process is involved. Therefore, sources of high-energy

γ-rays can also be possible sources of neutrinos (assuming a hadronic production mechanism

in the source), and by observing the gamma-ray activity of the source, its neutrino spectrum can be constrained. This has been the central premise of this work. Fig. 1.8 shows the SEDs of cosmic rays, high-energy neutrinos and γ-rays measured by various experiments put together. The favor ratio at the source described above can be skewed under certain circumstances. For sources with muon dampening, i.e. the sources where the muon is absorbed before it can decay, the ratio becomes νe : νµ : ντ = 0 : 1 : 0, since one νe and one νµ is lost for every interacting µ. An additional consideration is the neutrino flux from charm decays, known as the prompt flux. Decaying charmed mesons can also produce neutrinos. Since the charmed particles are very short-lived, they do not lose energy before decaying and the neutrinos are thus produced at higher energies. While for the conventional neutrinos flux: the flux from the decay of charged pions and kaons, the energy of the neutrinos is degraded due to their relatively longer lifetime, which allows for energy loss through collisions before the decay. However, the cross-section for the charmed production mechanism is much smaller than the conventional one. After the crossing point (Eν ∼ 500 TeV), the prompt flux starts to dominate over the conventional flux (Fig. 1.9). The flavor ratio of the charmed flux is νe : νµ : ντ = 1 : 1 : 0. The charmed decays also become important in the case of atmospheric neutrinos, where a distinction has to be made between the conventional neutrino flux [75] and the prompt flux [76] , especially at higher energies where the prompt flux component is dominant.

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1.3.1 Production

Figure 1.8: Neutrino SED inferred (i) using 8 years of IceCube data (tracks, upgoing) with 1σ uncertainty range (shaded) shown with the solid red line, and (ii) using 7 years of IceCube data (HESE sample) shown in magenta. This has been compared to the gamma-ray flux of unresolved extra-galactic sources from Fermi [72] (blue data) and ultra-high-energy cosmic rays [73] (green data). The multi-messenger interfaces are highlighted: (A) The joined production of charged pions (π±_{) and neutral pions (π}0_{) in cosmic-ray interactions leads to}

the emission of neutrinos (dashed blue) and gamma rays (solid blue), respectively. (B) Cosmic ray emission models (solid green) of the most energetic cosmic rays imply a maximal flux (calorimetric limit) of neutrinos from the same sources (green dashed). (C) The same cosmic ray model predicts the emission of cosmogenic neutrinos from the collision with cosmic background photons (GZK mechanism). Figure from [74].

The above described flavor ratios are valid only near the sources. Propagation through space changes these flavor ratios, so that the ratio detected on Earth differs from the one at source. For the case of νe : νµ : ντ = 1 : 2 : 0 at source, it is expected to be roughly

νe : νµ : ντ = 1 : 1 : 1on Earth. While for the case of νe : νµ : ντ = 0 : 1 : 0at source (muon damping), it is expected to be νe : νµ: ντ = 1.6 : 0.6 : 0.8at Earth [77]. The process through which this happens is known as neutrino oscillations. It is discussed in the next section.

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1.3.2 Oscillations

Figure 1.9: The conventional atmospheric neutrino fluxes for µ, νµ and νe compared to the prompt flux

calculation for νµ+ ¯νµ based on [76]. Figure from [76].

1.3.2 Oscillations

Neutrinos come in three known flavors, each associated with the producing lepton in their charged-current interactions. They were also asuumed to have zero mass until the 1970s, when experimental evidence began to contradict this view. Today, it is well understood that neutrinos are massive particles and hence, can undergo flavor transition as they propagate through space. This phenomenon is known as neutrino oscillations.

The problem with the massless view of the neutrino surfaced when the Homestake Mine ex-periment measured 1/3rd of the expected electron neutrino flux from the Sun [78]. This came to be known as the “solar neutrino problem”. Over the decades, several other experiments measured either a deficit of electron neutrinos or an excess of tau neutrinos [79, 80, 81, 82]. The plausible way to reconcile all the empirical evidence was to accept that the neutrinos changed flavor during propagation from the source to the detector. This inherently implied revising the massless treatment of the neutrino within the standard model [83].

Not only do neutrino have mass, but they oscillate because their mass eigenstates are distinct from their flavor eigenstates. The flavor eigenstates exist as the superposition of the

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1.3.2 Oscillations mass eigenstates: |ναi = X i=1,2,3 hU_αi∗ | νii (1.23) |νii = X i=1,2,3 hUαi| ναi (1.24)

where |ναi and |νii are the flavor and mass eigenstates respectively, and Uαi is the Pon-tecorvo–Maki–Nakagawa–Sakata matrix (PMNS), analogous to the CKM matrix that deter-mines quark mixing. Apart from the flavor mixing angles it aslo contains the CP violating phase δ. The probability of oscillation from any flavor α to β is given as follows:

P (να → νβ) = X i,j [UβiUαi∗U ∗ βjUαj]eiΦij (1.25) where Φij = (m2j− m2i)L/2E, or Φij = (∆m2)L/2E, with L as the propagation distance and E as the energy of the neutrino. Eqn. 1.25 makes it obvious that oscillations cannot be possible if neutrinos do not have mass. Moreover, the phenomenon depends not on the abso-lute masses of the neutrinos, but on the squared mass difference. Consequently, the absoabso-lute masses of the neutrinos are yet unknown and so are the mass orderings, whether they follow a normal heirarchy (NH, m1 < m2 < m3) or an inverted heirarchy (IH, m3 < m1 < m2). This

is the “mass heirarchy problem” in fundamental neutrino physics.

