Investigating UV Nightglow Within the Framework of the JEM-EUSO Experiments

Investigating UV nightglow within the framework of the JEM-EUSO Experiments

Frej-Eric Salomon Emmoth

Space Engineering, master's level 2020

Luleå University of Technology Department of Computer Science, Electrical and Space Engineering Investigating UV nightglow within the framework of the JEM-EUSO Experiments

Master Thesis Space Engineering, Instrumentation and Spacecraft

Author: Frej-Eric Salomon Emmoth

Supervisors: Dr. Toshikazu Ebisuzaki Chief Scientist at Computational Astrophysics Laboratory, RIKEN & Dr. Marco Casolino Team Leader, EUSO Team, Research Scientist at RIKEN

Examiner: Dr. Johnny Ejemalm Senior Lecturer, Lule˚aUniversity of Technology Acknowledgements I would like to ﬁrst extend my most sincere gratitude to Dr. Ebisuzaki and Dr. Casolino for giving me the opportunity to do my thesis at RIKEN within the JEM-EUSO Collaboration and their invaluable help during my time there. I am also very grateful for the people at the laboratory, many of whom I today consider my friends, who made my stay in Japan so much better. I look forward to once again see the mountains of Nagano.

I next want to give thanks to my family and friends for their constant support and encourage- ment. And especially to my brothers and my girlfriend who were always there to give me a push in the right direction. I am lucky to be surrounded by such great people.

Lastly, I want to express my gratitude to everyone for their patience with me, with special thanks to my supervisors and my examinator in this regard. Thank you. Abstract The main mission of the JEM-EUSO (Extreme Universe Space Observatory) Collaboration is to observe Cosmic Rays. These high energy particles come from a variety of sources and bombard the Earth all the time. However, the higher the energy, the lower the flux, and particles with an energy above 1018 eV (called Ultra High Energy Cosmic Rays or UHECRs) are so sparse that just a few might hit the atmosphere in a year. When CRs, and UHECRs, hit the atmosphere they cause what is called Extensive Air Showers, EAS, a cascade of secondary particles. This limits the effectiveness of ground based observatories, and that is where the JEM-EUSO Collaboration comes in. The goal is to measure UHECRs, by observing the fluorescence of the EAS from space. This way huge areas of the atmosphere can be covered and both galactic hemispheres can be studied.

Since the JEM-EUSO instruments are telescopes measuring in the near UV range, a lot of other phenomena can be observed. One of these applications is UV nightglow. Airglow in general are lights in the sky which are emitted from the atmosphere itself, while nightglow is simply the nighttime airglow. There are many uses of airglow, and one of these is as a medium to observe atmospheric gravity waves.

The aim of this thesis is to investigate how a space-based photon counting telescope, such as those of the JEM-EUSO Collaboration, can be used to measure disturbances in the terrestrial nightglow, to identify atmospheric gravity waves. To accomplish this, a theoretical basis for these interactions was explored and a simple scenario was built to explore the plausibility of measuring UV nightglow modulations. The aim was to see what variables would aﬀect a measurement, and how important they were.

Along side this, a calibration was conducted on one of the JEM-EUSO Collaborations instruments, the EUSO-TA (EUSO-Telescope Array). The goal in the end was to try and measure the night sky, to complement the calculations.

The investigation showed that the conditions during the measurement are very important to the measurement. This includes things like background intensity, nightglow activity, and magnitude/shape of the modulations. Of more importance though are the parameters which can be actively changed to improve the measurement, the most important of which is measurement time. It was concluded that a measurement of the nightglow modulation should be, under the right conditions, possible to do with a currently operating instrument, the Mini-EUSO, or similar instrument.

The calibration of the EUSO-TA involved a series of repairs and tests, which highlighted some strengths and weaknesses of the instrument. However, the calibration itself produced few workable results that in the best case scenario reduced the focal surface to an unevenly biased 2-by-2 Elementary Cell square. Unfortunately this would not be sufficient to do proper measurements with, but the process did point out shortcomings with the then involved sensors, as well as some problematic aspects of the software operating the instrument. Abbreviations List of abbreviations commonly used: AGN : Active Galactic Nuclei AGW : Atmospheric Gravity Wave ASIC : Application Specific Integrated Circuit BW : Bandwidth CPU : Central Processing Unit EAS : Extensive Air Shower EC : Elementary Cell EM : Electromagnetic JEM-EUSO : Joint Exploratory Mission Extreme Universe Space Observatory FD : Fluorescence Detector FS : Focal Surface FoV : Field of View FWHM : Full Width at Half Maximum GBR : Gamma Ray Burst GD : Ground Detector GTU : Gate Time Unit GZK : Greisen-Zatsepin-Ku’min HBI : Herzberg I HVPS : High Voltage Power Supply IR : Infrared ISS : International Space Station JEM : Japanese Experiment Module LED : Light Emitting Diode LEO : Low Earth Orbit LIDAR : Light Detection And Ranging MAPMT : Multi-Anode PMT MLT : Mesosphere-Lower Thermosphere OI5577 : Atomic Oxygen 5577 A˚ PAO : Pierre-Auger Observatory PCE : Photon Collection Efficiency PDM : Photon-Detection Module PMMM : Polymethyl-Methacrylate PMT : Photon Multiplier Tube SNR : Signal-to-Noise Ratio SR : Supernova Remnants SSD : Solid State Drive TA : Telescope Array UHECR : Ultra High Energy Cosmic Rays UV : Ultra Violet Contents

1 Introduction 1 1.1 Thesis Subject ...... 1 1.2 Thesis Scope ...... 1

2 Cosmic Rays 3 2.1 Ultra High Energy Cosmic Rays - UHECR ...... 4 2.1.1 Extensive Air Showers ...... 5 2.1.2 Current Experiments ...... 9 2.2 The JEM-EUSO Project ...... 12 2.2.1 Past and Current Projects ...... 14 2.2.2 Future Projects ...... 17 2.3 The Mini-EUSO ...... 21 2.3.1 Optics, Focal Surface, and Data Acquisition System ...... 22 2.3.2 Acquisition and Usage ...... 25 2.4 The EUSO-TA ...... 27

3 Nightglow 31 3.1 The Herzberg I Bands ...... 32 3.2 The OI5577 Green Line ...... 38 3.2.1 Comparison of OI5577 and HBI emission rate ...... 42 3.3 Atmospheric Gravity Waves ...... 44 3.3.1 Modulation of Airglow ...... 45 3.4 Ultraviolet Background ...... 46 3.5 Measurement Noise ...... 48

4 Measurement of Nightglow 49 4.1 Method and Analysis ...... 49 4.1.1 Expected Modulation of Herzberg I bands ...... 49 4.1.2 Geometry of Measurement ...... 52 4.1.3 Estimating Photon Counts ...... 54 4.1.4 Constructing a Scenario ...... 56 4.1.5 Applying the Scenario ...... 57 4.1.6 Eﬀects of Measurement Noise ...... 57 4.2 Results of the Estimation ...... 58 4.2.1 Singel Pixel ...... 58 4.2.2 The Scenario ...... 59 4.2.3 Full Frame ...... 60

5 Calibration of the EUSO-TA focal surface 64 5.1 Method of Calibration ...... 64 5.1.1 Initial Complications ...... 65 5.1.2 Mapping the Frame ...... 66 5.1.3 Mechanical Problems ...... 67 5.1.4 Orientation of the ECs ...... 70 5.1.5 Diﬀerences between ECs ...... 71 5.1.6 Diﬀerences between Pixels ...... 72 5.1.7 Reordering och ECs ...... 73

4 5.2 Results of the Calibration ...... 74

6 Conclusions 76 6.1 Discussion ...... 77 6.2 Future Work ...... 77 1 Introduction

This Master Thesis deals with using the Mini-EUSO space telescope, and similar instruments, to measure the modulations in ultraviolet nightglow in the lower thermosphere induced by Atmospheric Gravity Waves (AGWs). This work is a part of the larger JEM-EUSO Collaboration, an international cooperation which aims to broaden the understanding of the universe at large, but more speciﬁcally our understanding of Ultra High Energy Cosmic Rays. This thesis is divided in to two parts, one theoretical and one practical. The theoretical part contains the background and framework for the physics describing the phenomena of nightglow and how it can be used, while the practical part contains a description of the calibration process of the EUSO-TA instrument and measurements done at the Wako Campus of RIKEN, Saitama, Japan. This whole project was facilitated by the Computational Astrophysics Laboratory at the RIKEN research institute main campus, who provided the tools, facilities, and assistance necessary to complete these tasks.

1.1 Thesis Subject The main topic of this thesis is an investigation and estimation of measuring nightglow and the intensity modulations that occur due to atmospheric gravity waves breaking in the upper atmosphere. This is done within the framework of a photon counting instrument, the Mini-EUSO space telescope and therefore the focus lies on ultraviolet nightglow from the Herzberg I system. The goal is to investigate the feasibility and extent to which these modulations can be measured. The Mini-EUSO is currently mounted on the ISS (International Space Station) and is pointed in nadir direction, which aﬀects the amount of light measurable from nightglow. This is combined with a process of practical calibration with a similar instrument, the EUSO-TA, and a test measurement.

1.2 Thesis Scope Since the subjects of airglow and atmospheric gravity waves are quite broad on their own, the scope of this thesis has to be limited. The focus lies in investigating how the Mini-EUSO will measure modulations in the UV nightglow and establishing the signal-to-noise ratio to make these measurements possible, and as such these are the goals of this thesis:

• Identify and estimate the relevant parameters of nightglow emissions. • Investigate the eﬀect of AGWs on UV nightglow. • Estimate the UV background for a nadir pointing instrument in orbit.

• Model the measurement of UV nightglow modulations in the framework of the Mini-EUSO. • Calibrate and optimize the EUSO-TA instrument. • Make measurements with the EUSO-TA.

The model described in this thesis is limited, and the following are a few things not taken in to account, or not included:

• An atmospheric model of airglow in the lower thermosphere.

1 • A model of atmospheric gravity waves propagating in the atmosphere. • A developed radiative transfer model from the lower thermosphere to the ISS (or ground). • Measurement data from the actual Mini-EUSO instrument.

• An optimization of the measurement window in regards to spacial maximums in nightglow over the globe compared to the ISS orbital path.

2 2 Cosmic Rays

Cosmic rays are high energy particles, mainly composed of nuclei or protons, originating primarily from outside of our solar system. When they impact the Earth’s atmosphere they can produce ”showers”, cascades of secondary particles which propagate through the atmosphere. The way these showers are detected and how to relate their characteristics to that of of the incident cosmic ray will be discussed later on. Suﬃce to say, when measuring cosmic rays of higher energies it is usually these showers that are measured, since cosmic ray collisions at this level are rare enough to make direct measurements unrealistic. The energy range of these particles lies around a few GeV to 1018 eV, above which they start being called ultra high energy cosmic rays. It is theorized that UHECRs originate from outside of our galaxy, while the lower energy ones come from within the Milky Way.

The effect of cosmic rays were first measured in the early 20th century when scientists, trying to measure the natural radioactivity from the Earth surface, discovered that ionization-rate in air increased based on altitude. The phenomena at the time was called high altitude radiation but later gained the name cosmic rays. It was concluded that the ionization could not come from the ground itself and throughout the first half of the century the research continued, with many concurring experiments, which culminated in an article called Cosmic Ray Theory, by Rossi and Greisen (1941).[1] Since then, and with the foundation set, the nature of cosmic rays has continued to be investigated. While many of the mechanisms have been explored and described there are still many mysteries left unsolved about the limitations of cosmic rays and their origins.

Figure 1: Cosmic Ray Spectrum of energy above 1011eV multiplied by E2. Image is from Letessier- Selvon 2011 (Fig. 1).[2]

Figure 1 shows the cosmic ray spectrum as it is understood today. Low energy cosmic rays are much more plentiful and possible to measure directly, though notably they are susceptible to the magnetic

3 fields of the the heliosphere and geomagnetic field. In general the spectrum follows a power law function, E−α, marked by three changes above 1015eV (the knee), at 3 ∗ 1018eV (the ankle), and finally at 3 ∗ 1019eV (the cutoff ). Between these three points, the knee, ankle, and cutoff, the value of the spectral index α changes. At energies below the knee α = 2.7 and the flux decreases with a factor of 50 as energy increases by an order of magnitude, while above the knee α = 3.0 and the flux decreases with factor of 100. The index changes again above the ankle, until finally hitting cutoff.[2]

The changes in behaviour of the power-law function indicate a difference in how these cosmic rays are generated, how they reach Earth, and/or how they interact with the atmosphere. The general theory is that the cosmic rays below the knee comes from astrophysical objects such as supernova remnants (SR) or binary systems in our galaxy, where these particles are accelerated. The knee itself then supposedly indicates a limit to how much these mechanisms can accelerate the particles in question, and others suggest the limit might be particles escaping the galaxy due to their Larmor radius within the galactic magnetic field at these energies exceeding the thickness of the galactic disc, at least for protons. Thus, the magnetic field would not be able to contain the protons.[3]

The consensus is less certain about the origins of particles in the energy range above the knee and below the ankle. Part of the spectrum above the knee (until the so called 2nd Knee at 4 ∗ 1017eV) is suggested to be heavier nuclei which have yet to reach the required Larmor radius to escape the galaxy.[4] Above the ankle however, it is believed that cosmic rays are generated wholly outside of this galaxy by powerful events like active galactic nuclei (AGN), radio galaxies, and gamma ray bursts (GRB). It is the cosmic rays with energies above the ankle, the UHECRs, which are of interest to the JEM-EUSO Collaboration.

