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

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˚ a University 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 par- ticles with an energy above 10

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 fluo- rescence 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 mea- surement, and how important they were.

Along side this, a calibration was conducted on one of the JEM-EUSO Collaborations instru- ments, 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 mag- nitude/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 sim- ilar 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 work-

able 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

5.2 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 measure- ments 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.

• 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 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 10

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 mea- sure 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 experi- ments, 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 10

eV multiplied by E

. 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

fields of the the heliosphere and geomagnetic field. In general the spectrum follows a power law function, E

, marked by three changes above 10

eV (the knee), at 3 ∗ 10

eV (the ankle), and finally at 3 ∗ 10

eV (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 ∗ 10

eV) 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 10

eV, as shown in the bottom right corner of figure 1. The first cosmic ray shower measured to have an energy above 10

eV was discovered by an instrument that had an area of 8km

(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 10

eV (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 ∗ 10

eV 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 is about F

≈ 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 ∗ 10

eV, as they are particles ex- ceeding 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 produc-

tion by interacting with the cosmic microwave background. For particles to reach Earth with more

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 ∗ 10

eV. 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 Kolh¨ orster 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

seconds, decaying into EM radiation almost immediately, π

→ 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

seconds) and interact with the medium after a certain length, creating more pions, π

. 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 electro- magnetic 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

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 integrat- ing 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 N

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 electromag- netic 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 interac- tion. 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.

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

= N

ξ

(1)

where ξ

is in the Heitler model called the critical energy and in air usually occurs at 85 MeV.[11]

N

is the total EM shower size. The maximum penetration depth can be expressed as

X

= n

λ

ln(2) = λ

ln(E

/ξ

) (2) where n

is the number of splitting lengths for a shower that has reached critical energy, λ

is the the radiation length in the medium. The elongation rate is defined as

Λ = dX

dlog

E

(3)

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, λ

, of parti- cles is used, though in a similar fashion, to split the medium in to layers of λ

ln(2) and is assumed to be constant. The interaction produce 2N

charged poins and N

neutral pions. As explained earlier, the π

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, π

. As before the cascade

continues until π

energy reach critical levels ξ

, where they decay and produce mouns.

If the cosmic ray carries the energy E

, then the total number of charged pions after n interaction is N

= (N

)

. The energy of the primary particle, including the EM cascade from eq.(1), can then be expressed as

E

= ξ

N

+ ξ

N

(4) where ξ

is the critical energy of the hadronic shower and N

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

X

= X

+ λ

ln( E

3N

E

) → X

= λ

ln(2) + λ

ln( E

3N

E

) (5)

The elongation rate of the hadronic shower can be calculated to

Λ

= dX

d log

(E

) = d(λ

ln(2) + λ

ln(

chE^γ,e_c

))

d log

(E

) (6)

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 interactions by A singular nucleons each of energy

, and be described as offset with respect to proton showers

X

= X

− λ

lnA (7)

With the number of muons being

N

= N

A

(8)

2.1.2 Current Experiments

The most direct way to measure the properties of an EAS is to detect the resulting secondary par- ticles, 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 km

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.

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, mak- ing 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 km

. 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]

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.

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 in- struments 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 founda- tion 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]

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.

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

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.

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.

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 stan-

dard fluorescence detector in the nadir direction.

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.

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.

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 10

eV to date. It will have spatial resolution of 560 m and cover an area of 1.4 ∗ 10

km

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 indi- vidual 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.

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,

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.

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,

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.

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]

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.

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 > 10

. The detection efficiency is 30%.[20]

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.

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]

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.

3 Nightglow

As mentioned, one of the many applications of the Mini-EUSO instrument is the possibility to mea- sure 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 radi- ation 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 5577˚ A. Since then many differ-

ent 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

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 = 10

quanta cm

sec

column

.

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(

S) green line at 557.7 nm, as well as the Atmospheric and Herzberg Bands of O

. 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 nm to 520 nm.[27] For these wavelengths to be emitted, the molecular state of O

(A

Σ

) 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

O(

P ) + O(

P ) + M → O (A

Σ

) + M (9)

O

(A

Σ

) → O

(X

Σ

) + hv(Herzberg) (10) Here M refers to a third body, usually N

or O

. The notation used here is electronic state ordering in molecules (A

of the O

(A

Σ

) expression) and molecular term symbol (Σ

of the same expres- sion). 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 O

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 X

, can be expressed as

I(X

) = [X

]

τ (X

) (11)

where I is the emission rate in photons per cm

per second, [X

] denotes the density, and τ (X

) is the radiative lifetime. In the atmosphere the species will experience quenching, which is when another species, i, deactivates X

(reacts with, so that the excited state producing fluorescence is lost) and the rate of quenching can be described by using a coefficient, k

. From this, the rate of production of X

, in a steady state, can be expressed as

P (X

) = (1 + τ (X

) X

k

[i])I(X

) (12)

The rate of production (in cm

) 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 + X

+ i → X

+ i (13)

and give a simple production rate

P (X

) = X

k

[X][X

][i] (14)

→ X

k

[X][X

][i] = (1 + τ (X

) X

k

[i])I(X

) (15)

→ I(X

) = P k

[X][X

][i]

1 + τ (X

) P k

[i] (16)

From here, the emission rate is easily found. However, for a more applicable process the reactions would be

X + X

+ i → X

+ i (17)

X

+ i → X

+ i (18)

X

→ X

+ hv (19)

From this process, the emission rate would be

I(X

) = P k

[X][X

][i]k

P(k

+ P k

)[i] (20)

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 → O

+ M (21)

As mentioned before, this can produce a number of different emissions, depending on the state of

O

. The most relevant of these can be seen in figure 30.

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

Quenching of the O

state is expressed as

O

+ M → O

+ M (22)

And in the case of the HBI, HBII, and Chamberlain emissions, the next step is emission of light

O

→ O

+ hv (23)

Looking specifically at the HBI band

O(

P ) + O(

P ) + M → O

(A

Σ

) + M (24)

O

(A

Σ

) + M → O

+ M (25)

O

(A

Σ

) → O

(X

Σ

) + hv(HBI) (26) From these reactions, way of expressing the emission rate is (taken from Thomas, 1981 [31])

I

= k

[O]

[M ]

Q

(27)

, where k

is the rate of reaction, 

is the fraction of O

molecules produced in the A state and Q

is the quenching factor for the A state. This expression is essentially the same as X, just taking in to account 

. The quenching factor is given by

Q = 1 + τ X

k [X ] (28)

τ

is the radiative lifetime, and k

is the quenching rate by the i th component, X

. If the expression for quenching is expanded, taking into account the different quenching species O

and N

(and to a lesser extent [O]) the formula for the emission rate becomes

I

= k

[O]

([O

] + [N

] + [O])

1 + τ

(k

[N

] + k

[O

] + k

[O]) (29) 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), k

, was measured by Campell and Gray in 1973 [33]

k

= 4.7 ∗ 10

(300/T )

cm

mol

s

(30) The quenching rate, described by Q

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

Q

= 1 + 1.1 ∗ 10

[O

] (31)

or

Q = 1 + 2.75 ∗ 10

[N ] (32)

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])

k

+ Rk

= 5.3 ± 0.5 ∗ 10

cm

s

(33) Where the constant R is the ratio of [N

]/[O

], 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.