Abundance and icy grain chemistry of O 2 Abundance and icy grain chemistry of O 2 in comet 67P/Churyumov-Gerasimenko in comet 67P/Churyumov-Gerasimenko

Abundance and icy grain chemistry of O ₂ Abundance and icy grain chemistry of O ₂ in comet 67P/Churyumov-Gerasimenko in comet 67P/Churyumov-Gerasimenko

through ROSINA/DFMS measurements through ROSINA/DFMS measurements and molecular and kinetic modelling and molecular and kinetic modelling

Andrew Gibbons Andrew Gibbons

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Acknowledgements

First of all, I would like to thank my thesis promoters, Nathalie Vaeck and Johan De Keyser without whom this thesis project would not have existed. Thank you both for having taken time from your busy schedules in order to read through this manuscript.

Nathalie, merci tout d’abord de m’avoir accueilli au sein de ton équipe malgré ma demande tardive et le fait que je n’avais pas fait mon mémoire avec toi. Merci de m’avoir confié ce sujet de thèse très varié qui a fait en sorte que je ne m’ennuyais jamais. ^* Merci également pour l’énorme liberté de recherche que tu m’as ac- cordée pendant ces années, et enfin pour ton encadrement dès que j’en avais besoin, de la préparation du FRIA jusqu’à la rédaction.

And Johan, thank you for trusting me and hiring me during the first year of my thesis, and for agreeing to be my co-promoter for my second FRIA attempt. Thank you for sharing your office and for always having taken the time to discuss the many calibration issues we encountered despite your very busy schedule, and for your always insightful suggestions and ideas.

Je tiens à remercier mon laboratoire à l’ULB, le CQP, à commencer par ses directeurs Michel(s) et Piet pour leur bonne humeur et pour avoir entretenu une si bonne ambiance dans le labo durant ces années. Merci à Jérôme pour ton aide avec KROME, l’intégration des constantes cinétiques et fortran, et pour ton avis sur les approximations farfelues que je faisais. Merci à Cédric et à Antoine d’avoir fait le travail de mise en place de KROME et du réseau chimique qui m’a fait gagner beaucoup de temps. Merci aussi à Cédric d’avoir sup- porté mon encadrement pas toujours idéal jusqu’au bout! Merci à Daniel pour les nombreuses fois où ton aide experte m’a sorti d’un pépin avec Linux. Et merci à Julie de toujours répondre à nos innombrables ques- tions administratives. Je tiens particulièrement à remercier Thibaut, Nicolas, Ludovic, Jérôme, Milaim et Cédric pour leur compagnie lors de diverses conférences qui a rendu les soirées moins monotones. Merci enfin à tous les membres du CQP, les anciens comme les nouveaux, pour la bonne humeur, les midis, les parties de ping-pong endiablées et les nombreux afterworks, drinks et tartes d’anniversaire qui ont animé nos semaines.

Un merci tout particulier va à mes voisins de bureau Thibaut et Nicolas pour m’avoir tenu compagnie pen- dant toutes ces années (ou presque), malgré le fait qu’il paraîtrait que je sois bruyant... Merci pour les dis- cussions, l’entraide, les réarrangements de bureau et évidemment tous ces midis aux combos Seven Days + Tavernier. En particulier, merci à Nico pour les discussions geek, de m’avoir donné goût au jeu de société et pour l’organisation d’innombrables après-midis et soirées jeux, et merci à Thibaut de toujours être partant pour être mon compagnon de boisson et de toujours avoir une oreille prête à l’écoute (c’est sans doute parce qu’elles sont bien dégagées!), que ce soit pour des discussions scientifiques ou personnelles. Merci aussi à Adrien qui n’a pas pu rester aussi longtemps que nous dans le bureau, mais qui reste à la fois un membre spirituel du bureau, un collègue spirituel de Nico et un frère spirituel de Thibaut.

I would also like to thank the Space Physics Division of BISA. First of all, thank you to Frederik for introducing me to DFMS and for the many discussions we have shared regarding calibration issues and results, as well as for proof reading part of this manuscript. Merci à Romain et Gaël pour les discussions inspirantes et vos encouragements. Merci à Fabien de toujours s’occuper des tâches que personne d’autre ne veut faire. Finally, thank you to you all for your company during the Co-I meeting in Gruyères.

* La preuve, je suis resté un an de plus!

Thank you to Eva-Maria Krammer and Martine Prévost for teaching me the ropes of classical molecular dy- namics during my Master thesis.

I would next like to thank the members of my thesis accompanying committee, Thierry Visart de Bocarmé, Vinciane Debaille and Romain Maggiolo, for following the advancement of my thesis over the years. Thank you as well to Martin Rubin, Emmanuël Jehin, Martine Prévost and Pierre-François Coheur for agreeing to be part of the thesis jury and reading this work.

My thanks also go to all of the engineers, technicians and scientists involved in the Rosetta mission and the ROSINA instrument consortium in particular for the past 25 years, without whom we would not dispose of such wealth and quality of data. Thank you also to the ROSINA team at Universität Bern for handling the day-to- day supervision of ROSINA during the Rosetta mission and ensuring each instrument functioned properly, as well as for the calibration campaigns carried out with the DFMS instrument copy. Thanks also go to BELSPO for financing the Rosetta project in Belgium. This work would not have been possible without the computing resources provided by the “Consortium des Équipements de Calcul Intensif” funded by the F.R.S.-FNRS and the Walloon Region and in particular the Vega cluster hosted at ULB. I furthermore acknowledge financing pro- vided by a FRIA research grant of the F.R.S.-FNRS, an Additional Researcher’s Grant issued by BELSPO (contract WE/35/007), a Van Buuren-Jaumotte-Demoulin grant and a De Brouckère-Solvay travel grant without which this thesis could not have been undertaken.