2.1 Ultra High Energy Cosmic Rays - UHECR Ultra high energy cosmic rays, UHECRs, are the subset of cosmic rays with energies above 1018eV, as shown in the bottom right corner of ﬁgure 1. The ﬁrst cosmic ray shower measured to have an energy above 1019eV was discovered by an instrument that had an area of 8km2 (the Volcano Rach air-shower array in New Mexico, Linsey et al., 1961 [5]). In 1963 this was followed by a shower measured to have 1020eV (Linsey, J. [6]) and during the coming years more such events were reported. The highest reported event so far took place in 1991, the so called ”Oh-My-God” particle with an energy of 3.2∗1020eV and was measured by the Fly’s Eye Detector in Utah, USA.[7]

As mentioned earlier, the flux of these energetic events is very low. The flux at the ankle for example 2 is about Fankle ≈ 3 particles/km /year/sr, which might help put their scarcity in perspective. As the energy of the particle increases the flux decreases. This means that large areas of the night sky must be monitored by instruments hoping to investigate these cosmic rays. It is these high energy particles which are of greatest interest, especially those above 5 ∗ 1019eV, as they are particles exceeding the Greisen-Zatsepin-Ku’min limit.

The Greisen-Zatsepin-Ku’min effect The Greisen-Zatsepin-Ku’min effect, GZK effect for short (which was developed independently by three scientists, Zatsepin and Ku’min, and Greisen), places a theoretical limit to the energy a particle travelling through the universe should be able to maintain, without losing energy to pion production by interacting with the cosmic microwave background. For particles to reach Earth with more

4 energy than the GZK limit permits, the source needs to be close enough (R ≈ 100Mpc) to not have time to interact with the CMB.[8] [9]

Particles with higher energies are still reported however, which suggests the existence of accelerators within the 100 Mpc radius that have yet to be discovered, but there are still no veriﬁed answers. To ensure that these high energy measurements are not a case of measurement errors, consistent and exhaustive measurements must be done on UHECRs above 5 ∗ 1019eV. This would provide a proper statistical basis for these phenomena.

2.1.1 Extensive Air Showers Extensive Air Shower, EAS, were first observed in 1938 independently by two scientists, Werner Kolhörster and Pierre Victor Auger. As previously mentioned, EASs are cascades of ionized particles and electromagnetic radiation following a cosmic ray interaction with an atom in the atmosphere. This first interaction then produces energetic hadrons, usually in the form of pions, 1/3 of which are neutral and the others charged. Neutral pions are generally unstable with a mean life-time of 8.4 ∗ 10−17seconds, decaying into EM radiation almost immediately, π0 → γ + γ. These photons in turn decay into electron-positron pairs through pair-production who then produce more photons through the process of Bremstrahlung, and the EM side of the shower continues to propagate thusly. The charged pions produced in the first interaction are relatively longer lived (mean life-time of 2.6∗10−8 seconds) and interact with the medium after a certain length, creating more pions, π+,−,0. Finally, when an energy threshold is reached where the hadronic interaction cannot be sustained, the charged pions decay into muons and neutrinos by the following process: π+ → µ+ + ν and π− → µ− + ν. In this whole process the primary particle goes on to interact with more nuclei in the atmosphere, each adding to the cascade.

Figure 2: Schematic of Extensive Air Showers, illustrating the evolution of hadronic and electromagnetic cascades. From Letessier-Selvon, 2011 (FIG. 3).

When measuring EASs two characteristics are important: the number of particles observed and the depth of the shower maximum. These two quantities relate to the total energy of the shower

5 and the primary’s mass respectively (though the mass can also be estimated from the electron-to- muon ration). The number of particles can be measured by utilizing arrays of detectors on the ground, using for example scintillation detectors to measure the electromagnetic component of the shower and calorimeters for the hadronic components. By measuring particle density and integrating the lateral density distribution, the total number of particles in the shower can be estimated. The direction of the EAS can be determined using the arrival time of the particles in each detector.[3]

The depth of the shower maximum can be measured by detecting the resulting Cherenkov light produced by the shower using sky-facing PMTs or by measuring the ﬂuorescence light emitted by N2 molecules in the air which are excited by the EAS process.

Electromagnetic Shower The EAS is comprised of two parts, one electromagnetic and one hadronic. The pure electromagnetic cascade was ﬁrst described by Walter Heinrich Heitler in 1954.[10] In this model the cascade consists of a binary tree, where at each step the particles interact, either through Bremsthralung in the case of electrons/positrons or through pair production for the photons, and produces secondary particles of the same energy.

Figure 3: Schematic views of Heitlers model of an electromagnetic cascade. From Matthews, 2004 (Fig. 1).[11]

The particles considered in this model are only photons, electrons, and positrons. Energy of the secondary particles in each step is assumed to be equal to that of the parent particle before interaction. A few simpliﬁcations are made for this model, mainly that the cross sections of the processes are taken to be independent of energy and that the loss of energy due to collisions can be ignored.[2] If the primary particle in the Heitler model is an electron, the interaction comes in the form of Bremsstrahlung and the radiated photon will have half of the energy, and the electron will keep the remaining energy and continue on its way until the next interaction, which will transpire in similar fashion. The photon will, as mentioned, produce an e−, e+ pair with its energy split equally between them.

6 To not go into too much detail about the theory, some main quantities will be explained here. The energy of the instigating article can be expressed as

e e E0 = Nmaxξc (1) e where ξc is in the Heitler model called the critical energy and in air usually occurs at 85 MeV.[11] e Nmax is the total EM shower size. The maximum penetration depth can be expressed as γ e Xmax = ncλrln(2) = λrln(E0/ξc ) (2) where nc is the number of splitting lengths for a shower that has reached critical energy, λr is the the radiation length in the medium. The elongation rate is deﬁned as dX Λ = max (3) dlog10E0

Hadronic Shower The other part of the EAS is the hadronic shower, which can be modelled in a similar way as the EM part. Instead of considering the radiation length in the medium, the interaction length, λI , of particles is used, though in a similar fashion, to split the medium in to layers of λI ln(2) and is assumed to be constant. The interaction produce 2Nπ charged poins and Nπ neutral pions. As explained earlier, the π0 quickly decay in to photons and starting EM cascades, while π± pass through a layer and then interact.

Figure 4: Schematic views of hadronic cascade. From Matthews, 2004 (Fig. 1)[11]

In the primary hadronic interaction charged pions, π±, and neutral pions, π0. As before the cascade ± π continues until π energy reach critical levels ξc , where they decay and produce mouns.

7 If the cosmic ray carries the energy E0, then the total number of charged pions after n interaction is n Nπ = (Nch) . The energy of the primary particle, including the EM cascade from eq.(1), can then be expressed as π e π µ E0 = ξc Nmax + ξc Nmax (4) π e where ξc is the critical energy of the hadronic shower and Nmax is the total hadronic shower size. The depth of the hadronic shower maximum is more complex than in the case of the purely EM shower but an approximation can be made (or rather, has been made by Matthews, 2005 [11]) based on the EM evolution

p p air E0 p pr air E0 Xmax = X0 + λr ln( γ,e ) → Xmax = λI ln(2) + λr ln( γ,e ) (5) 3NchEc 3NchEc The elongation rate of the hadronic shower can be calculated to

p pr air E0 dX d(λ ln(2) + λr ln( γ,e )) Λp = max = I 3NchEc (6) d log10(E0) d log10(E0)

Larger Atoms A ﬁnal detail in the model is to expand it to include nuclear primaries. Assuming there is a nucleus with atomic number A, then an interaction of this nucleus can be viewed as a superposition of E0 interactions by A singular nucleons each of energy A , and be described as oﬀset with respect to proton showers A p Xmax = Xmax − λrlnA (7) With the number of muons being A p 1−β Nµ = NµA (8)

8 2.1.2 Current Experiments The most direct way to measure the properties of an EAS is to detect the resulting secondary particles, using big surface detectors. These try to measure the primary direction by the timing of the incoming signals, the total energy of the cosmic ray particle by estimating the particle densities, and the mass of the primary particle by measuring the depth of the shower. However, since the secondary particles are concentrated in the direction of the shower, the EAS cannot be measured if the detector is too far away from the incident shower. It is therefore necessary to build large detectors sites, using several instruments in tandem such as Scintillators, water Cherenkov tanks, muon detectors, Cherenkov telescopes, etc. Couple this with the low ﬂux of UHECRs and the need for large arrays becomes apparent. Also, while not secondary particles of the EAS, Cherenkov light can be measured as a result of the charged particles of the shower passing through the atmosphere at higher velocities than the phase velocity of light in that medium. This light however is also quite directional and is best measured from the ground.

Another way to measure an EAS is to look for the fluorescent light produced by the shower. Such light comes from air molecules excited by the shower and radiate isotropically. This light can be measured from far away and can give information about the amount of particles in the shower, composition parameters, the evolution of the EM cascade, and shower profile as well as penetration depth. There are however limitations of ground-based detectors measuring the fluorescent light, mainly in regards to duty cycle. They can only operate during night and require good weather, which puts the duty cycle at around 10%.[2]

Today there are two operational observatories, but all are ground-based. These are the Pierre Auger Observatory in Argentina and the Telescope Array in USA. These give a good overview of how UHECRs are investigated from the ground.

The Pierre Auger Observatory The Pierre Auger Observatory, PAO, is named after the French physicist Pierre Victor Auger and is a detector array located in the Mendoza Province, Argentina. At 3000 km2 it is the worlds largest detector array dedicated to UHECR measurements. The project has its beginnings in the late 1990s and construction started in the year 2000. Oﬃcial operations commenced in 2008 but it began making measurements as early as 2006, during the construction.

9 Figure 5: The Pierre-Auger Observatory, where each dot represents a surface detector stations. The ﬂuorescence detectors are shown with their ﬁeld of view, as well as the two laser facilities, CLF and XLF. From The Pierre Auger Collaboration, 2015 (Fig. 1). [12]

The array consists of water Cherenkov particle detector stations and air fluorescence telescope, making it a hybrid detector using both surface detectors (SD) and fluorescence detectors (FD). It has 1660 SDs spaced at 1500m from each other in a triangular grid, and 24 FDs in groups of 6 lining the perimeters of the array, as can be seen in figure 5.

The SDs used at the PAO, water Cherenkov particle detectors, consists of a water tank with a sealed liner with a reﬂective inner surface. The tank, with its 3.6 m diameter, contains 12000 l of puriﬁed water and is surrounded by by photomultiplier tubes, PMTs, looking down though clear polyethylene in to the tank. When particles from the EAS pass through the water at relativistic speeds, Cherenkov light is produced. This light is then measured by the PMTs.

The hybrid design of the array allows not only for cross-checking between the instruments but also better characterizing of the shower.[12]

The Telescope Array The Telescope Array, TA, in Utah covers an area of about 730 km2. The TA project was started by members from the HiRes and AGASA projects, and construction began in 2003 and started it its oﬃcial operation in 2008. Similar to PAO the TA is a hybrid detector. However, instead of using water Cherenkov particle detectors it uses plastic scintillation detectors. The array contains 507 of these SDs and 3 FD clusters of 12-14 FDs, which are complemented with LIDAR systems to monitor atmoshperic conditions. The SDs are laid out in a square grid of 1200 m distances, with the SD sites surrounding them.[13]

10 Figure 6: The Telescope Array layout, where the surface detector units are represented by the dots and the FDs are located on the perimeters of the site. From University of Utah, Telescope Array.[14]

The scintillator detectors work by measuring the secondary particles produced in an EAS event as they pass through the detector. The detector device is made out of a scintillating material that emit UV light when excited by passing secondary particles. This light is then gathered through optical ﬁbres into a PMT. Through GPS timing the results in each SD can be compared, so that the direction of the UHECR can be calculated. By recording the number of SDs hit and the signals they produce, the UHECR’s energy can be found.

One of the JEM-EUSO Collaboration’s instruments, the EUSO-TA is installed at the TA, in front of the Black Rock Mesa Fluorescence Detector (BRM-FD) station, which provides external triggers which are often used to calibrate the EUSO-TA.

11 2.2 The JEM-EUSO Project The JEM-EUSO Collaboration consists of several projects and instruments. EUSO, which stands for Extreme Universe Space Observatory, is a collaboration between 19 countries and 93 research institutes to investigate the upper limits of the cosmic ray spectrum. The many instruments that fall under the umbrella of the JEM-EUSO Collaboration are applications of the main design philosophy, each with their own scientific goals while also acting as testing ground for future project. These instruments work as fluorescence detectors, with the ultimate goal of measuring UHECRs from space. By making measurements from space, the instrument can cover areas which are unattainable with ground-based arrays. Due to the orbit of the ISS the instrument is also able to increase the duty cycle of the detector, from the 10% of the ground-ground based FD, up to 20% since the JEM-EUSO instrument does not suffer under the constraints of the day/night cycle. In their function, they can also serve many other purposes, which will be discussed later on.[15]

The roots of the JEM-EUSO project was a proposed satellite experiment called Satellite Observa- tory of Cosmic Ray Showers, or SOCRAS, which never got oﬀ the ground due to the technological demands being out of reach at the time.[16] The idea for the project was revived in 1995 under the name Maximum-energy Air-Shower Satellite, MASS, by Yoshiyuki Takahashi. Over the years this turned into the JEM-EUSO project, and feasibility studies and prototypes started in 2001. The project was delayed until 2006, and was scaled down from an entire satellite to an instrument which could be mounted on the ISS. But from 2006 onward the work has been steadily going, with many instruments being produced with their own, but related, scientiﬁc objectives while building on the common design of the JEM-EUSO Collaboration.

Some of the past and current experiments will be described here, and that includes the EUSO- Balloon, EUSO-SPB (section 2.2.1), Mini-EUSO (section 2.3), and EUSO-TA (section 2.4). A few of the future experiments are also discussed, including the EUSO-SPB2, K-EUSO, and JEM-EUSO (section 2.2.2).