As stated before, as a consequence of propagation over long astronomical distances, neu-trinos have sufficient time to oscillate such that the flavor ratio observed from astronomical sources on Earth is expected to be equal for all flavors, especially due to the large mixing angle θ23 [85]. Neutrino telescopes like IceCube, ANTARES and the upcoming KM3NeT can

be an important tool to study neutrino oscillations and mass heirarchy. These telescopes can differentiate between tracks created by interacting νµ and cascades created by all interacting flavors, and subsequebtly obtain the flavor ratios by comparison to Monte Carlo simulations. A scan of the likelihood profile of flavor composition on Earth, as measured by IceCube [84], is shown in Fig. 1.10.

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1.4. NEUTRINO ASTROPHYSICS

Figure 1.10: Scan of the likelihood profile of neutrino flavor composition on Earth by IceCube [84]. The x indicates the best-fit while + is the best-fit for an older analysis. The various production scenarios are indicated by the color markers: pion production (red), muon-damped (orange) and neutron decay (green).

The detection techniques and the challenges encountered in observing astrophysical neu-trinos, and the major experiments involved are covered in Chapters3 and 4.

1.4 Neutrino Astrophysics

Neutrino astrophysics involves the understanding of astrophysical phenomena by observing neutrinos from cosmic sources. Until recently, the only observed extra-terrestrial sources of neutrinos were the Sun and the SuperNova SN1987A in the Large Magellanic Clouds (LMC), whose burst neutrinos were observed by the Kamiokande II, IMB and Baksan observatories [86,87]. But the first tentative association of a blazar with a very high energy neutrino event, and an excess of neutrinos from the direction of that source in archival data was announced by IceCube and other collaborations in 2017 [28,12]. This observation has breathed new life into the prospering field of neutrino astrophysics and put blazars on the forefront of the search for sources of VHE neutrinos.

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1.4.1 Neutrinos from Galactic sources

Neutrino astrophysics involves three major energy ranges. In the energy range of ∼ 10 MeV - 100 GeV, the Sun and the burst neutrinos from Super Novae are expected to be the major sources. Additionally, this range is also the one to which neutrino oscillation and mass heirarchy studies are most sensitive, due to the dominant atmospheric neutrino flux. The energy range from roughly 100 GeV - 30 TeV is expected to be dominated by the emission from galactic sources like SuperNova Remnants. Above ∼ 30 − 50 TeV, the emission from extra-galactic sources like AGNs and GRBs is thought to be the dominating component in the neutrino SED [38].

This section provides some detail about the neutrino emission potential of the candidate galactic and extra-galactic sources. AGNs and blazars are described in the next chapter (see chapter 2).

1.4.1 Neutrinos from Galactic sources

Confinement arguments suggest that protons above PeV energy will escape from the galaxy. This puts an upper bound on the energy that the neutrinos from galactic sources can achieve, of ∼ 50 − 100 TeV [70]. Since the distances involved are significantly less than the those for extra-galactic sources, the same event rate can be expected from a much less luminous galactic source than for an extra-galactic one. Neutrino telescopes also have a high sensitivity in this energy range. The observation of Galactic point-like neutrino emission in the energy range 50-100 TeV is therefore exceedingly interesting since it will confirm the hadronic nature of the pevatrons [88].

SuperNova Remnants

SuperNova Remnants are expanding shells of hot gas and matter left over from the collapse and subsequent explosion of a massive star. The relativistic matter shocks the surrounding inter-stellar environment and these shocks are prime sites for acceleration of particles. Conse-quently, SNRs are believed to be the foremost candidates for producing galactic cosmic rays. Observations by imaging Cherenkov telescopes in the last couple of decades have provided

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strong evidence for the capability of SNRs to produce VHE γ-rays upto a few TeVs, but to establish their potential as cosmic ray accelerators upto PeV energy, observations of γ-ray flux upto 100 TeV from these sources is required. No SNR with a photon SED reaching these energies has been observed so far. Most of the galactic supernovae are old. SNRs can only produce PeV cosmic rays for the initial few hundred years of their life (in the Sedov phase), after which they cool down and do not have enough energy, or the accelerated protons diffuse outside the vicinity of the SNR and cannot be traced back to it [88]. Thereafter, their emission will be dominated by electrons and positrons.

Figure 1.11: The combined diffuse neutrino flux from SNRs and HyperNova Remnants (HNRs) with the associated uncertainty band (shaded) estimated by [89]. The data points represent the IceCube 3-year diffuse flux and the solid black line shows the best fit.