Je termine par remercier ma famille et mes amis, que ce soient les chimistes, ceux que je connais depuis le secondaire et avant ou tous ceux que j’ai rencontrés depuis, pour toutes les soirées, weekends, anniversaires et autres activités plus ou moins alcoolisées qu’on a faites ensemble durant ces années qui m’ont permis de décompresser quand j’en avais besoin.

Merci enfin et plus que tout à Aurélie pour tes encouragements, ton affection et le bonheur que tu m’apportes jour après jour.

月が綺麗ですね。

Abundance and icy grain chemistry of O ₂ in comet 67P/C-G

through ROSINA/DFMS measurements and molecular and kinetic modelling Summary of the thesis

p Thesis context and goals y This thesis has been carried out jointly in the Quantum Chemistry and Photo- physics laboratory (CQP) of the Université libre de Bruxelles (ULB) and in the Space Physics division of the Royal Belgian Institute for Space Aeronomy (BISA) under the direction of Profs. Nathalie Vaeck and Johan De Keyser respectively. The global goal of the thesis is to improve our understanding of the abundance and the interaction between icy grains and O 2 in the atmosphere (or coma) of comet 67P/Churyumov-Gerasimenko (67/C-G). O 2 was recently and unexpectedly discovered at 3.8 ± 0.85 % of the abundance of water in the coma of comet 67P/C-G by Bieler et al. (2015, Nature 526, 678). A first pillar of the thesis therefore consists in the study of the variability of the O 2 abundance in the coma of 67P/C-G as well as the possible detection of a distributed source of O 2 (i.e. part of O 2 being released from grains in the coma). In addition, astrochemical models require very specific conditions to be able to reproduce the amounts of O 2 observed in molecular clouds (e.g. Hincelin et al. (2011, A&A 530, A61)). Combined with the fact that O 2 is trapped in water ice in the nucleus of 67P/C-G, this could indicate that the exchanges between the gas phase and the icy mantles of dust grains in molecular clouds are poorly understood, as was recently pointed out by Rawlings et al. (2019, MNRAS 486, 10–20). As a result, a second pillar of the thesis consists in determining the desorption rate of O 2 from the surface of icy grains using computational chemistry methods. This desorption rate can be of use both for estimating the contributions of a distributed source in the coma and in astrochemical models. O 2 has the added benefit of being a small molecule, which makes the initial development of the methodology slightly less complex. The third pillar of this thesis consists in modelling the chemistry of the coma of 67P/C-G using a relatively simple kinetic coma model. The calculated desorption rate of O 2 is included in this model in order to assess the impact of a potential distributed source on the abundance ratio of O 2 . In order to give context to such a model, the results of the ROSINA/DFMS mass spectrometer aboard the ESA Rosetta spacecraft that studied comet 67P/C-G between 2014 and 2016 are used to generate the initial conditions of the coma model as well as to interpret its results.

p DFMS data treatment improvements y The ROSINA/DFMS data calibration process implemented at BISA is

based on the simultaneous calibration of sets of mass spectra measured during the same day. In order to

do so, it relies on the recognition of mass peaks from one spectrum to another. This was a semi-automatic

process since shifting of peak positions across the detector caused misidentifications that needed to be cor-

rected manually. Several effects of instrumental and environmental origin were corrected during the course of

this thesis. In particular, a deconvolution procedure allowed for the reconstruction of doubled peaks on the

detector resulting from an instrumental effect. Varying temperatures in the instrument cause the mass peaks

to arrive at different positions on the detector of DFMS. Correcting for this environmental effect ultimately

rendered a reliable fully automated calibration process possible at BISA. The influence of temperature on the

shape of the peaks was studied in order to improve the accuracy of the results. The absolute sensitivities

of DFMS towards most relevant coma species were re-evaluated using updated experimental and theoreti-

cal values, again yielding more accurate results. Finally, a method was implemented to use total neutral gas

pressure measurements in the coma to normalise the abundance ratios determined using DFMS. This method

was developed at Univerität Bern (UniBE), but had not yet been implemented at BISA. This implementation

allows for a direct comparison between the results obtained independently at BISA and UniBE in order to

cross-validate them.

abundance ratio value of 3.3 ± 1.8 % is found, which is in agreement with the published value by Bieler et al..