Common Design of the JEM-EUSO Instruments Most JEM-EUSO instruments follow a common design philosophy, relying on it to form the foundation of the instruments. The main one of these is the focal surface of the detector, but many share the optics, electronics and software.

The focal surface consists of Photo-Detection Modules, which are made up of many PMTs. The PMTs are arranged as sets of 64 pixels in a Multi-Anode Photo-Multiplier Tube, MAPMT, who in turn are arranged in a 2x2 pattern to create a Elementary Cell, EC. Each PDM contains 3x3 of these ECs, which puts the pixel count of a PDM at 2304. The number of PDMs per instrument diﬀers, but the main instrument, JEM-EUSO, is planned to have 137 PDMs.[15]

12 Figure 7: Focal surface of the JEM-EUSO instruments and breakdown of the PDM. From ”JEM- EUSO: Extreme Universe Space Observatory onboard Japanese Experiment Module” (Fig. 4.3.1- 1).[15]

The PMTs used, or more accurately, the MAPMTs used are from Hamamatsu Photonics and are named Hamamatsu-R11265U. These MAPMTs have ﬁlter attached to them, to limit the light en- tering the PMTs to a certain band in the UV range which matches the light from air ﬂuorescence produced in an UHECR event, which comes from excited nitrogen molecules who produce light in the range of 300 nm to 420 nm.[2]

Figure 8: Hamamatsu-R11265U MAPMT, without the BG3 ﬁlter glued on. From ”JEM-EUSO: Extreme Universe Space Observatory onboard Japanese Experiment Module”, (Figure 4.3.2.1-1).[15]

Each MAPMT and their channels are connected to a port at one of the three ASICs boards mounted behind the focal surface. These ASICs boards connect to a multiplexer board which interfaces the incoming data to a processing board. The boards and their function will be discussed in more detail in section 2.3.2.

13 Figure 9: Electronics of the EUSO-TA, a) showing the ECs connected to the ASIC boards, and b) showing the ZYNQ and multiplexer.

Another shared aspect of the design is the optics. Many of the instruments use a Fresnel Lens, which is much more compact and lighter than a conventional convex lens of similar diameter. The size of the lens used for each instrument is of course determined by the the size of the instrument itself.

Figure 10: Fresnel lenses of, a) the Mini-EUSO, and b) the JEM-EUSO.

2.2.1 Past and Current Projects EUSO-Balloon One of the earliest instruments of the JEM-EUSO project was the EUSO-Balloon, launched in 2014. The balloon was launched from Canada, the Timmis Stratospheric Balloon Base, and meant to test important systems of the JEM-EUSO design. Aside from the EUSO PDM and the two Fresnel lenses making up a refractive optical system, it was launched with a full compliment of auxiliary

14 instruments, such as an IR camera. At an altitude of around 40 km, with a single PDM coupled with a field of veiw of 12◦ and a flight-time of about 8 hours, it measured an area too small to expect UHECR detections. Instead the EUSO-Balloon was complimented with a helicopter firing lasers in the view-path of the instrument. This was done so that some valid measurements could be done with the instrument and test the capturing algorithm.

Figure 11: EUSO-Balloon Payload.

The flight resulted in a measurement of the Earth’s UV background over the flight path, and thus fulfilled one of its main scientific goals. The full measurements (both the UV and IR) can be seen in figure 12.

15 Figure 12: Results from the EUSO-Balloon ﬂight. Top showing UV results and bottom showing IR. From Miyamoto, 2016 (Figure 3).[17]

EUSO-Balloon Launched 2014-08-24 Mass 250 kg Optics Fresnel lens, 1 m x 1 m Detector Hamamatsu-R11265-103-M64 Num. of MAPMTs 36 Num. of Pixels 2304 Spatial Resolution 1 km Temporal Resolution 2.5 µsec FoV 12◦

Table 1: EUSO-Balloon speciﬁcations.

EUSO-SPB The EUSO-SPB, or EUSO Super Pressure Balloon, was an improvement on the previous EUSO- Balloon. While close to the previous design, the EUSO-SPB tested a few improvements such as a new triggering algorithm. It also came with a smaller silicon PDM which was tested during the ﬂight. As with the previous project, it came with an IR camera. Launched in 2017 from New Zealand, it was expected to have a ﬂight-time of 30-40 days but was cut short due to a helium leakage and only lasted 12 days before it went down in the ocean. Luckily about 30 hours of measurements were recovered, with data from a height of 33 km.

16 Figure 13: EUSO-SPB1, test with a laser pointed upward across the EUSO-SPB1 detector ﬁeld of view. Each of the 8 frames shows a 2.5µs exposure. From Wiencke, 2017 (Figure 9).[18]

EUSO-SPB Launched 2017-04-24 Mass 1230 kg Optics 2 Fresnel lenses, 1 m * 1 m Detector 1 Hamamatsu-R11265-113-M64-MOD2 D1,Num. of MAPMTs 36 D1,Num. of Pixels 2304 Spatial Resolution 1 km Temporal Resolution 2.5 µsec FoV 11.1◦ Detector 2 Hamamatsu-S13361-3050AS-08 D2,Num. of MAPMTs 4 D2,Num. of Pixels 256

Table 2: EUSO-SPB1 speciﬁcations

2.2.2 Future Projects EUSO-SPB2 The EUSO-SPB2 will be a follow-up on the two previous balloon experiments, but with an expanded arsenal of instruments. Planned for 2022, it will have three telescope with diﬀerent functions. One will measure Cherenkov radiation in the horizontal plane, another upwards facing, and lastly a standard ﬂuorescence detector in the nadir direction.

17 Another big diﬀerence is the lack of lenses for the optical system, instead using mirrors. The focal surfaces are also planned to be shaped diﬀerently, the ECs aligned in a more linear fashion.

EUSO-SPB2 Planned Launch 2022 Telescope 1 Cherenkov, hor Optics Fresnel lens, 1 m x 1 m Detector Hamamatsu-R11265-64 Num. of MAPMTs 52 Num. of Pixels 3328 Spatial Resolution 2◦ Temporal Resolution 2.5 µsec FoV 3.5◦ x 45◦ Telescope 2 Cherenkov, up Optics Fresnel lens, 1 m * 1 m Detector Hamamatsu-R11265-64 Num. of MAPMTs 52 Num. of Pixels 3328 Spatial Resolution 2◦ Temporal Resolution 2.5 µsec FoV 3.5◦ x 45◦ Telescope 3 Fluorescent Detector Hamamatsu-R11265-113-M64 Num. of MAPMTs 36 Num. of Pixels 2304 Spatial Resolution 0.2◦ Temporal Resolution 2.5 µsec FoV 3.2◦ * 28.8◦

Table 3: EUSO-SPB2 speciﬁcations.

K-EUSO K-EUSO, KLYPVE-EUSO (itself a acronym for Kosmicheskie Lichi Predl’no Vysokikh Energii) is a future project which is planned to be launched in 2022 to the ISS. Started by the SINP MSU (skolbeltsyb Institute of Nuclear Physics Lomonosov Mooscow State University), it had a preliminary design done by 2012. However, it was not able to reach the specifications needed for its scientific goals. Therefore, in 2013 the KLYPVE project started a collaboration with the JEM-EUSO project, and the development continued from there. The current design uses a Schmidt optic and MAPMT focal surface. It is planned to cover a FoV of 40◦, with an aperture of 2.5 m and focal length of 1.7 m. An overview of the current design can be seen in figure 14.

18 Figure 14: Concept of the K-EUSO design, ”FS” being the focal surface.

K-EUSO Planned Launch 2022 Mirror Spherical, 4 m radius Detector Hamamatsu-R11265-M64 Num. of MAPMTs 1872 Num. of Pixels 119808 Spatial Resolution 0.11◦ Temporal Resolution 2.5 µsec FoV 40◦

Table 4: K-EUSO speciﬁcations.

19 JEM-EUSO The JEM-EUSO is the long-term main project of the collaboration. The goal is to mount this instrument on the ISS, though at the moment there is no set date when this project will commence. It will be using the common JEM-EUSO design but in a much more ambitious scale. With 137 PDMs it will not only be the largest of the JEM-EUSO telescopes, but the most ambitious attempt at observing UHECRs with energies above 1020eV to date. It will have spatial resolution of 560 m and cover an area of 1.4 ∗ 10 5km2 when facing nadir. It will also be able to tilt 30◦, which will allow for an even greater area of observation.

Figure 15: Footprint of the ﬁeld of view of the JEM-EUSO. The blue proﬁle is nadir mode, the white and yellow is when JEM-EUSO is tilted by an angle of 20 and 30 degrees, respectively. From Adams, 2015 (Fig. 5). [16]

The focal surface on this instrument will be slightly curved and not set into a square, as can be seen in ﬁgure 16. As with the other JEM-EUSO instruments, the JEM-EUSO focal surface is planned to be made up PDMs, each with 3 x 3 ECs, each EC containing 2 x 2 MAPMTs, and lastly each MAPMT holding 64 channels (pixels). At 137 PDMs, it would put the pixel count to 315648 individual pixels.

The optics module will consist of three Fresnel lenses and an iris. The first lens, the one facing space, will be a curved doublet Fresnel lens, the second will be a diffractive lens and serve to reduce vignetting factor and as a chromatic corrector, and the third will be another curved doublet lens to focus incoming light to the focal surface, as seen in figure 16.

20 Figure 16: Conceptual design of the JEM-EUSO. From ”JEM-EUSO: Extreme Universe Space Observatory onboard Japanese Experiment Module”, (Fig. 4.2.1-1).[16]

Aside from the optics and the focal surface, there will also be IR cameras and LIDAR to assess atmospheric conditions.

JEM-EUSO Planned Launch Not Determined Mass 1153 kg Dimensions 2.97 m x 3.35 m x 3.63 m Optics Circular, 2.5 m Detector Hamamatsu-R11265-M64 Num. of MAPMTs 4932 Num. of Pixels 315648 FoV 30◦ Spatial Resolution 0.074◦ Temporal Resolution 2.5 µsec

Table 5: JEM-EUSO speciﬁcations.

However, as of this moment it is uncertain when work will continue on the JEM-EUSO instrument.

2.3 The Mini-EUSO The Mini-EUSO is the latest of the JEM-EUSO projects, and is now mounted on the ISS in the Russian Zvezda Service Module. It is thus the ﬁrst JEM-EUSO instrument to measure cosmic rays from orbit, and it represents a big step forward for the JEM-EUSO Collaboration, both as a tech- nical demonstration and in the broader investigation of cosmic rays. While its main objective is measuring UHECRs, it also pursues several other scientiﬁc objectives such as meteors, space debris, atmospheric phenomena, and the UV background of Earth. Together with the main UV detector,

21 the Mini-EUSO also carries visible light and infrared cameras for monitoring weather conditions for measurements. These cameras operate in 1500-1600 nm and 400-780 nm respectively.

Figure 17: Mini-EUSO, mechanical body with subsystems.

The instruments contain many of the common designs of the JEM-EUSO projects, namely the PDM focal surface, Fresnel lens optics, ASICs and ZYNQ combination electronics, and the accompanying software.

2.3.1 Optics, Focal Surface, and Data Acquisition System Starting with the optics, the Mini-EUSO has two round Fresnel Lenses, both double-sided and weigh- ing 0.8 kg, and can be seen in both figure 10 and figure 17. They are made of a UV-transparent material called polymethyl-methacrylate (PMMM) and have a diameter of 25 cm with a thickness of 11 mm. The effective focal length of the optics is 300 mm and has a FoV of 44◦.[19]

The Photon Collection Efficiency (PCE) of the optics, defined as the number of photons which arrive in one pixel size divided by the number of photons incidentupon the front lens, have been estimated using ray tracing simulations, for singular pixels. Three wavelengths were used for the photon beams, 337 nm, 357 nm, and 391 nm, each with equal intensity. This results in the a behaviour which can be seen in figure 18.

22 Figure 18: The photon collection eﬃciency in 1 pixel as a function of the angle at which photons enter the ﬁrst lens. From Capel, 2017 (Fig. 4, top).[19]

As can be seen in figure 18, the efficiency of the optical system for one pixel is around 0.5 but drops after 20◦. The the nadir facing UV-transparent window in the Zvezda module, with a transparency of around 86%, is not included in the simulated instrument response as it will not affect the results, just slightly increase the thresholds.

The focal surface consists of a single PDM, that is a net of 3x3 ECs, or 6x6 MAPMTs (the Hama- matsu R11265-M64), giving a total of 2304 pixels (each pixel being a single PMT). Each MAPMT is covered by a 2 mm thick BG3 Bandpass UV ﬁlter (allowing only UV light trough), with a center wavelength of 365 nm and FWHM of 146 nm. The MAPMTs are powered by a high-voltage power supply (a Cockroft-Walton HVPS), similar to what can be seen in Figure 19.

Figure 19: Caption

The Data Acquisition System of the instrument is made up of the electronics reading the PMTs, the PDM Data Processing (PDM-DP), and a CPU. That is, each MAPMT is connected to an Ap-

23 plication Specific Integrated Circuit (ASIC, a SPACIROC3 ASICs to be specific) board, and each ASIC board is connected to 6 MAPMTs. This gives 6 ASIC boards in total, assembled in a similar fashion to can seen in figure 20. The ASIC boards preamplify and digitize the pulses from the PMTs every 2.5 µs or every 1 gate time unit (GTU). The ASICs are in turn read by a Zynq board (Xilinx Zynq XC7Z030) which is part of the PDM-DP. The Zynq performs triggering and time- stamping on the now digitalized data, sorted in 128 GTU frames, and if it passes the trigger is sent to the CPU for further processing and storage. The CPU is also responsible for the controlling the sub-systems, housekeeping, operational modes, and data reading from the NIR and VIS cameras. The Zynq board, or more specifically the Kintex7 Field Programmable Gate Array (FPGA) on the board deals with alot of the data handling, such as buffering, configuration of the ASICs, triggering, synchronization, and interfacing with the CPU. The Zynq board also controls the HV applied to the PMTs.The triggering and sorting is done to reduce the required data transfer, which without triggering would amount to 0.96 Gbyte/second.