Another possibility is to trace a clear signature of π-decay in the spectrum of the SNR [90]. This has been obtained in the case of SNRs IC 443 and W44 (see Fig. 1.3). Their SED is best explained assuming a π0_{-decay component at low-energies (E < 200 MeV) [}₉₁_].

However there is no confirmation at high energies since their spectrum falls of sharply on ac-count of them being old supernovae. Within the TeVCat catalog, 15 SNRs are reported with a spectrum arriving upto the TeV band [64]. While none of them has a spectrum ranging upto 100 TeV or above, sources producing γ-rays upto ∼ 10 TeV are known. These can be useful

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to constrain the overall contribution of SNRs to the total cosmic ray spectrum. It still remains unknown however, if SNRs are responsible for cosmic rays upto the knee of the spectrum, and future imaging air Cherenkov telescopes (IACT) like CTA will see this as one of their main goals. Fig. 1.11shows the total diffuse neutrino flux estimated from stellar remains like SNRs and HyperNova Remnants (HNRs). It is compared to the 3-year diffuse flux measured by IceCube. Also, since most of the galactic sources including SNRs lie in the Southern hemisphere, the upcoming KM3NeT neutrino observatory, with an improved sensitivity in the range of 1-100 TeV, will be in a great position to observe neutrinos from these sources and settle the question of cosmic ray acceleration by SNRs.

Pulsar Wind Nebulae and MicroQuasars

PWNe are formed by strongly magnetized winds of leptons through which pulsars lose their energy. These are extended sources, and the most popular among them is the Crab nebula. While leptnic emission is sufficient to explain the spectrum of most PWNe, some PWNe have also been detected in TeV, motivating the possibility of possible hadronic emission in their environment [92].

Microquasars are a special case of X-ray binaries, where one of the components is a donor and the companion is an accretor. Matter from the donor is accreted through the accretion disk, and whenever a large concentration of matter falls in, the accretor releases this energy in the form of X-rays, often times through jets which are aligned with the axis of rotation of the disk. These relativistic jets, just like the jets of pulsars and AGNs, can be potential sites to accelerate particles. The emissions detected from microquasars are consistent with synchrotron emission of electrons, and their jets are radio bright, with emissions extending upto optical wavelengths. However, occasional TeV photons have also been detected from microquasars by HESS and MAGIC [93, 94]. Thus, the possibility of particle acceleration in these jets upto TeV energy is not ruled out. Both photo-pion and p-p interaction possibilities in microqusar jets have been considered. [95] consider the possibility of photo-mesonic production of TeV neutrinos in the jets of microqusars, while [96,97] estimate the neutrino flux from two

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sources assuming p-p interaction of non-relativistic protons from the wind of the donor and the accelerated protons from the jet of the companion. IceCube and ANTARES have placed limits on the neutrino flux from these sources of the order of 10−11_{− 10}−12 _{T eV}−1_cm−2_s−1

[98, 99].

Galactic Center & Fermi Bubbles

The galactic center and the central molecular zone (CMZ, r < 200 pc from the galactic center) are hotbeds of extreme astrophysics, having large concentration of mass and presence of hot gas with density upwards of 103 _cm−3_{. The region is already known to emit in γ-rays and is}

an important portion of the sky in indirect dark matter searches. Recently, HESS announced PeV γ-ray photons originating from the central SMBH Sag A∗ _[₁₀₀_{] and evidence for TeV}

emission from an SNR (G 0.9+0.1) close to the Center has also been found [101]. This, and the correlation of the high energy γ-ray emission from the CMZ with the integrated intensity of CS emission (a compund of Carbon-Sulphur), reported by HESS, presents a strong argument in favor of hadronic origin of these γ-rays [102]. There is also new evidence for the presence of PeV particles in the very center of the galaxy [103], and although the current acceleration rate might not be sufficient to account for any major fraction of the galactic cosmic rays, increased activity from Sag A∗ _{in the past > 10}6 years can cause a significant flux of PeV

cosmic rays and TeV neutrinos [103]. The Galactic Center and CMZ are thus very important in the estimation of total galactic neutrino flux.

Fermi bubbles are massive bubble-shaped structures, named after their discovery by the Fermi-LAT collaboration due to their strong γ-ray emission. They protrude out on both sides of the galactic disk, 50◦ _{in latitude and 40}◦ _{wide [}₁₀₄_{]. They show a hard spectrum in gamma}

(∝ E−2_{) with a cutoff at 110 (±50) GeV [}₁₀₅_{]. The low latitude bubbles seems to have a}

higher cutoff and a harder SED for γ-rays. Their origin is not yet known. The most relevant hypotheses propose a formation through past episodes of accretion into the central SMBH or a nuclear starburst within the last 10 Myr [104]. Hadronic scenarios explaining the γ-ray emission from Fermi Bubbles have also been proposed [106,107,108,109,110], but similar to the leptonic models, they fail to explain all features of the spectrum. ANTARES upper limits

Analyzing the High Energy Activity of Candidate Neutrino Emitter Blazars to Constrain their Observability through Deep Sea/Ice Cherenkov Telescopes

Dipartimento di Fisica “E. Fermi”