However, the average O 2 abundance ratio determined for the same period Bieler et al. studied is 5.0 ± _{2.8 %.}

This value falls outside the uncertainty margins provided in this paper. This difference in large part results from the use of the updated sensitivity of DFMS towards O 2 that differs by about 22 % from the one deter- mined using values given by the National Institute for Standards and Technology (NIST). The variability of the O 2 abundance is studied during the mission, showing that there seem to be 2 regimes corresponding to abun- dance ratios of < 2 % and > 5 % relative to water. No evidence of a potential distributed source of O 2 is found in the data, even though Rosetta described elliptical orbits around comet 67P/C-G towards the end of its mission in order to better highlight any such distributed sources. However, a strong latitudinal dependence of the O 2

abundance ratio was observed during this period, leading to the determination of a region of O 2 -rich ice on the nucleus. This same region does not show higher levels of O 2 at other times during the mission.

p Interaction of O 2 and an amorphous water ice surface y In order to model the desorption of O 2 , a slab of crystalline water ice was built from lattice parameters of the I h water ice structure. While the structure of water ice in the nucleus of comets is still under debate, it has been clear for decades that most of the ice in the interstellar medium is amorphous in structure. For this reason, a slab of amorphous ice was generated using a protocol adapted from the literature by rapidly quenching the melted crystalline structure using the LAMMPS software. Given the amorphous and therefore highly statistical nature of the surface, the adsorption geometry of O 2 is first studied approximately using classical molecular dynamics. In order to better describe its interactions with the amorphous ice, a potential of the O 2 –H 2 O complex was fitted to ab initio calculations of the potential energy surfaces of the system. The CCSD(T) coupled cluster method was used in conjunction with the aug-cc-V5Z Dunning basis set in order to accurately represent the long-range van der Waals interac- tions of the molecules. The dynamics of O 2 on the amorphous ice surface were simulated using the O 2 –H 2 O interaction potential determined on the basis of these computations. Though O 2 can move between more stable adsorption sites on the ice, its position remains globally stable during the simulations. The binding energy of O 2 determined from its interaction potential with water and its adsorption geometry is surprisingly close to recent experimental measurements reported by He et al. (2016, Astrophys. J. 823, 56). A more thor- ough description of the adsorption dynamics of O 2 could not be carried out in the frame of this thesis due to time constraints, but the binding energy of O ₂ can nonetheless be used to study the impact of a theoretical distributed source of O 2 in the coma.

p Chemical modelling of the coma of 67P/C-G y A relatively simple chemical coma model was used in order

to study the chemistry of the coma of 67P/C-G as well as the impact of a theoretical distributed source of O 2 .

This model consists in a Fortran90 embedding script simulating the physical conditions of the coma and that

calls the KROME astrochemical simulation package in order to compute the chemistry of the system for each

cometocentric distance. The KROME package uses the kinetic rate equation approach to solve the chemical

network of the system. The model simulates an expanding cometary atmosphere using radial expansion and

distance-dependent temperature and velocity profiles. A full reaction rate database consisting in over 4000

reactions involving over 300 different ionic and neutral species was generated using the KIDA database. The

initial densities of the coma simulation were determined from the DFMS measurements. Unsurprisingly, the

model shows that little chemistry occurs in the coma of 67P/C-G, given the generally low abundances resulting

from this weakly outgassing comet. The model is used to verify assumptions that were made during the cali-

bration of the detector gain and the fragmentation pattern of the hydrogen halides. The impact of artificially

increasing the densities and including a distributed source of O 2 is also explored.

Contents

I Introduction and objectives 1

1 INTRODUCTION TO COMETS AND MODELLING ³

1.1 A brief introduction to comets . . . . 4

1.2 The Rosetta mission . . . 10

1.3 A brief introduction to coma modelling . . . 12

1.4 O ₂ in astrophysical environments . . . 13

GOALS OF THE THESIS ¹⁶ II Mass spectrometry of comet 67P/C-G 17 2 ROSINA/DFMS INSTRUMENT DESCRIPTION AND CALIBRATION ¹⁹ 2.1 Rosetta/ROSINA . . . 20

2.2 The inner workings of ROSINA/DFMS . . . 22

2.3 Nominal calibration of DFMS spectra . . . 29

2.3.1 Mass scale calibration – L2 → L3 . . . 31

2.3.2 Intensity calibration – L2 → L3 . . . 32

2.3.3 Peak fitting – L3 → L4 . . . 37

2.4 Extracting coma densities from ion fluxes – L4 → _L5 . . . 39

3 IMPROVEMENTS TO THE CALIBRATION OF DFMS SPECTRA ⁴³ 3.1 Instrumental effects . . . 44

3.1.1 Commanded mass 0 . . . 44

3.1.2 Peak doubling . . . 46

3.2 Environmental effects . . . 48

3.2.1 Temperature-dependent ion beam shift . . . 48

3.2.2 Temperature-dependent peak shape . . . 61

3.3 Determining DFMS densities through COPS normalisation . . . 65

3.3.1 Data filtering . . . 65

3.3.2 Determining density ratios from DFMS data . . . 68

3.3.3 COPS data calibration process . . . 74

3.3.4 Normalising DFMS data to the COPS pressure . . . 76

i

4.1 Cross validation of the density determination . . . 82

4.1.1 Absolute densities of the 4 main gas phase species . . . 82

4.1.2 Abundance ratios of O 2 , CO 2 and CO to H 2 O . . . 88

4.2 The O 2 /H 2 O ratio in the coma of 67P/C-G . . . 93

4.2.1 Literature overview of the O 2 /H 2 O ratio in comets . . . 93

4.2.2 Average ratio and long-term variability . . . 93

4.3 Variability of the O ₂ /H ₂ O ratio . . . 97

III Modelling the chemistry of a coma 105 5 O 2 INTERACTION WITH WATER ICE ¹⁰⁷ 5.1 Classical molecular dynamics . . . 108