The storage consists of solid state disks (SSDs) that are stored on the ISS. Each supply mission up to the ISS bring new SSDs to the instrument and take the data-ﬁlled SSDs back down. The reason this method is used instead of telemetry is because of the large data amounts the measurements produce. It would not be feasible to transfer the data via telemetry in a realistic time-frame.

Figure 20: ASIC boards of the EUSO-TA. Disassembled from the rest of the instrument.

24 Mini-EUSO Launched 2019-08-22 Mass 30 kg Dimensions 0.37 m x 0.37 m x 0.62 m Optics Circular, 0.25 m Detector Hamamatsu-R11265-M64 Num. of MAPMTs 36 Num. of Pixels 2304 FoV 44◦ Spatial Resolution 6.11 km Temporal resolution 2.5µsec

Table 6: Mini-EUSO speciﬁcations.

2.3.2 Acquisition and Usage When the photon pulse is read from the PMTs by the ASICs, which is then read by the Zynq which puts the GTU in a 128 GTU sized buﬀer. The data is put through three triggers by the Zynq.

During the reading of these 128 GTUs, the data is read as the number of photons detected in each pixel during 1 GTU, and an average is made of 8 GTUs, a so called moving average. This average is compared with a calculated average over an entire buffer of 128 GTUs, and if the threshold is reached the 128 GTUs centered around the trigger point are stored in to a buffer between the Zynq and CPU. This is the first trigger stage.

The next trigger takes an average of 128 GTUs, adds a weight (named P) to the average and then compares it to a calculated limit (generally also weighted, named N). If the weighted average reaches the limit, the 128 GTUs are stored in the buﬀer.

Lastly, the third ”trigger” is not used as a trigger, rather, it is used to store the average of the 128 GTUs. The 128 GTUs are labelled differently depending on which trigger it has gone through, a packet from the first trigger being named L1, from the second trigger L2, and from the third L3. These three different data streams differ in their temporal resolution and are used for different purposes. L1 being mainly for cosmic rays has a resolution of 2.5 µs, L2 for Transient Luminous Events or TLEs (320 µs), and L3 being used for continuous readout (40.96 ms).[19]

25 Figure 21: Schematic of the trigger system used by the Mini-EUSO.

When the data has been stored by the CPU will be in a raw format. When the SSDs are taken down to Earth, the data will then be converted to a more workable format, specifically ROOT TTrees. The ROOT framework is the main software framework which Mini-EUSO data is processed. A ROOT TTree is structured somewhat as a ”tree” with ”branches” and ”leaves”, each being defined by the user. The EUSO TTree architecture can be seen in figure 22, but will not be discussed at depth in the report. Suffice to say, it is within the ROOT framework which post-measurement data processing is handled.

26 Figure 22: Visualisation of the ROOT TTree data structure.

For the EUSO TTree, the branch represents the measured photon counts of a frame (held in a 4D array) and has 3200 entries. The array has the following structure: PhotonValue = array[CCB number][PDM number][x pixel][y pixel] (CCB is the Cluster Control Board).

2.4 The EUSO-TA The EUSO-TA is a ground based telescope situated at the Telescope Array in Utah, USA. There it is acting as one of the TAs fluorescence detectors, and working in coordination with other instruments at the site. With a single PDM and two Fresnel lenses, it was one of the first pathfinding projects and has been operational since 2015.

As with other JEM-EUSO instruments, the focal surface is covered with a band pass ﬁlter which in this case allows UV in the 290nm-430nm range to pass through. The MAPMTs used are the R11265-M64 from Hamamatsu, and the FoV of each MAPMT is 0.2◦ x 0.2◦. The total FoV of the detector is 10.6◦ x 10.6◦. Since the instrument is designed for single photon detection, each pixel has a gain of > 106. The detection eﬃciency is 30%.[20]

27 Figure 23: Optics (left) and focal surface (right) of the EUSO-TA. From Bisconti, 2016 (Fig. 2).[21]

Figure 24: The EUSO-TA detector, inside the small building, in front of the Black Rock Mesa Fluorescence Detector station, Telescope Array. From Bisconti, 2016 (Fig. 1).[21]

Four data acquisition campaigns in the year 2015 and one in 2016 were done with the external trigger provided by BRM-FDs, for a total of about 140 hours.

28 Figure 25: UHECR event detected on May 13th, 2015, by the EUSO-TA. It shows counts per pixel per GTU on the full PDM. From Bisconti, 2016 (Fig. 4).[21]

Figure 26: Same event as in ﬁgure 25 detected by the BRM-FDs, in horizontal coordinates, where each circle represents one PMT of the BRM-FDs. the red rectangle indicates the EUSO-TA ﬁeld of view. From Bisconti, 2016 (Fig. 4).[21]

29 EUSO-TA Installed March 2013 Start of Operation 2015 Mass 30 kg Dimensions 0.37 x 0.37 x 0.62 m Optics Square, 1 x 1 m Detector Hamamatsu-R11265-M64 Num. of MAPMTs 36 Num. of Pixels 2304 FoV 10.6◦ x 10.6◦ Spatial Resolution 0.2◦ x 0.2◦ Temporal resolution 2.5µsec

Table 7: EUSO-TA speciﬁcations.

As it has been in operation for quite a while, the EUSO-TA has produced a number of interesting results, and continues to be a testament to the designs viability.

30 3 Nightglow

As mentioned, one of the many applications of the Mini-EUSO instrument is the possibility to measure ultraviolet airglow in the upper atmosphere. Since the Mini-EUSO is mounted on the ISS, it will not suffer the same amount of attenuation of the already relatively faint UV emissions caused by the atmosphere. Understanding the mechanisms behind airglow is key to understanding what the atmosphere in the lower thermosphere is made of, how it interacts, and what effects it can have. This can help with monitoring things like solar activity and, in this particular case, atmospheric gravity waves. Atmospheric gravity waves, in turn, are vital in their role of transferring energy, momentum, and chemical species between the different atmospheric layers and in the subsequent influence on upper atmosphere winds, turbulence, temperature and chemistry. Understanding the relation between gravity waves and airglow is vital, and it will be discussed at greater length in section 3.3.1.

Figure 27: Taken from the ISS on the 7th October 2018. From NASA.[22]

Airglow is, generally speaking, the many faint emissions of light in the atmosphere caused by radiation from the Sun, cosmic rays, and chemiluminescence. Usually airglow is separated from Aurora since airglow is a broader term, but they are related.

Airglow was most likely ”discovered” before the 1800s, but the first person to measure it was L. Yn- tema in 1909, who called the phenomena Earthlight.[23] Airglow was distinguished from other lights in the sky by two major factors: higher intensity of light at the horizon which could not be explained by scattered star/galactic light, and higher brightness in other directions than the Milky-Way. The first kind of airglow to be identified was the Green Line of oxygen at 5577A.˚ Since then many different airglows have been identified all over the EM-spectrum and with that a greater understanding of their underlying mechanisms has been reached. Today some of the more important subjects of study are: verification of mechanisms, verifying validity of empirical models, seasonal variations of

31 intensity, latitudinal variations in intensity, finding new lines, effects of lunar activity, effect of solar activity, airglow behaviour on other bodies in the solar system, and influence of ozone depletion. [24]

Airglow spectra can be grouped in to three categories: lines, band systems, and continuums. For this thesis two types are investigated, one line and one bands. Airglow is usually measured in terms of intensity and the unit adopted for this purpose is the Rayleigh, R, and is deﬁned as: 1R = 106 quanta cm−2 sec−1 column−1.

As the name suggests, nightglow is simply airglow occuring at night (its counterpart being called dayglow). These emissions cover almost the entire electromagnetic spectrum, from X-rays to ra- diowaves. Some of the common oxygen-species are the OH Meinel bands, the visible O(1S) green line at 557.7 nm, as well as the Atmospheric and Herzberg Bands of O2. Like many processes in the atmosphere, the nightglow intensity is highly variable in both time and space and as such it is hard to predict the emission rate at an particular moment.

While nightglow is a very broad term, the UV nightglow is of particular interest here due to it being within the bandwidth of mini-EUSO’s detector, 330 to 400 nm. There are several kinds of nightglow in this range, such as the Herzberg I, II, and III systems, as well as the Chamberlain system (to mention the most relevant ones).[25] However, only the Herzberg I system will be discussed since it dominates in intensity in the Earths atmosphere. The HBI system emissions are in the spectral range of 240 nm to 520 nm. Luckily the intensity of the HBI system peaks around 300 to 400 nm, which is suitable for the Mini-EUSO bandwidth.[26]

Figure 28: The airglow spectrum showing 300nm to 440nm, showing most of the Herzberg I band. From Broadfoot and Kendall, 1968.[26]

3.1 The Herzberg I Bands Named after the German physicist Gerhard Herzberg, these bands are a product of a forbidden transition of excited oxygen molecules (meaning the transition is not allowed by the system’s selection rule and has a high probability of not occurring), and produce light with wavelengths of about 240 3 + nm to 520 nm.[27] For these wavelengths to be emitted, the molecular state of O2(A Σu ) is necessary, but how this state is attained in the atmosphere was a topic of debate during the middle of the 20th century. Today the generally accepted model for the transition is a three-body recombination of ground state oxygen atoms

1 1 3 + O( P ) + O( P ) + M → O2(A Σu ) + M (9)

32 3 + 3 − O2(A Σu ) → O2(X Σq ) + hv(Herzberg) (10)

Here M refers to a third body, usually N2 or O2. The notation used here is electronic state ordering 3 3 + + in molecules (A of the O2(A Σu ) expression) and molecular term symbol (Σu of the same expression). A proper explanation of these are not necessary at this level, suffice to say they describe different states of the molecule. Anyways, this reaction produces emissions in the 250-400 nm range and parts of the spectrum fits the bandwidth of the Mini-EUSO (330 nm to 400 nm).

Figure 29: Energy level diagram for O2 and O emissions. From Johnston, 1993 (Fig. 1).[28]

The emission rate of the Herzberg I band can be and has been measured, usually with an intensity of around 300 R. However, it is necessary to understand the mechanisms behind the emission rate, so that the eﬀect of AGWs can be understood and characterized.

General Formulation A general formulation of emission rates in airglow was presented in R.A. Young, 1968.[29] The ∗ emission of an excited state of a species, call it Xn, can be expressed as

∗ ∗ [Xn] I(Xn) = ∗ (11) τ(Xn)

3 ∗ ∗ where I is the emission rate in photons per cm per second, [Xn] denotes the density, and τ(Xn) is the radiative lifetime. In the atmosphere the species will experience quenching, which is when ∗ another species, i, deactivates Xn (reacts with, so that the excited state producing ﬂuorescence is lost) and the rate of quenching can be described by using a coeﬃcient, ki. From this, the rate of ∗ production of Xn, in a steady state, can be expressed as

∗ ∗ X ∗ P (Xn) = (1 + τ(Xn) ki[i])I(Xn) (12)

33 The rate of production (in cm−3) is due to chemical reactions and have several dependencies, such as the species involved, temperature, and density. By ﬁnding the production rates dependence on [i], the rate coeﬃcients (k) can be found.

In nightglow, associative reactions are most often three-bodied

∗ X + Xn−1 + i → Xn + i (13) and give a simple production rate

∗ X P (Xn) = ki[X][Xn−1][i] (14)

X ∗ X ∗ → ka[X][Xn−1][i] = (1 + τ(Xn) ka[i])I(Xn) (15)

P ∗ ka[X][Xn−1][i] → I(Xn) = ∗ P (16) 1 + τ(Xn) ka[i] From here, the emission rate is easily found. However, for a more applicable process the reactions would be ∗ X + Xn−1 + i → Xn + i (17)

∗ Xn + i → Xn + i (18)

∗ Xn → Xn + hv (19) From this process, the emission rate would be P ∗ ka[X][Xn−1][i]kc I(Xn) = P P (20) (kc + kb)[i] This can be complicated further if there are more intermediate steps in the process, but this is suﬃcient for now.

Application on Herzberg I The state producing HBI emission is one possible outcome of the three-body recombination. The three-body recombination of oxygen, in general terms is usually expressed as

∗ O + O + M → O2 + M (21)

As mentioned before, this can produce a number of diﬀerent emissions, depending on the state of ∗ O2. The most relevant of these can be seen in ﬁgure 30.

34 Figure 30: Schematic over the three-body recombination of atomic oxygen occurring in the atmosphere. From Swenson, 1989 (Fig. 1).[30]

∗ Quenching of the O2 state is expressed as

∗ O2 + M → O2 + M (22) And in the case of the HBI, HBII, and Chamberlain emissions, the next step is emission of light

∗ O2 → O2 + hv (23) Looking speciﬁcally at the HBI band

1 1 3 + O( P ) + O( P ) + M → O2(A Σu ) + M (24)

3 + O2(A Σu ) + M → O2 + M (25) 3 + 3 − O2(A Σu ) → O2(X Σq ) + hv(HBI) (26) From these reactions, way of expressing the emission rate is (taken from Thomas, 1981 [31])

2 k2[O] [M]1 IHBI = (27) QA

, where k2 is the rate of reaction, 1 is the fraction of O2 molecules produced in the A state and QA is the quenching factor for the A state. This expression is essentially the same as X, just taking in to account 1. The quenching factor is given by X QA = 1 + τA kAi[Xi] (28) i

35 τA is the radiative lifetime, and kAi is the quenching rate by the ith component, Xi. If the expression for quenching is expanded, taking into account the diﬀerent quenching species O2 and N2 (and to a lesser extent [O]) the formula for the emission rate becomes

2 k2[O] ([O2] + [N2] + [O])1 IHBI = (29) 1 + τA(kA,N2 [N2] + kA,O2 [O2] + kA,O[O]) which is an expression that is more familiar, and also recognisable from other authors.[32] [25]

Emissions occur at an altitude of 80-100 km in the lower thermosphere and are mainly dependent on the atomic oxygen density and temperature (which contributes to the rate of reaction).[31]

Figure 31: Oxygen, Herzberg I and OI5577 green line. From the left: 5577 up, Herzberg up, Oxygen density down and up, 5577 down. The dotted line is measurement data and the solid line is calculations. Image taken from Thomas, 1981 (Fig. 1).[31]

It is clear that the species densities aﬀect the emission rate but to see the temperature dependence, the rates of reaction needs to be expanded. The rate of reaction of eq.(24), k2, was measured by Campell and Gray in 1973 [33]

−33 2 6 −2 −1 k2 = 4.7 ∗ 10 (300/T ) cm mol s (30)

The quenching rate, described by QA, which is also aﬀected by densities of the quenching species, has had its rates of reaction estimated by Thomas 1981 [31] to be

−12 QA = 1 + 1.1 ∗ 10 [O2] (31) or −13 QA = 1 + 2.75 ∗ 10 [N2] (32)

36 It should be noted that these estimations are quite old and other estimations have been done later on, using a slightly diﬀerent quenching parameter (from Melo 1997 [32])

−11 −3 −1 kA,O2 + RkA,N2 = 5.3 ± 0.5 ∗ 10 cm s (33)

Where the constant R is the ratio of [N2]/[O2], assuming an = 0.03 and a high brightness ratio between HBI and HBII.

Observing special events in the atmosphere and how they change the surrounding airglow can help in understanding how these events interact and aﬀect the atmosphere. One such special event is AGWs, which are of special interest here. One way of trying to understand how AGWs interact would be to look for already observed UV modulations which match the signs of an AGW and use those studies as a reference. Unfortunately, such studies have not been done in any extensive way, one of the few reported measurements coming from M. N. Ross et. al., in 1992.[34]

The experiment used a wideband UV camera mounted on a spacecraft in LEO, and took place in early 1988. The images showed clear variability in the HBI emissions, the regular variations being caused by AGWs, and shows that HBI is deﬁnitely aﬀected by gravity waves, which was expected. In one such structure, the measured amplitude of the oscillations at 95 km was measured to be around 10%. The horizontal wave structure in question was measured to be close to 15 km in wavelength.

Figure 32: Herzberg I UV nightglow, with a horizontal wavelength which is comparable to gravity wave driven structures. From Ross, 1993 (Fig. 1). [34]

While promising, the ﬁndings of this article is only a single result from an experiment made in 1988. To show that HBI modulations from AGWs can be observed (and give meaningful information) other examples of nightglow modulation can be used as reference, to compliment the ﬁndings of Ross et. al.. One such species are of particular interest, the OI5577 Green Line.

37 3.2 The OI5577 Green Line The green line of 557.7 nm is a nightglow in the visible range occurring at the same altitudes as HBI. Being in the visible spectrum makes it more easily measured from ground as it doesn’t suffer the same amount of absorption in the atmosphere as HBI. It is an easier target for investigating AGWs and extensive observations have been done. The green line was the first airglow emission to be identified and can be found in most parts of the night sky. In fact, it was common enough to called permanent aurora by L. Yntema and nonpolar aurora by Rayleigh.[24]

Figure 33: Visible green airglow over Auvergne, France on 13th August, 2015. From Clame Reporter / CC BY-SA (https://creativecommons.org/licenses/by-sa/4.0).[35]

The cause of OI5577 emissions is a two-body mechanism of atomic oxygen, but as mentioned before the reaction which precedes it is the same three-body recombination which produces the Herzberg emissions 3 3 ∗ O( P ) + O( P ) + M → O2 + M (34) ∗ 3 1 O2 + O( P ) → O( S) + O2 (35) O(1S) → O(1D) + hv(557.7nm) (36) ∗ M represents a third body, either N2 or O2. O2 represents an excited state of O2, suspected to be 1 − O2(c Σu ) which is the species that radiates the Herzberg II band. The quenching reactions have been omitted here, since they all look the same, but there are 5 quenchings in total: [O2], [N2], [O] for the ﬁrst reaction and [O2], [O] for the second. The emission rate can be expressed in a similar way to HBI, adding another step in the reaction. For the sake of comparison, the volume mission rate of OI5577 here is also taken from R.J. Thomas, 1981 [31]

38 3 k2k3[O] [M]2 IOI = (37) Q1Q2/τ2 X Q1 = 1 + τ1 k1i[Xi] (38) i X Q2 = 1 + τ2 k2i[Xi] (39) i

Here k2 and k3 are the rate of reaction for the reactions seen in eq.(34) and eq.(35) respectively, 2 is ∗ 1 the fraction of atoms produced in the precursor state O2,τ2 the radiative lifetime of O( S), τ2 the ra- ∗ 1 ∗ diative lifetime of O2, while Q1 and Q2 describes the quenching rate of the two states O( S) and O2.

Expanding the expression gives

3 k2k3[O] ([O2] + [N2] + [O])2 IOI = (40) (1 + τ1(k1,O2 [O2] + k1,N2 [N2] + kO[O]))(1 + τ2(k2,O2 [O2] + k2,O[O]))/τ2 In the same fashion as with HBI, the emission rate is mainly dependent on the densities and rates of reaction (k2,k3). The rates of reaction is mainly dependent on the temperature. They have been observed to be (from V.Yu Khomich 2008 [25])

−33 3 −1 k2 = 5.5 ∗ 10 (200/T ) ∗ 2(cm s ) (41) −12 3 −1 k3 = 1 ∗ 10 cm s (42) The quenching rates of reaction have been estimated to

−11 3 −1 k2,O = 5.0 ∗ 10 exp(−305/T )(cm s ) (43)

−12 3 −1 k2,O2 = 4.3 ∗ 10 exp(−865/T )(cm s ) (44)

The k1,∗ rates are not included since there is still some uncertainty to the precursor state produced in the three-body recombination. Anyways, the relevant piece of information to take from this is OI5577’s dependence on atomic oxygen density, the factor [O]3 and the coeﬃcients of rate of reactions which are dependent on temperature.

Modulation of OI5577 As one of the most common airglows in the night sky it is natural that it has been thoroughly observed, which has led to the identification of modulations caused by atmospheric gravity waves. One such instance can be seen in figure 31 taken from R.J. Thomas 1981 paper, where the modulations are especially clear on the downleg measurement of the green line. Imaging of the airglow show these wave patterns quite clearly, as in figure 34, where the intensity of the light is being modulated, causing the wave patterns.

39 Figure 34: AGW wave patterns in OI5577 emissions, taken by the Andes Lidar Observatory, 29th October, 2013. From Vargas, 2018 (Fig. 1).[36]

By measuring the amplitude of the modulations and understanding how the wave aﬀects the intensity, important information can be obtained about the energy and propagation of the wave. The imaging of the modulation patterns is also necessary for the characterization of the wave. All this will be discussed in more detail a bit later. For now, only the eﬀects on the intensity will be shown.

The amount of modulation is of course dependent on the gravity wave that causes it, but a couple of examples can be seen in ﬁgure 35 and 36.

40 Figure 35: All-sky image of OH (left) and OI5577 (right) emissions with AGW structures. From Snively, 2009 (Figure 1).[37]

Figure 36: Gravity waves structures observed over Panhala, India, 18th February 2001. Taken with an all-sky camera, OI5577 (top) and OH (bottom). From Mukherjee, 2003 (Fig. 2).[38]

41 Normal ﬂuctuation usually lie around ±5%, like reported in N. Iwagami et. al. 2005, J.B. Snively et. al. 2010, and S. Perwitasari et. al. 2015 to name a few.[39][37][40] There are also a number of simulations conducted such as T. Horinouchi 2004 which shows intensity ﬂuctuations.[41]

Figure 37: Simulation of the eﬀects of AGWs breaking in the lower thermosphere and the results on OI5577 nightglow. From Horinouchi, 2004 (Fig. 5h).[41]

3.2.1 Comparison of OI5577 and HBI emission rate It is difficult to find a clear relation between HBI and OI5577 in regards to how their emission rate would react to disturbances in temperature and density. This is mainly due to the rates of reaction and quenching factors having a complicated dependence on temperature. Looking at them in the simple form 2 k2[O] [M]1 IHBI = (45) QA 3 k2k3[O] [M]2 IOI = (46) Q1Q2/τ2 it is clear that the atomic oxygen density differs by a factor of [O], but there is a more difficult to interpret relation with the temperature. In the article from R. Thomas, 1981, there is a comparison between the deduced quenching factors for the green line and HBI, seen in figure 38.

42 Figure 38: Altitude profile of quenching factors. The deduced data are in dots and a fit as a solid curve. From left to right: the profiles are green line upleg, green line downleg, and Herzberg upleg. From Thomas, 1981 (Fig. 3).[31]

These proﬁles take into account the measured intensities and [O], as well as models of [N2], [O2], and T . The quenching factors of OI5577 and HBI are not entirely dissimilar above 95 km. In the −3 article a temperature dependence of T is assumed for the k2k3 factor, based on another study (Weinstock, 1978 ).[42] If temperature variations are relatively small, k2k3 would be insensitive to it’s temperature dependence.

Since the quenching by O2 and N2 is strong over the observed region, the emission rate of HBI is 2 3 + 1 proportional to [O] . And while the O2(A Σu ) is not the precursor to the O(S ) state, the green line intensity has been observed to be proportional to the HBI intensity times a factor of [O]. This gives a relationship between the two emission intensities, where a is some constant [31][25][43]

3/2 IOI = a(IHBI ) (47)

This is not a universal relationship, but can arise during conditions like those observed by R. Thomas 1979.[31]

Volume Emission Rate to Intensity The total column emission rate, or intensity, is calculated by integrating the volume emission rate over the column (height) Z ∞ Icolumn = Ivolume(Z)dZ (48) 0

Here Icolumn will have the unit Rayleigh, R, and is usually referred to as the intensity.

43 3.3 Atmospheric Gravity Waves While it has been shown that the most variable factors aﬀecting the emission rates of HBI and OI5577 are the densities and temperature, it is important to show how AGWs aﬀect the atmosphere and consequently nightglow emissions.

Atmospheric gravity waves, AGWs, are the result of perturbations of the atmospheres equilibrium. They propagate because the buoyancy of the atmosphere pushes air upwards while gravity pushes it downwards. AGWs are a key mechanism of the Earth’s atmosphere since it is an important way momentum and energy is transferred to the upper atmosphere, and it drives meridional circulation which results in colder summer- and colder winter mesopause.[44] These are a few examples, but their eﬀect is quite widespread.

The sources of AGWs are manifold though most are a result of disturbances in the troposphere. Some examples would be flow over topography, convective systems in the atmosphere, and jets. The amplitude of the waves grow as they move upwards in the atmosphere, due to lower densities at higher altitude. When the waves finally dissipates they dump their momentum and energy into the surrounding atmosphere. This drives a variety of processes, from local generation of turbulence to large scale circulations. These large scale circulations caused by AGW breaking is especially promi- nent in the MLT region at 50-130 km, where the waves reach high amplitudes. This forcing of large circulation can reach a global scale and are a big part of how Earths atmosphere functions. It is also in the lower thermosphere that many interesting airglow occur, and when the AGW breaks at those altitudes they create wave patterns of thermal and density fluctuations.

A proper model will not be presented in this thesis, but a few key characteristics of the AGW will be discussed. Starting in one end, the AGW can be described by the vertical momentum ﬂux, FM , and vertical energy ﬂux, FE. They are related to a few intrinsic parameters of the AGW (taken from G.R. Swenson et. al. 2003 [45])

2 * 0 2+ −ρ0λzg T FE = 2 (49) λhτBV N T

m FM = FE (50) ρ0ω where ρ is the atmospheric density, T the temperature, T 0 the temperature perturbation, g the grav- itational acceleration, N the Brunt-Vaisala frequency, λz the vertical wavelength, λh the horizontal wavelength, m the vertical wave number, τBV the Brunt-Vaisala period, and ω the wave frequency. The wave potential energy is expressed as * + g2 T 0 2 F = (51) P 2N 2 T

The wave will dissipate either through damping or by the formation of convective or dynamical instabilities. The This wave dispersion can be described by the convective stability parameter (N 2, or the Brunt-Vaisala frequency squared)

g dT g N 2 = + (52) T dz Cp

44 here z refers to the altitude and Cp is the adiabatic lapse rate. Using the dispersion relationship, the vertical wave number m, can be expressed

2 2 2 2 (N − ωI ) 2 ωI 1 m = 2 2 k + − 2 (53) ωI − f γgH 4H where γ is the ratio of specific heats, k is the horizontal wave number, f is the inertial frequency, and H is the scale height. The vertical wave number can be used to acquire the vertical wavelength 2π λ = (54) z m And in similar fashion, the horizontal wavelength is given by 2π λ = (55) x k To relate the gravity wave driven fluctuations in intensity-weighted temperature (hT i) to the fluctuations in vertically integrated intensity (hIi), a transfer function called Krassovsky’s ratio is used[46] hI0i / hIi hηi = (56) hT 0i / hT i

3.3.1 Modulation of Airglow Conventionally, an adiabatic and windless atmosphere is assumed, where a monochromatic wave perturbation is added to the background temperature T0 at position (x, z) [47] 2πx 2πz T (x, z, t) = T0(x, z, t) + Acos( + − ωI t) (57) λx λz Here A is the amplitude of the wave. The vertical displacement δz of an air parcel from its equilibrium height z + δz can be expressed as [48]

T (x, z, δz) ≈ T (x, z) + (Γad − Γ)δz (58)

0 where Γad and Γ are the local and adiabatic lapse rates. The perturbed density, ρ , at a height z can be expressed at equilibrium height z + δz as ρ0(x, z) = ρ(x, z, δz) ≈ ρ(x, z)e−(γ−1)δz/H (59) Given a ρ0, the speciﬁc density perturbations of the involved gases can be found through [N ]0 [O ]0 ρ0 2 = 2 = (60) [N2] [O2] ρ

Because N2 and O2 are well mixed below the mesopause, their perturbations can be associated with the atmospheric density perturbation. However, since [O] is dependent on both temperature and density, and because it is not uniformly distributed in the lower thermosphere, its perturbation is given by [49] [O0] ρ0 1 + DH T 0 = −DH + (61) [O] ρ γ − 1 T where D = d(ln[O])/dz which is the inverse of the local scale height of unperturbed [O]. The unperturbed density proﬁles for [O2], [N2], and [O] can be found in ﬁgure 39.