5.1.1 Classical total force field . . . 109

5.1.2 Integration algorithm . . . 110

5.1.3 Statistical ensembles . . . 111

5.1.4 Limitations of the method . . . 112

5.2 Using CCSD(T) to derive an O ₂ –H ₂ O interaction potential . . . 113

5.2.1 The CCSD(T) method in quantum chemistry . . . 113

5.2.2 Literature data . . . 116

5.2.3 Refining the available data . . . 116

5.2.4 Fitting the PEC to a classical potential . . . 118

5.3 Simulations and results . . . 120

5.3.1 Generating a slab of amorphous water ice . . . 122

5.3.2 Interaction of O 2 with the ASW . . . 124

6 MODELLING THE COMA OF 67P/C-G ¹³¹ 6.1 Simulated astrophysical environment . . . 132

6.2 Reaction rate database . . . 134

6.3 Initial densities measured in the coma . . . 137

6.3.1 Interlude – absolute sensitivities of all species in the coma . . . 138

6.3.2 Absolute densities . . . 143

6.4 Coma simulations . . . 149

6.4.1 Coma simulations from about 10 km . . . 149

6.4.2 Coma simulations closer to the nucleus . . . 153

6.4.3 Coma simulations from perihelion estimates . . . 157

6.4.4 Coma simulations with estimated distributed source . . . 160

IV Concluding remarks and perspectives 167

GLOBAL CONCLUSIONS 169

PERSPECTIVES ¹⁷³

ii

Bibliography 175

Appendix 201

A DFMS CALIBRATION ADDITIONAL INFORMATION A1

A.1 L2 Raw data file . . . A1 A.2 MCP gain and pixel gain correction for channel B . . . A7

B ENVIRONMENTAL EFFECTS ADDITIONAL INFORMATION B1

B.1 Peak shift additional information . . . B1 B.2 Peak shape variation additional information . . . B3 B.3 DFMS density additional information . . . B8

C MASS SPECTROMETRY OF 67P/C-G ADDITIONAL INFORMATION C1 C.1 Absolute densities additional information . . . C1 C.2 O ₂ ratio additional plots . . . C2

D HYPERSPHERICAL EXPANSION OF THE O 2 –H 2 O PES D1

D.1 Rydberg potentials . . . D2 D.2 Expansion moments . . . D6 D.3 Hyperspherical harmonic interaction potential . . . D7 D.4 Mathematical details . . . D9

E POTENTIAL FITTING ADDITIONAL CALCULATIONS E1

E.1 Distance between CM(H 2 O) and O(H 2 O) . . . E1 E.2 Conversion to Cartesian coordinates . . . E2 E.3 Potential fitting routine . . . E7

F MOLECULAR DYNAMICS ADDITIONAL INFORMATION F1

F.1 Center and wrap the system . . . F1 F.2 O 2 –surface interaction energy . . . F2

G CHEMICAL MODELLING G1

G.1 Coma embedding code . . . G1 G.2 Double Gaussian fit parameters for initial densities . . . G3 G.3 Photoreactions of the reaction network . . . G6 G.4 Chemical model additional plots . . . G11

H LIST OF COMMUNICATIONS H1

H.1 Peer-reviewed papers . . . H1 H.2 Oral contributions . . . H2 H.3 Poster contributions . . . H3 H.4 Teaching activities . . . H5

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List of Figures

1.1 Comet C/1995 O1 (Hale-Bopp) in 1997 after its perihelion passage . . . . 3

1.2 Simplified schematic representation of the steps of solar system formation . . . . 4

1.3 Relative production rate ranges measured for volatile molecules in cometary atmospheres through remote sensing . . . . 6

1.4 Solar magnetic field lines and inclination of the magnetic equator . . . . 7

1.5 Parker spiral described by the current sheet of the solar magnetic field . . . . 7

1.6 Draping of the interplanetary magnetic field lines around a comet as a result of the formation of the diamagnetic cavity . . . . 8

1.7 Overview of Rosetta’s journey . . . 10

1.8 Heliocentric and cometocentric distance profiles for the Rosetta mission . . . 11

2.1 High resolution spectrum of mass to charge ratio 32 amu . . . 19

2.2 Photographs of DFMS, RTOF and COPS . . . 20

2.3 Comparison of the characteristics of DFMS and RTOF. . . 21

2.4 Detailed schematic view of ROSINA/DFMS . . . 22

2.5 Ionisation chamber and transfer lens of DFMS . . . 23

2.6 Electrostatic analyser of DFMS . . . 24

2.7 Magnetic analyser of DFMS . . . 25

2.8 Acceleration voltage as a function of the mass to charge ratio and the magnet temperature. . . 26

2.9 Zoom system of DFMS . . . 27

2.10 Ion trajectories arriving on the MCP/LEDA detector with and without zoom . . . 27

2.11 Detector head of DFMS . . . 28

2.12 Schematic representation of the MCP/LEDA detector . . . 29

2.13 Ion trajectories arriving on the MCP/LEDA detector with post acceleration . . . 29

2.14 L2 mass spectrum taken with ROSINA/DFMS and the MCP/LEDA detector . . . 31

2.15 Offset correction for channels A and B . . . 33

2.16 background-corrected L2 mass spectrum taken with ROSINA/DFMS and the MCP/LEDA detector 34 2.17 MCP gain values for channel A . . . 35

2.18 PGC factor deterioration for channel A and GS 16 . . . 36

2.19 PGC factor for channel A and all gain steps for the centre area of the MCP . . . 36

2.20 L3 mass spectrum taken with ROSINA/DFMS and the MCP/LEDA detector . . . 37

2.21 MCP/LEDA schematic peak shape . . . 37

2.22 L4 fitted mass spectrum taken with ROSINA/DFMS and the MCP/LEDA detector . . . 38

2.23 Relative sensitivity fit . . . 40

3.1 Variability of the peak maximum position of CO 2 + as a function of the magnet temperature . . . 43

3.2 L2 commanded mass 0 spectrum . . . 44

3.3 Laboratory reproduction of the charge leakage of the LEDA integrated circuit . . . 44

3.4 L2 spectrum showing the H ₂ O ⁺ peak at a wrong commanded mass . . . 45

3.5 Example spectra showing doubled peaks . . . 46

v

3.8 L2 mass spectra recorded at CM 36 on 31/10/2014 . . . 49

3.9 Temperature profiles of the LEDA, the magnet and the cover motor of DFMS fo the month of October, 2014 . . . 50

3.10 Example mass spectrum measured at commanded mass 44 amu . . . 50

3.11 Peak maximum position of the CO 2 + ion signal on the MCP/LEDA detector for channel A on four separate days . . . 51

3.12 Temporal profiles of the magnet temperature and the peak maximum position for 4 days . . . . 52

3.13 Magnet temperature from 30/08 to 05/09/2014 . . . 52

3.14 Temporal variations and phase space plots of the magnet temperature and peak maximum po- sition for 31/10/2014 . . . 53

3.15 Peak maximum position of the CO 2 + ion signal on the MCP/LEDA detector for channel A as a function of the magnet temperature . . . 53

3.16 Peak maximum position of the CO 2 + ion signal on the MCP/LEDA detector for channel A as a function of the cover motor temperature . . . 54

3.17 Schematic representation of the Rosetta reference frame . . . 54

3.18 Smoothed solar aspect angles of Rosetta from 01/08/2014 to 30/09/2015. . . 55

3.19 Schematic representation of the optimal measuring geometry of DFMS . . . 55

3.20 Modelled illumination function of DFMS . . . 56

3.21 DFMS cover motor temperature and modelled illumination function from 01/08/2014 to 30/09/2015 57 3.22 Orbits described by the Rosetta spacecraft around comet 67P/Churyumov-Gerasimenko (67P/C-G) during the months of September and October 2014 . . . 58

3.23 Peak maximum position variation of the CO 2 + peak in high and low resolution during the first half of August 2014 . . . 58

3.24 Schematic representrations of the possible physical causes for the observed peak position vari- ations . . . 59

3.25 L2 mass spectra recorded at CM 36 for 31/10/2014 corrected for peak shift . . . 60

3.26 Raw and corrected sum spectra of CM 36 on 31/10/2014 . . . 61

3.27 Illustration of the increase in the signal to noise ratio when computing the sum of 24 spectra recorded at commanded mass 36 . . . 61

3.28 Variation of the H 2 16 O ⁺ peak shape, the LEDA temperature and the magnet temperature as a function of time . . . 63

3.29 Variation of the H 2 16 O ⁺ peak shape and the magnet temperature for October 2014 . . . 63

3.30 Variation of the H 2 16 O ⁺ and H 2 18 O ⁺ peak shapes as a function of time . . . 64

3.31 Schematic representation of the DFMS FOV and the off-pointing angle . . . 66

3.32 Dimensions of comet 67P/C-G . . . 67

3.33 Schematic representation of the closest measuring geometry allowing the entire comet to be in view . . . 67

3.34 Schematic representation of the maximum acceptable off-pointing allowing the entire comet to be in view . . . 67

3.35 Schematic representation of the maximum acceptable off-pointing in order for part of the comet to be in view . . . 68

3.36 Maximum acceptable off-pointing in order for the entire or part of the comet to be in the FOV of DFMS as a function of cometocentric distance . . . 69

4.1 O 2 abundance ratio mapped on the nucleus of 67P/C-G close to the end of the mission . . . 81

4.2 Total gas phase density in the coma of 67P/C-G . . . 82

4.3 Total and individual densities of water, CO 2 , CO and O 2 for the duration of the mission . . . 83

vi

4.4 Comparison of total density and the densities of the 4 main cometary obtained at BISA and at

UniBE . . . 88

4.5 Comparison of density ratios obtained at BISA and at UniBE . . . 89

4.6 Electron impact ionisation cross section data provided by the NIST website for O 2 . . . 90

4.7 Comparison between the values obtained for r _{O 2} at BISA and at UniBE for two time periods . . . 92

4.8 O 2 to water ratio from August 12 ^th 2014 to September 5 ^th 2016 . . . 95

4.9 Density of O 2 as a function of the density of H 2 O . . . 96

4.10 Density profiles of O 2 and H 2 O and cometocentric distance of Rosetta . . . 97

4.11 O 2 ratio during the end of mission elliptical orbits . . . 98

4.12 Latitude–longitude plot of the O 2 abundance ratio in the coma in August 2016 . . . 100

4.13 Interpolated O 2 abundance ratio latitude–longitude map . . . 100

4.14 Rosetta spacecraft view of 67P/C-G on 4 arbitrary occasions . . . 101

4.15 O 2 abundance ratio measured in the coma of 67P/C-G mapped to the nucleus of 67P/C-G . . . . 103

5.1 Top view of O 2 on an amorphous water ice surface . . . 107

5.2 Schematic view of the different bonded interaction types and their associated coordinates in a classical force field . . . 109