45 Figure 39: Unperturbed temperature and number density proﬁles based on MSIS90. From Liu and Swenson, 2002 (Figure 2).[45]

If the magnitude of the perturbation in temperature and density can be found, they can be applied to the expression for the volume emission rate of the nightglow.

3.4 Ultraviolet Background When the intensity modulations are measured, it will be against the full UV background, and not simply the unperturbed emission rates of the airglow. Aside from nightglow there are many other sources of UV light in the night sky and the major contributor (outside of nightglow) is moonlight. For a nadir facing instrument starlight and such are not a problem.

Luckily there are direct measurements of the moonlight intensity and the near UV background for a Earth facing instrument, and at this point it is sufficient to use this data. The main source that will be used here is G.K. Garipov et. al. 2005 which describes the measurements taken from the MSU ”Tatiana” satellite.[50] The range which these measurements are taken is 300 - 400 nm, which corresponds well with the Mini-EUSO bandwidth. The effect of the Moon on the UV background can be seen in figure 40.

46 Figure 40: Showing the average UV intensity as a function of the moon phase. From Garipov, 2005 (Fig. 6).[50]

Examples of the UV background without moonlight is shown in ﬁgure 41, where the markings α, β, and γ show peaks in intensity resulting from three cities (Mexico City, Houston, and Los Angeles).

Figure 41: Example of UV background intensity on moonless nights, two diﬀerent circulations. The satellite passed a few cities marked by α (Mexico City), β (Houston), and γ (Los Angeles). From Garipov, 2005 (Figure 5).[50]

These measurements show a low level at less than 108 photons cm−2 sr−1 sec−1. To compare this, the UV background during a full Moon can be seen in ﬁgure 42, where the low levels lie around 109 photons cm−2 sr−1 sec−1.

47 Figure 42: Example of UV background intensity during full moon. From Garipov, 2005 (Figure 7).[50]

The minimum UV background measured was at 3 ∗ 107 photons cm−2 sr−1 sec−1, measured above the ocean and Siberia on moonless nights. The highest lies at 3 ∗ 109 photons cm−2 sr−1 sec−1, during the full Moon.

To achieve the best possible results while measuring the Herzberg intensity modulations, they should take place over oceans or empty, not ice-covered, landmasses, and during moonless nights. It is important to remember that a major part of the UV background comes from the nightglow itself.

3.5 Measurement Noise Aside from the UV background, a photon counting telescope using PMTs will experience shot noise (or Poisson noise). This is a quantum noise eﬀect, related to the discreteness of photons and electrons. The Poisson probability mass function is

λne−λ P (K = n) = (62) n! where K is the event of a photon reaching the sensor (n = 0, 1, 2, ...), λ is the expected value of K (and also its variance), and e is Euler’s number. The signal to noise ratio in regards to shot noise is

S¯ λ SNR = = (63) N σ where S¯ is the mean pixel value, N the noise in the signal, and σ the standard deviation of the shot noise which is equal to the square root of the average number of events λ λ √ SNR = √ = λ (64) λ It should be noted that in this case the SNR only concerns the amount of photons reaching the detector, since more noise may arise from the detector itself.

48 4 Measurement of Nightglow

With the basic theory for the intensity modulations of the Herzberg I band and their cause established, it is time to estimate how (and if) the Mini-EUSO and subsequent instruments will measure it. There are several things to consider when calculating the amount of photons from the modulations that will reach the focal surface and be registered, such as the geometry of the measurement, eﬃciency of the instrument, noise levels, crossing times for the frame, the capturing mechanism of the Mini-EUSO, and a few other that will be discussed.

4.1 Method and Analysis The ﬁrst thing that should be established is the size of the signal that is expected to be measured. That means a range for the intensity and the modulations should be determined.

4.1.1 Expected Modulation of Herzberg I bands Intensity Range of Herzberg I The emission rate of the Herzberg I band has been measured by many but, as has been mentioned earlier, the nightglow emissions are produced by atomic and ionic recombination which means it is highly variable both spatially and temporally. Whatever value is used for this estimation will be a token value, representative of what could be measured, and not necessarily what will be measured.

In R.R. Meier’s Ultraviolet Spectroscopy and Remote Sensing of the upper Atmosphere the Herzberg I band is put at a column emission rate of IHBI = 225 R.[27] As is made clear there, the spectrum itself (shown in this thesis in ﬁgure 43) is taken from Hennes 1966 while the absolute values are obtained through normalization of data taken from Huﬀman et. al. 1980, and thus the values presented are representative of moderately active conditions near the equator at midnight in March of 1978 and April 1979.

49 Figure 43: Composite UV nightglow spectrum, with the O2 spectrum taken from the Hennes (1966) Measurement. From Meier, 1991 (Fig. 9).[27]

A higher value was reported by R.J. Thomas 1981, though this was a single rocket measurement. The intensity measured there was IHBI = 500 R, and can be seen in ﬁgure 31 from earlier. This might be considered a high value, somewhat bigger than normal. The measurement was done at White Sands Missile Range (New Mexico, USA) on 11th of July, 1977.

On a similar note, in an article by C.P. Philbrick et. al. 2018 about the LAICE instrument, they comment on the integrated intensity which they believe to lie at IHBI = 600 R, a value taken from Huﬀman and Jursa’s Handbook of Geophysics and Space Environment 1985.[51][52]

A value in between was reported in Johnston and Broadfoot’s 1993 article where the Herzberg I band intensity was measured to be IHBI = 340 ± 15 R.[28] These measurements where made in 1990-1991 in Tucson, Arizona, USA.

Naturally there are other reported values for the band intensity, but for this thesis the following range seems suﬃcient IHBI = 300 − 600R (65)

Modulation of Herzberg I For the modulation it is necessary to take into consideration the kind of AGW that would cause them. AGWs can range from localized waves with relatively small wavelengths and energy to globally propagating AGWs with country-spanning features. Therefore, inspiration is taken from a number of more easily identiﬁed AGWs.

In Perwitasari et. al. 2015, a concentric AGW was measured by observing the modulations in OH airglow and OI5577 airglow. The modulations in the OI5577 airglow was observed by an all-sky camera at Rikebetsu on the 18th of October 2012. The wave parameters were estimated manu-

50 ally from successive measurement, which put the horizontal wavelength of the OI5577 emissions at λh = 51 km and phase velocity of vph,h = 96 m/s. The emissions were assumed to take place at 96 km with an intrinsic wave period of 9.2 minutes. The vertical wavelength was calculated to λz = 42 km and the modulations in the OI5577 appear to be around ±5%.[40]

Figure 44: AGW structure in OI5577 emissions, observed with the all-sky camera at Rikubetsu on 18th October 2012. From Perwitasari,2015 (Figure 3).[40]

Figure 45: Mapped data from the CGW event on the 18th October 2012. Fitted with estimated circle and center position. From Perwitasari,2015 (Figure 4).[40]

Other measurements of OI5577 modulations, described in M.P. Hickey et al, 1997, were taken at Arecibo, Puerto Rico, during January 1993. Two monochromatic gravity waves were observed. The modulation by the two waves was measured to about ±3.9% + 0.7% and 2.9% + 0.5%. The waves themselves had a horizontal wavelengths of 39 and 28 km, with phase velocities of 26 and 37.5 m/s.

51 Figure 46: OI5577. A complex wave pattern in the OI green line resulting from the intersection of two gravity waves, over Arecibo on 21st January, 1993. The image has been ﬂat-ﬁelded to enhance the wave structure. From Hickey, 1997 (Figure 1).[53]

On the larger end, there have been simulations of planetary waves in the MLT region which resulted in up to 34% variation in green line intensity.[54] However, for the purpose of this report a more modest range of modulations would be better

0 I5577 = 0 − 10% (66)

4.1.2 Geometry of Measurement The next thing to consider is the geometry of the measurement, and how that aﬀects the amount of photons reaching the aperture. The relevant parameters are: area of the Mini-EUSO aperture (AEUSO), altitude of the ISS (hEUSO) and emission layer (hHBI ), the spatial and temporal resolution of the Mini-EUSO (lp and τe), area observed by an entire frame (Aframe), loss due to atmospheric attenuation (µatm), crossing time for the measurement (τx), solid angle of the aperture (ΩEUSO), and lastly the eﬃciency of the instrument (η).

The Mini-EUSO is a nadir facing instrument which will be operating from the ISS. Area of the aperture, spatial resolution, and temporal resolution can be found or derived from table 6. From an altitude of about 410 km, it will observe the nightglow, emitting at an altitude of around 95 km, each frame covering an area of about 260 ∗ 260 km2. Each pixel of the focal surface will count the incoming photons from an area of 6.11 ∗ 6.11 km2 every 2.5 µsec, which will be gathered into a complete frame, 1 GTU. The general gist of the geometry can be seen in ﬁgure 47.

52 Figure 47: Rough sketch of the geometry of the measurement.

Since usual AGW horizontal wave velocity is very low compared to τe, the main constraint for the timing of the measurement is the crossing time, i.e. the time it takes for the Mini-EUSO’s area of measurement to move one lp in the direction of the orbit

lp τx = (67) vF rame

The speed of the frame, vF rame, is determined by the altitude of the measurement, altitude of the instrument and the orbital velocity of the instrument

vEUSO ωEUSO = → vF rame = ωEUSO(hHBI + REarth) (68) hEUSO + REarth Setting the instrument on the ISS and assuming an measurement altitude of 95 km results in

τx ≈ 0.84s (69) The effect of atmospheric attenuation is slightly more difficult to estimate, but since the measurement is done from space, it will be significantly smaller compared to ground-based measurements. Because of this, the effect on the measurement will be minimal, and will be neglected. However, in the future a proper analysis of the attenuation should be included.

The solid angle between the radiating source and the aperture is deﬁned as

AEUSO AEUSO −6 ΩEUSO,⊥ = 2 = = 6.4 ∗ 10 sr (70) r hEUSO − hHBI While not tied to the geometry, the quantum eﬃciency of the detector (η) also needs to be discussed. It is deﬁned as the number of photons which arrive in one pixel size divided by the number of photons incident upon the front lens. For the Mini-EUSO, a worst case scenario would be 8%, and a higher estimate would lie at around 35%.

To summarize the parameters of the geometry:

53 Geometric parameters Parameter Value Unit 2 2 AEUSO 12.5 π m hEUSO 410 km hHBI 95 km 2 lp 6.11 km 2 2 Af 260 km τe 2.5 µs τx 0.84 s −6 ΩEUSO,⊥ 6.4 ∗ 10 sr

Table 8: Parameters related to the geometry of the measurement.

As a side note, the unit Rayleigh which is used when referring to the column emission rate is deﬁned as such 1 R = 1010ph ∗ m−2 ∗ (column) ∗ s−1 = 4π ∗ L Where L is the radiance in 1010ph m−2 s−1 sr−1.

4.1.3 Estimating Photon Counts One Pixel

A single pixel will cover an area of Ap and count all incoming photons within the Mini-EUSO bandwidth for each GTU. The ﬁrst step is then to calculate the intensity from the entire area of Ap by integrating the intensity over the area ZZ Ip(t) = I(x, y, t)ΩEUSO(x, y)dxdy (71) xp,yp where x and y are the directions in the plane, t the time in seconds, I(x, y, t) the total column emissions at (x, y), and ΩEUSO(x, y) is the solid angle of the aperture from (x, y).

Since dIp(t)/dt is very small on the timescale of τe or even τx, its time-dependence can be ignored for now, outside of the starting conditions. The solid angle of the measurement is deﬁned as A Ω(x, y) = EUSO (72) r(x, y)2

The distance to the aperture from each point affects the solid angle. For a more accurate estimation the difference of solid angles in different corners of a pixel should be accounted for, but for this calculation the error is negligible. Therefore, the solid angle of each point in the pixel is assumed to be constant.

The area measured by a single pixel sets the limit of how accurately the column emission rate can be measured. In this case, the column emission rate in the area will be the average calculated from the amount of counts registered by the pixel over the time frame. This means that each area of about 32 km2 will be a single data point for the instrument. That way, the individual points of column

54 emission rate I(x, y) is set to equal the average of all the column emission rates in the pixel, Ip,avg. eq.(71) can then be expressed as ZZ Ip = Ωp Ip,avgdxdy = ΩpIp,avgAp (73) xp,yp where Ap is the area covered by the pixel. It should be noted here that Ip is not the measured intensity of the area, but the amount of photons radiated from the area reaching the aperture per unit time. The amount of measured photons in a pixel is then

Np = ηIpτe (74)

, taking the eﬃciency of the instrument in to account.