5.3 Schematic representation of the PES of a protein . . . 112

5.4 Equilibrium structure of the O 2 –H 2 O molecular complex . . . 116

5.5 Coordinate system used to describe the O 2 –H 2 O interaction potential. . . 117

5.6 Potential energy curves of the O 2 –H 2 O molecular complex . . . 118

5.7 Total and individual potential fits to the CCSD(T)/AV5Z data points . . . 121

5.8 Top-view illustration of periodic boundary conditions applied in the x y plane . . . 122

5.9 Slab of crystalline I h water ice obtained by replicating the unit cell . . . 123

5.10 Slab of amorphous water ice obtained through classical molecular dynamics simulations . . . . 124

5.11 Top view of the amorphous water ice surface with an O 2 molecule adsorbed . . . 125

5.12 Example root mean square deviation of the water ice during a production run . . . 126

5.13 Simulation snapshots showing O 2 changing adsorption sites . . . 126

5.14 Extract of the interaction energy of O 2 with the ASW surface during production run 1 . . . 127

5.15 Interaction energy of O 2 with the ASW surface for all runs . . . 127

6.1 Effect of a distributed source on the O 2 relative abundance profile . . . 131

6.2 Temperature and velocity profiles of the coma model . . . 134

6.3 Composite solar flux used to integrate photoreaction cross sections . . . 136

6.4 Comparison between cubic spline and PCHIP interpolation . . . 136

6.5 Total EII cross sections of linear alcohols . . . 140

6.6 Normalised densities of the 4 main coma species for both sets of initial conditions . . . 149

6.7 Density of water and abundance ratios of O ₂ , CO ₂ and CO for both sets of initial conditions . . . 150

6.8 Abundance ratios of 8 secondary coma species for both sets of initial conditions . . . 152

6.9 Abundance ratio profiles of the ionised 4 main coma constituents . . . 152

6.10 Modelled total neutral gas density as a function of cometocentric distance . . . 154

6.11 OH abundance ratio relative to water . . . 155

6.12 Abundance ratio of O 2 relative to water as a function of the cometocentric distance . . . 156

6.13 Abundance ratios of X relative to HX for both sets of initial conditions . . . 156

6.14 Modelled abundance ratio of X and HX relative to water . . . 157

6.15 Density of water and abundance ratios of O 2 , CO 2 and CO for estimated perihelion densities . . 159

6.16 OH abundance ratio relative to water near perihelion . . . 160

6.17 Abundance ratio profiles of the ionised 4 main coma constituents near perihelion . . . 160

6.18 Distributed source of O 2 and its effect on the total O 2 density in the coma model . . . 162

vii

A.2 PGC factor deterioration for channel A and GS 8 . . . A8 A.3 PGC factor deterioration for channel A and GS 9 . . . A8 A.4 PGC factor deterioration for channel A and GS 10 . . . A8 A.5 PGC factor deterioration for channel A and GS 11 . . . A8 A.6 PGC factor deterioration for channel A and GS 12 . . . A8 A.7 PGC factor deterioration for channel A and GS 13 . . . A8 A.8 PGC factor deterioration for channel A and GS 14 . . . A8 A.9 PGC factor deterioration for channel A and GS 15 . . . A8 A.10 PGC factor deterioration for channel A and GS 16 . . . A8 A.11 PGC factor deterioration for channel B and GS 8 . . . A9 A.12 PGC factor deterioration for channel B and GS 9 . . . A9 A.13 PGC factor deterioration for channel B and GS 10 . . . A9 A.14 PGC factor deterioration for channel B and GS 11 . . . A9 A.15 PGC factor deterioration for channel B and GS 12 . . . A9 A.16 PGC factor deterioration for channel B and GS 13 . . . A9 A.17 PGC factor deterioration for channel B and GS 14 . . . A9 A.18 PGC factor deterioration for channel B and GS 15 . . . A9 A.19 PGC factor deterioration for channel B and GS 16 . . . A9 B.1 Unzoomed peak maximum position of the CO 2 + ion signal on the MCP/LEDA detector for channel A B1 B.2 Illustration of the fit of the peak maximum position of the CO 2 + ion signal on the MCP/LEDA

detector for channel A as a function of magnet temperature for four separate days . . . B2 B.3 Raw solar aspect angles of Rosetta from 01/08/2014 to 30/09/2015. . . B2 B.4 Variation of the peak shape of various ions for October 2014 . . . B7 C.1 Comparison between the values obtained for r _{O 2} at BISA and at UniBE for the month of October

2014. . . C2 C.2 O 2 ratio as a function of time and subspacecraft longitude . . . C3 C.3 Latitude-longitude plot of the O 2 abundance ratio in the coma in October 2014 . . . C3 D.1 Coordinate system used to describe the H 2 O–O 2 interaction potential. . . D1 D.2 Representation of the 18 leading configurations used to build the potential energy function . . D2 D.3 Fifth order extended Rydberg potential fits for all 15 non redundant leading configurations of

the H 2 O–O 2 system . . . D3 D.4 Rydberg fitted curves corresponding to the 15 leading configurations . . . D4 D.5 Expansion moments for each of the α configurations and total isotropic term of the interaction

potential . . . D7 E.1 Coordinate system and H 2 O geometry . . . E1 G.1 Density of water and abundance ratios of O 2 , CO 2 and CO for both sets of initial conditions