Entire Frame The data captured by a single pixel does not provide enough information to identify nightglow modulation. The instrument will not distinguish between photons from the Herzberg I band and other sources within the bandwidth. By gathering the counts of each pixel, a frame can be put together for that GTU.

Figure 48: Example of a EUSO-TA frame. A test-measurement taken 2019-01-08.

It is with these frames that patterns and events might be identiﬁed. One frame covers an area of 260 km2 which allows for the detection of a wide range of AGWs.

55 While the JEM-EUSO instruments operate using the ROOT framework and structures its data in a speciﬁc way, the number of photons reaching the detector during a speciﬁc GTU can be expressed in a more generally understood way, as a vector of Np

Nˆ = [N1,N2,N3, ...] (75) When discussing simpliﬁed scenarios, this way of expressing the frame will be easier to use. Later when measurement data will be handled, it will be done within the ROOT framework.

4.1.4 Constructing a Scenario To give a general idea of how modulations might look to the instrument, a realistic example can be constructed. As previously stated, a proper model for AGW effects on the upper atmosphere is beyond the scope of this project and thus a simplified model, or plausible scenario, will suffice. In this scenario, identified AGWs will be used as reference, similar to the ones already presented. At the same time other necessary components will be included, such as the UV background intensity as seen from space. While these parameters are not from the same time and place, a case will be presented as to why they might be used in constructing the scenario.

A continuous function of the total column emission rate is used, I(x, y). The emissions consist of three parts: the Herzberg I emissions, the modulations of these emissions, and the background which is not aﬀected by the AGW (measurement noise will be considered later on)

I(x, y) = Ibg + IˆHBI (x, y) (76) Since the instrument does not differentiate between the sources of the incoming photons, the UV background and the normal Herzberg I emissions are combined in to Ibg, while the modulations are described by IˆHBI (x, y). Here Ibg is assumed to be constant, while IˆHBI (x, y) depends on the AGW propagation and energy. Instead of using the wave equations for the AGW and from there attaining the amplitude of the modulation at a specific point, an approximation of the horizontal structure of a reference AGW is used. In this case, it is the data presented in Perwitasari et. al. 2015, seen in figure 44 and 45. The wave discussed here is a concentric gravity wave observed at Rikubetsu, Hokkaido, Japan 2012-10-18. Observed there was a single concentric wave packet. Taking inspiration from the predictable horizontal wave pattern, the modulation for a scenario such as this could be described by ˆ 0 ˆ IHBI (x, y) = IHBI ΦAGW (x, y) (77) ˆ 0 where ΦAGW (x, y) describes the wave pattern and IHBI is the maximum amplitude of the modulation. The pattern in this case is a concentric wave packet such that

2 − (r−ct) Φˆ AGW (x, y) = Φˆ AGW (r) = 1 − e λx (78) r = x2 + y2 wher c is the wave speed and t is the time. The amount of modulation measured in the OI5577 emissions lie around 5%, which mean that the modulation of Herzberg I emissions should lie at around 3%.

Lastly, the positioning of the frame should be set. Since the frame extends over a 2502 km2 area, only a small slice of the wave packet will be captured. In this scenario, the frame will be centered towards the northeast, similar to the measurements taken in Rikubetsu.

56 This will result in a I(x, y) mapping in the form of

(r−ct)2 0 − λ I(x, y) = Ibg + IHBI (1 − e x ) (79)

4.1.5 Applying the Scenario The scenario presented is continuous, and it has to be applied to how the instrument will measure it. Eq.(73) describes the amount of photons reaching the aperture bound for a pixel and by multiplying with the GTU of the instrument and the quantum eﬃciency, the number of registered photons for each pixel is ZZ 0 ˆ Np = Ωp η τe (Ibg + IHBI ΦAGW (x, y))dxdy = Ωp η τe Ip,avg Ap (80) xp,yp

Nˆ = [N1,N2,N3, ...] Nˆ can then be mapped, and compared with the scenario and the original measurement.

4.1.6 Eﬀects of Measurement Noise Each pixel will experience shot noise and standard deviation of the noise will be the square root of the incoming number of photons p σp = Np (81) It is very important that the signal of the measured modulation is larger than the shot noise, since a small signal to noise ratio will make the modulations indistinguishable from the noise.

57 4.2 Results of the Estimation 4.2.1 Singel Pixel To know if the effects of the modulations would be measurable, the signal (intensity difference due to modulation) is compared to the noise level (standard deviation of the shot noise). Using eq.74 for a single pixel, the amount of photons captured by a pixel in this scenario could, based only on the UV background, be around Np = 625 − 2085 This would be over a period of 512 GTUs and assuming a quantum efficiency of 8%. The range for the number of photons due to modulation, assuming modulations of 3%, over the same period is

Nmod = 5 − 14

These are somewhat meagre results. Here the signal-to-noise ratio would be: SNRmin = 0.18 and SNRmax = 0.3, just using the minimum and maximum of the pixel counts. While the exact numbers are a bit arbitrary, the ratios are important. Since HBI activity is limited at 600 R, a big part of the background intensity will at a certain point be dominated by non HBI light, thus increasing the noise. The best result is reached when the background is equal to the HBI intensity: SNRbest = 0.44.

There are ways of improving the SNR. While a bigger modulation would be good, it is out of the direct control of the instrument. What can be improved is the quantum eﬃciency of the instrument (η), and the amount of time measuring. Increasing the eﬃciency to 35% yields the following results

Np = 2740 − 9120

Nmod = 20 − 61

This results in the following: SNRmin = 0.30, SNRmax = 0.63, SNRbest = 0.92. To further improve the SNR, the number of GTUs measured can be increased. 512 GTUs is the smallest measurement possible with the Mini-EUSO, but the upper limit is set by the crossing time. The number of GTUs which ﬁt in during the crossing time is quite a lot, but the instrument cannot measure for so long continually. But, by simply increasing the measurement time by a factor of 5 (2560 GTUs) the SNR improves drastically

Np = 13683 − 45610

Nmod = 102 − 305

This gives: SNRmin = 0.86, SNRmax = 1.42, SNRbest = 2.05. Even if the quantum eﬃciency was at 8%, it is still possible to have an acceptable SNR during normal conditions by adjusting the measurement time. However, the conditions do set a limit to what kind of AGWs that can be measured. A period of τx would not be enough time to measure modulations during a full moon for example.

As mentioned earlier, the numbers are somewhat arbitrary, but it is possible to get an indication of what kind of waves are actually measurable and under what conditions.

58 4.2.2 The Scenario Using the intensity equation, eq.(79), the scenario was constructed. It is important to remind of the fact that this is not a model of the night sky or of the event, it is a simpliﬁed scenario to showcase how a limited measurement might take place using the Mini-EUSO or similar instruments. Since the night sky in general is much more varied at such scales as in ﬁgure 49, the frame of the instrument is a bit more reliable simply due to its scale. The image is centered at in the north-east, with the epicenter of the AGW as origo.

Figure 49: CGW scenario for good conditions.

The modulation of HBI was set to 3% which results in a deviation from the background level of 8.4 R, which is 1.4%. As a side not, the reason the deviation is not 3% is due to the bandwidth of the instrument being smaller than the spectrum of the HBI band. The background intensity however was measured in the same BW. A comparison with the CGW used as inspiration can be seen in ﬁgure 50.

Figure 50: Comparison of the scenario with the CGW used as reference.

59 4.2.3 Full Frame The total amount of photons counted by each pixel was calculated using eq.(80). Starting with the worst case, assuming low background emissions, low HBI activity, but a modulation of ±3% as in the inspiration for the scenario. This resulted in the frame seen in ﬁgure 51.

Figure 51: Frame with all pixels: low background intensity, low HBI activity, modulation of 3%, over 512 GTUs.

The wave pattern is barely visible, which is expected due to the bad conditions. The pattern deteriorates with a smaller modulation. If the AGW produced a modulation of around ±1% with the same background intensity, the calculation resulted in a dim pattern, seen in ﬁgure 52.

Figure 52: Frame with all pixels: low background intensity, low HBI activity, modulation of 1%, over 512 GTUs.

Here the wave pattern is not distinguishable at all. This means that to overcome the eﬀects of shot noise, whose magnitude depends mainly on the background radiation, there is a limit to when the

60 wave pattern becomes invisible to the instrument. In this scenario under the current conditions, a modulation of ±4% yields an easily identiﬁable pattern, as seen in ﬁgure 53.

Figure 53: Frame with all pixels: low background intensity, low HBI activity, modulation of 4%, over 512 GTUs.

Adjusting the Parameters A few settings were explored, the ﬁrst of which was a higher background radiation environment. As reported in G.K. Garipov et. al. 2005, the minimum background levels measured was at around 3 ∗ 107 photons cm−2s−1sr−1 during moonless nights. The maximum levels measured was 108 photons cm−2s−1sr−1, which would yield the frame seen in ﬁgure 54.

Figure 54: Frame with all pixels: high background intensity, high HBI activity, modulation of 3%, over 512 GTUs.

61 With a high UV background a high HBI activity can be expected and the wave pattern is visible. However, since high background levels outpace the maximum HBI activity, the noise levels increases in relation to the signal, and it becomes worse during moonlit nights, where G.K. Garipov et al. 2005 reports values up to 3 ∗ 109 photons cm−2s−1sr−1, which resulted in a more diﬀuse frame, seen in ﬁgure 55.

Figure 55: Frame with all pixels: during full moon, high HBI activity, modulation of 3%, over 512 GTUs.

The variables that holds the strongest eﬀect on the ability to measure the modulations are the magnitude of the modulation and the level of background radiation. The strength of the nightglow and the background intensity go hand in hand but only so far, since there are other sources adding to the background as well. For the best results, an AGW would have to cause modulations of at least ±2% or more, during moonless nights where HBI emissions compromise most of the UV background, seen in ﬁgure 56.

62 Figure 56: Frame with all pixels, best case: medium background intensity, high HBI activity, modulation of 3%, over 512 GTUs.

At worst the AGW is weak and only cause slight modulations of around ±1% or less, during a full moon, with high background emissions but low HBI activity, seen in ﬁgure 57.

Figure 57: Frame with all pixels, worst case: Full Moon, low HBI activity, modulation of 1%, over 512 GTUs.

The SNR is most easily improved by having larger modulations. Lastly, the efficiency of the instrument itself affects the measurement greatly. The 8% efficiency used here is a worst-case scenario.

63 5 Calibration of the EUSO-TA focal surface

While the measurements of nightglow modulation is supposed to be done by the Mini-EUSO and other space based telescopes, other instruments of the JEM-EUSO collaboration may also be considered, since they function similarly and are based on the same concept. Aside from the main usage and implementation (ground-based vs. space-based) the EUSO-TA is very similar to the Mini-EUSO in design, though the Mini-EUSO is newer and more optimized. Since the EUSO-TA was available to use for testing, it would suﬃce for this project. However, before it could be used to make measurements it had to be assembled and calibrated. Unfortunately its arrival to RIKEN was delayed and no measurements of the night sky was made. It did however give the opportunity to gain a more in-depth understanding of the instrument. The following is a description of the process of assembly and calibration of the EUSO-TA at RIKEN Wakoshi campus.

5.1 Method of Calibration The calibrations were done using a black-box and a Lambertian screen with ﬁlters, as can be seen in ﬁgure 58.

Figure 58: Set-up of the EUSO-TA inside the black-box. When a test was conducted much of the extra material (such as foam, paper, and bubble wrap) was removed, and the lid was sealed.

The lambertian screen was first checked using another UV photon counter, the ”METER LUCE” sensor. This was to ensure that the Lambertian screen was sufficiently isotropically radiating. The setup for this check can be seen in figure 59.

64 Figure 59: Sketch of the set-up for the calibration of the lambertian screen. The instrument METER LUCE was used to measure the UV light. where the strength of the UV light was measured every 10 degrees, starting from the left (facing the screen). The same angles where measured with and without the screen on, so that the background light could be accounted for.

5.1.1 Initial Complications After ﬁrst assembling the focal surface and electronics of the EUSO-TA, the initial test measurement in the black-box gave the frame shown in ﬁgure 60.

Figure 60: Initial test measurement after the ﬁrst assembly, with several ECs not responding.

This showed that several ECs (what later will be named EC3, EC4, EC7, EC8, EC9) were not activated during the tests. Priority was put on repairing the 5 ECs that were not responding.

65 5.1.2 Mapping the Frame To start, each EC was activated in turn, to see if they would work independently of each other. During this, the frame was mapped and all the ECs were named, and each high voltage connector was numbered, seen in ﬁgure 61 and 62.

Figure 61: Mapping of the frame, with each EC getting a speciﬁc number.

Figure 62: The HVPS output numbered to their speciﬁc EC. Used to troubleshoot the HVPS, and controlling the outputs.

As seen in the initial measurement EC3, EC4, EC7, EC8, and EC9 where not working. All ECs were tested while connected to the HV2 contact, which was known to be working. Diﬀerent ASIC connections were also tested, but gave the same result, which suggested that the problem was neither in the software or electronic boards.

66 This prompted a check for mechanical problems, and thus the focal surface was again disassembled, one EC at a time.

5.1.3 Mechanical Problems When disassembling the focal surface a few things were of concern: the connectors to the high voltage output, the cabling to the EC boards, the connectors on the EC boards to the cables, the EC boards themselves, and the MAPMT connections to the EC board. This was done for each EC. The connector layout for the HV can be seen in ﬁgure 63.