without photoreactions . . . G11 G.2 X and HX abundance ratios relative to water modelled using previous versions of the coma

model and the reaction network . . . G12

viii

List of Tables

2.1 High resolution zoom factors of DFMS depending on the commanded mass . . . 31

3.1 DFMS measuring sequence on October 23 ^rd , 2014 . . . 45

3.2 Ions detected by DFMS around m/Z = 16, 17 and 18 amu . . . 46

3.3 Fragmentation fractions of CO 2 measured in the DFMS instrument copy . . . 69

3.4 Absolute sensitivities and fragmentation fractions of the main 4 gas phase coma species . . . . 70

3.5 Absolute sensitivities of neon for all three emission currents of DFMS . . . 71

3.6 Partial and total EII ionisation cross sections of O 2 and CO for a 45 eV collision energy . . . 71

3.7 EII ionisation cross sections and theoretical and experimental absolute sensitivities for several atomic and molecular species . . . 73

3.8 COPS pressure correction factors . . . 75

4.1 Examples of average coma compositions and relative uncertainties of densities in the coma . . 87

4.2 Total electron impact ionisation cross sections, fragmentation fractions and absolute sensitiv- ities of O 2 . . . 90

4.3 Measurements of the O 2 /H 2 O ratio in comets . . . 94

4.4 Measurement times and O 2 abundance ratios corresponding to the 4 comet views . . . 101

5.1 Intramolecular geometrical parameters of H 2 O and O 2 . . . 117

5.2 Interaction potential parameters of the O 2 –H 2 O molecular complex . . . 118

5.3 Optimised Buckingham fitting parameters . . . 120

5.4 Interaction energy of O 2 with the ASW surface for each production run . . . 128

6.1 Ions of interest that can be produced as a result of the ionisation of several molecules . . . 138

6.2 Absolute sensitivities and fragmentation fractions for most species detected in the coma of 67P/C-G . . . 141

6.3 Fragmentation fractions of different molecules leading to the same ion . . . 142

6.4 ion fluxes, literature comparisons of the abundance ratios and initial densities on October 19 ^th 2014 . . . 144

6.5 ion fluxes, literature comparisons of the abundance ratios and initial densities on October 20 ^th 2014 . . . 145

6.6 Initial densities of the chemical model . . . 148

6.7 Parameters used to derive the abundance of OH in the coma . . . 155 B.1 Partial EII cross sections of n – C 4 H 10 at 45 eV . . . B8 C.1 Relative uncertainties of densities in the coma using a larger error on theoretically determined

sensitivities . . . C1 D.1 Corrected Rydberg fitting parameters of the 15 leading configurations of the O 2 –H 2 O molecular

complex . . . D5 G.1 Peak fit parameters of ions measured on October 19 ^th and 20 ^th 2014 . . . G5 G.2 Photoreactions of the chemical network . . . G10

ix

1

PART I

Introduction and objectives

1 INTRODUCTION TO COMETS AND MODELLING ³

1.1 A brief introduction to comets 4

1.2 The Rosetta mission 10

1.3 A brief introduction to coma modelling 12

1.4 O ₂ in astrophysical environments 13

GOALS OF THE THESIS 16

.

INTRODUCTION TO 1

COMETS AND MODELLING

Figure 1.1: Comet C/1995 O1 (Hale-Bopp) in 1997 after its perihelion passage. Both the plasma tail (blue) and the dust tail (yellow- white) are visible. Picture by Rhemann [1].

« Progress in astrochemistry rests upon a diversity of observational, experimental, and theoretical skills and a broad knowledge of chemistry and astronomy. »

– A. Dalgarno

T his thesis is the result of multi-disciplinary work overlapping the fields of cometary science, mass spectrometry, computational chemistry and astrochemical modelling. The purpose of this chapter is to provide the reader with the nec- essary introduction required to grasp the goals of the thesis, rather than being a comprehensive review of comets and modelling. Many additional references are given for further reading, while each subsequent chapter will have its own short introduction to its subject matter. This chapter will first cover the sub- ject of cometary science, detailing the formation and composition of comets as well as explaining how they interact with the solar radiation and the solar wind.

This interaction triggers the birth of cometary activity and the generation of the gas and dust cometary atmosphere (or coma). A short introduction to the European Space Agency Rosetta mission to comet 67P/Churyumov-Gerasimenko (67P/C-G) is then given.

Next, the essentials of astrochemical modelling are

provided to the reader, before detailing the current state of the detections of O 2 in the solar system and the interstellar medium (ISM), as well as its probable origin in comets.

(Watermark on the front page of PART I reproduced from Ip [2])

Outline

1.1 A brief introduction to comets 4

1.2 The Rosetta mission 10

1.3 A brief introduction to coma modelling 12

1.4 O 2 in astrophysical environments 13

3

1.1 A brief introduction to comets

p Formation and origin y Comets are celestial objects that were formed around 4.6 billion years ago, as the molecular cloud that would eventually result in the solar system collapsed on itself (Figure 1.2). This grav- itational collapse resulted in the formation of a protosolar nebula, and later a protoplanetary disk due to the rotation of the gas and dust. The accretion of dust grains in the protoplanetary disk gradually led to the formation of planetesimals that eventually coalesced to form the current planets of the solar system. Bodies that accreted beyond the snowline ² of the sun contain icy material. Part of these icy bodies were ejected towards the outer reaches of the solar system after their orbits were perturbed by the giant planets, even- tually leading the formation of the Oort cloud spanning from about 20.000 to 150.000 AU from the sun. The remaining icy bodies on far orbits were grouped together as the orbits of the giant planets shifted, resulting in the formation of the Kuiper belt between 30 and 50 AU from the sun [3]. These two regions, the Oort cloud and the Kuiper belt, are the two main reservoirs of comets in the solar system. The existence of the spherical Oort cloud was postulated in order to explain the random inclination of the orbits of long period comets [4,5].