Figure 63: Layout of the HVPS connection.

EC1 had no cable or connector problem, but the board had a wire connecting the grounds of each MAPMT, which had come loose due to problems with the soldering. This was later re-soldered, and secured. The EC itself didn’t seem to suﬀer from any other problems.

67 Figure 64: EC1, loose GND wire which was re-soldered to the EC PCB.

EC2 had no outstanding problems, some dirt and dust was removed from between the EC boards.

On EC3 it was discovered that the cables connecting the HV to the EC was oﬀ. The anode of the EC was connected to 40 V and the 40 V dynode was connected to ground, which was the reverse of what can be seen in ﬁgure 63.

Figure 65: Wires going in to the HVPS socket, where the 40 V wire and GND wire had been switched.

The cables were rewired to be according to the schematic. After that some light cleaning was done of the EC. The problem with the cables was also the fault with EC7, EC8, and EC9 and the rewiring of the HV cables was done to all of them.

68 EC4 had a slightly different problem, as the HV connection to the first dynode was broken. The cable was checked, as well as the connector to the HV supply, which left the EC board connection. This connection, the pin soldered to the board, seemed to have broken off. Unfortunately the connecting pins from the HV cable had already been glued (done to provide stability to the structurally weak pins and solder joints, as well as providing EMI protection). Reattaching/resoldering the connection would require the glue to be removed, which would at the time (with the resources at hand) mean a remake of all the connections between the cables and the EC board, and thus EC4 was excluded from further testing.

Figure 66: Caption

EC5 and EC6 was working as intended before and only required a bit of cleaning. EC7, EC8, and EC9 suffered the same cabling problem as EC3 and were fixed in the same way. After this, and testing the ECs individually again, the focal surface was reassembled. With all ECs except EC4 working again, another test measurement in the dark-box was done, see figure 67.

Figure 67: All ECs working, except EC4.

69 5.1.4 Orientation of the ECs As can be seen in figure 67 there are some odd patterns in the frame, and very uneven responses from different areas. One specific set of patterns could be seen at the middle ECs (EC2,5,7), and by covering the right half of the focal surface (upper half in the frame) with a thicker material, the frame seen in figure 68 was captured.

Figure 68: Comparison of the uncovered (left) and half-covered (right) focal surface. The middle ECs are clearly orientated the wrong way.

The cutoff can be seen from the reaction of the covered MAPMT’s, but the middle strip EC’s are clearly oriented the wrong way, turned 90◦ counter clockwise. EC2 and EC3 were switched to see if it had something to do with the ECs themselves, but it made no difference as seen in figure 69.

Figure 69: Normal order (left) and EC2/EC8 switched (right). The orientation of the ECs did not change, and the sensitivity of the MAPMTs stay consistent for the ECs (not dependent on position in the frame.

This is a software problem related to how the MAPMTs are mapped into the resulting frame, and could not be easily ﬁxed at that time. This artefact simply had to be accounted for when continuing.

70 5.1.5 Differences between ECs To see how the focal surface responded to different kinds of illuminations, a more direct light was applied, instead of the Lambertian screen. A UV LED was used in the setup seen in figure 70.

Figure 70: UV LED setup. For each measurement reﬂecting surfaces were removed and the LED was pointed towards the EC which was being checked.

The LED was pointed to diﬀerent parts of the focal surface with varying voltages applied to it. By directing it towards the right side facing the FS (physically up, EC123), a reaction could be seen on that side, with some PMT’s saturating at relatively low eﬀects. When pointing the same LED with at the left side facing the FS(physically down, EC789), that side light up but with some MAPMTs reacting much more than other.

Figure 71: Focusing the UV LED on EC2 (left) and EC8 (right). Each side responds in the right way. Two things to notice: diﬀerence in MAPMT sensitivity within and between ECs, and excessive reaction from the lower left MAPMT in EC3.

While these differences are problematic, the test showed that the EC’s were at least reacting appro- priately, aside from the earlier orientation problem, to light incoming from different directions and lighting up different areas.

71 By swapping the ECs around slightly, to different ASIC connections, it was concluded that the problem was not with the electronic boards, since the moved ECs acted just as before (as seen in figure 69. That meant the differences in sensitivity could stem from two different directions: uneven gain in the pixels/MAPMTs/ECs and/or problems with the 1 photon threshold of each pixel. Due to the then current software that the EUSO-TA was operating with, the gain of each pixel could not be calibrated, and it was not possible (in an easy way that is) to change it for each MAPMT or EC. Instead, the threshold had to be examined for the pixels with low photon counts.

5.1.6 Differences between Pixels While the large differences between ECs and MAPMTs within the ECs were worrying, the larger problem was the differences between the pixels within the same MAPMT. This is because the MAPMT is a discrete instrument which means that the pixels (individual PMTs) of the MAPMT are operating on the same supply. Thus if there is something inherently wrong with a pixel, it cannot be replaced. To investigate this, a few measurements were done in S-curve mode. By looking at the S-curves of singular pixels it was apparent that some were not working properly, see figure 72.

Figure 72: S-curve of a non-responsive pixel.

The S-curve displays the behaviour of the PMT, the registered photon count to the voltage across the dynodes. It is through this that the threshold voltage for a single count, an important property if any meaningful measurement is to be done, can be visualized and analyzed. The main problem shown in figure 72 is the absence of a bridge pattern, right before the peak (marked in red). It is within the span of this bridge that the threshold needs to be. Compare figure 72 with a functioning pixel seen in figure 73.

Figure 73: S-curve of a responsive pixel.

The PMT is working as intended, and a proper bridge is located right before the peak. Without a proper threshold voltage, the pixels will not count incoming photons as regularly, which brings

72 down the sensitivity. This was unfortunate since the main way of calibrating the pixels was through setting of the threshold voltage.

5.1.7 Reordering och ECs While large parts of the focal surface was performing sub-par, other parts were still functioning. To make the most out of these, the focal surface was rearranged so that the sensitive MAPMTs could be consolidated in to functioning area, which could act as a mini focal surface. The resulting frame can be seen in ﬁgure 74.

Figure 74: Remapping of the frame, with each EC marked. Not an actual measurement, this is an illustration of the new EC ordering.

As can be seen, the sensitive part of the focal surface is still quite spread out. This was due to not being able to rotate the ECs (they were mechanically ﬁxed to a certain rotation). However, in case of an actual measurement the more closely and centered sensitive area will have a higher chance of recording and making sense of an event.

73 5.2 Results of the Calibration After the reshuffle of the ECs within the frame, the final test measurement using the EUSO-TA resulted in figure 75.

Figure 75: Caption

Due to the issues discussed earlier, the attempts to improve the reaction of the pixels did not succeed, and instead the properly reacting ECs were clustered so that a smaller, functioning, focal surface could be achieved.

Inspecting the S-curve of the functioning pixels show the threshold voltages, see ﬁgure 76.

Figure 76: S-curves of pixels (29,26) and (4,20), as examples of relatively good pixel behaviours. Both have clear plateaus before the spike.

While the non-responding pixels showed the behaviour seen in ﬁgure 77.

74 Figure 77: S-curves of pixels (44,28) and (42,4), as examples of bad pixel behaviours. Both have clear plateaus before the spike.

Lower light from the lambertian screen produce the results shown in ﬁgure 78.

Figure 78: Caption

No measurement of the night sky was attempted, mainly due to time constraints and the somewhat poor performance of the PDM.

75 6 Conclusions

Measurement of Nightglow The HBI emissions vary spatially and temporally, but an often reported value is around 300 Rayleighs. Not all of the HBI spectrum is taken into account, due to the bandwidth of the instrument. While very few measurements of HBI modulations have been made, many articles report OI modulations due to AGWs at ±2 − 5%, but it can be larger. As the OI emissions are closely related to HBI, 2/3 under normal conditions there is a IOI = a(IHBI ) relationship between OI emission rate and HBI emission rate.

From this, an assumption can be made about the extent of the HBI modulations. The background UV emissions used as reference here were measured, as described in Garipov et. al. 2005 [50], over a year and only represents the background during that year. What it does show however is a possible UV background, the lower end at 107 all the way up to 1010 photons cm−2 sec−1 sr−1. Other measurements have been done, and similar values have been reported (Barbier et. al. 2005 ).[55]

The wave-equation, eq.(78), in this scenario was used to describe a concentric horizontal wave pattern, inspiration taken from the AGW recorded in Perwitasari et. al. 2015. Important to note is that it does not describe an AGW itself.

The conditions during the measurement is obviously important to the measurement, things like background intensity, Herzberg I activity, and magnitude/shape of the modulations. By changing these parameters as discussed in chapter 5, we can see the change in visibility of the pattern for a set scenario. As expected the visibility is low when intensities are low, modulations small, and AGW λh narrow, since shot noise dominate the diﬀerences between pixels. At high intensities it is a question of strength of the modulation against the background intensity. This gives us an idea of what kind of AGWs are identiﬁable, and under what conditions. It is also possible to rule out more unlikely occurrences. Since the HBI intensity and the background intensity are related (HBI emissions being a part of the measured UV background), the chance of the HBI intensity being low while UV background being high is unlikely (aside from moon nights) and vice versa.

However, of more importance are the parameters which can be actively changed to improve the measurement. As was shown in section 4.2.1, the most important is measurement time. By extending it, the eﬀects of shot noise can be minimized. A higher quantum eﬃciency will of course improve the measurement, but it is not quite as easily solved.

With all the limitations in mind, I believe that a measurement of the nightglow modulation should be, under the right conditions, possible to do with the Mini-EUSO or similar instrument. However, since both AGW- and nightglow activity are sporadic both spatially and temporally, a longsighted observation regime needs to be considered to increase the likelihood of capturing these events. Not only does a suﬃciently strong AGW need to be propagating, but there must also be enough NUV nightglow activity in the aﬀected area, during nighttime with the least bright moonphase, to obtain meaningful data.

Calibrations of EUSO-TA The calibration produced few workable results that in the best case scenario reduced the focal surface to an unevenly biased 2-by-2 EC square. Unfortunately this would not be suﬃcient to do proper

76 measurements with, but the process did point out shortcomings with the then involved sensors, as well as some problematic aspects of the software operating the instrument. Many, if not most, PMTs in the focal surface appeared unresponsive with an inability to properly conﬁgure the threshold, as seen by their S-curves. It was also touched upon in chapter 5, the fact that the sensitivity of individual PMTs within the same MAPMT varied widely. This means that each PMT in a MAPMT reacts very diﬀerently under the same voltages.

At the same time, it was not possible to adjust the gain of individual pixels and thus the diﬀerences in sensitivity could not be compensated for. Lastly it should be mentioned that some of the ECs were mapped incorrectly by the software, resulting in a few MAPMTs being ﬂipped in the produced image.

6.1 Discussion The scenarios constructed in this thesis are hypothetical, drawing inspiration from real and recorded events as a basis, and adapting simplified versions to show how a measurement by the Mini-EUSO or other space-based JEM-EUSO instruments could look. As such it is important to reiterate that these scenarios are not simulations derived from atmospheric models, that the assumed level of nightglow intensity is not coupled to the UV background intensity used, which are in turn not the levels of intensity recorded during the AGW events used as base for the scenarios (and in fact, no UV intensities in the relevant bandwidth were measured during these specific AGW events). The aim was to see if modulations could be measured using the JEM-EUSO instrument. And in that sense, the calculations show that it is possible. By varying the parameters, some of the limits of visual identification have been found. It has been shown that such measurements allow for the characterization of AGWs by observing nightglow modulations, under the right conditions.

The calibration of the EUSO-TA focal surface unfortunately did not involve too much actual calibration, rather a series of mechanical repairs and reshuffles. While most of the ECs were functioning (except for EC4) they proved difficult to calibrate to a usable state. In the end the most responsive ECs were sorted into a 2-by-2 square, but even those were quite uneven. Due to this and to time constraints the instrument was never fully assembled at RIKEN and no test measurements of the night sky were made. Instead, the calibration process acted as a troubleshooting and identified problems with the instrument as it was there at RIKEN.

6.2 Future Work There are several things that in the future could be undertaken to further investigate the UV nightglow modulations. First and foremost is actual measurements of the nightsky in the 300-400 nm range, and capturing HBI modulations. This does rely on having a steady enough scientiﬁc basis, which is the second undertaking. A more thorough investigation of the relationship between the emission rates of OI and HBI, as well as other related nightglows, would be necessary to properly estimate the strength of HBI modulations based on the modulations of related nightglows. Together with that, an AGW model should be explored and a simulation of HBI modulations should be done for a varied set of conditions. Through this, a valid estimation of the modulation strength should be possible.

77 A more updated measurement of the UV background emissions should also be taken into consideration, and applied to any future calculations regarding measurements with an JEM-EUSO instrument.

Lastly, eﬀects of spatial and temporal positioning of the measurement needs to be investigated, with regards to the emission strengths, noise/light pollution, and AGW activity. Nighttime measurements are obviously preferable, but taking into account the Moon phase and other factors such as surface albedo (land/ocean/ice), backscattering from clouds, surface emissions like cities, latitudinal/longitudinal variations in nightglow strength, temporal cycles in nightglow strength, and so on can improve the conditions for observations.

When it comes to the EUSO-TA instrument, it would be interesting to see if it can be utilized for measuring HBI modulations. While ground-based and not as suited as the Mini-EUSO for measuring nightglow, it could still be possible to see modulations. For this a function of the radiative transfer is necessary, as well as consideration of the geometric limitations from being ground-based, among other things. If all this could be accomplished, it could be very fruitful to implement a ground-based study of near UV nightglow and use it as a reference, as well as a jump-oﬀ point for further study of nightglow and AGW phenomena.

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