These long period comets are deviated towards the sun as the result of a gravitational perturbation of neigh- bouring stars. The existence of the Kuiper belt was postulated shortly after the discovery of Pluto, one of the largest Kuiper belt objects [6]. Comets ejected from the Kuiper belt due to a gravitational perturbation of one of the giant planets have shorter periods and have less elongated orbits [7]. Due to their smaller orbits, they remain captured by the solar gravitation field, which is not the case of all long period comets.

Figure 1.2: Simplified schematic representation of the steps of solar system formation. Top left: the presolar nebula or molecular cloud starts to collapse on itself due to gravitation. A protostar is formed from this collapse and the matter cloud surrounding the protostar forms the protostellar nebula. Top right: Protoplanetary disk formed by the rotation of the protosolar nebula. Bottom left: due to gravity, planetesimals start accreting mater and form gaps in the protoplane- tary disk. The outer edges of the disk contain many icy bodies, among which are the comets that will later be ejected to

the Kuiper belt and Oort cloud. Bottom right: fully formed stellar system. Image credit: Bill Saxton, NSF/AUI/NRAO.

Due to their low mass, comets have undergone little internal heating since the time of their formation. Fur- thermore, they have since resided in the cold outer reaches of the solar system. As a result, comets are among the most pristine bodies in the solar system. However, the cometary grain samples returned by the

2 i.e. the radial distance beyond which water remains frozen.

1.1 A brief introduction to comets 5

Stardust mission in 2006 have made it increasingly clear that cometary material has undergone more process- ing than what was previously assumed and probably originates from a wide range of heliocentric distances [8]. Although some degree of thermal and/or radiative processing of cometary material therefore inevitably occurred between the formation of the icy dust grains that compose them and their ejection to the Oort cloud or Kuiper belt, the similarities between the composition of cometary and interstellar dust grains bear witness to the fact they have retained some compositional and structural characteristics from the solar nebula stage of the solar system [3,9]. Comets therefore remain some of the more pristine material in the solar system, and cometary ices in particular must either have remained frozen from the solar nebula or have formed beyond the snowline in the early protoplanetary disk phase [10].

p The composition of comets y Comets are bodies that are composed of a mixture of refractory material and ices. In the second half of the 20 ^th century, comets were viewed as consisting mainly of ice among which refractory grains are trapped, after the image depicted by Whipple [11], hence the given analogy of comets being “dirty iceballs”. However, this image has been questioned by more recent results of the Giotto (1985), Deep Impact (2005) and Rosetta (2014–2016) missions among others to comets 1P/Halley, 9P/Tempel 1 and 67P/C-G respectively. Keller [12] has first suggested that comet 1P/Halley might be closer to an “icy dirtball”

than a “dirty iceball” based on the determination of a dust to ice ratio larger than one. This image has been confirmed over the years by a much larger than unity dust to ice ratio observed at comet 9P/Tempel 1 by the Rosetta spacecraft after the Deep Impact projectile collided with it [13,14]. More recently, the Rosetta space- craft also measured dust to ice ratios between 2.4 and 8 at comet 67P/C-G [9,15–18]. Furthermore, very little exposed surface ice was detected on comet 67P/C-G [19,20]. This is the current vision of comets, i.e. solid refractory objects from which ice sublimates somewhere below the surface.

The refractory fraction of comets has been revealed by the Stardust mission results to be very diverse, consist- ing in various types of chondrules as well as forsterite, calcium-aluminum inclusions or silicates among many others (for a comprehensive review, see Brownlee [8]). The composition of cometary ices has been studied for decades through the analysis of the composition of the gas in the cometary atmosphere – or coma – by both remote sensing and in situ measurements during space missions. The four main cometary volatiles are H 2 O, CO 2 , CO and O 2 , but many more secondary species were detected in the atmospheres of comets [21–23].

Figure 1.3 shows the coma species detected in comets through remote sensing prior to the Rosetta mission.

Since then, other coma molecules have been detected through in situ measurements at comet 67P/C-G, among which the noble gases Ar [24], Kr [25] and Xe [26], molecular nitrogen [27] and oxygen [28], the hydrogen halides HF, HCl and HBr [29], chloromethane [30], sulfur-bearing species [31] and glycine [32].

p The solar wind and the interplanetary magnetic field y The solar wind is a flux of ions and electrons ejected from the hot solar corona. This stream of particles permeates the solar system and can interact with the molecules of the coma, which is why it is important to understand its fundamental characteristics. The com- position of the solar wind, though not exactly identical to it, roughly reflects the elemental composition of the sun, i.e. mostly H ⁺ and about 5 % He ²⁺ as well as heavier highly ionised atoms that together account for about 0.1 % of the solar wind [33]. At a heliocentric distance of 1 AU, the solar wind moves at a speed of roughly 400 km · _s ^-1 . At this distance, the solar wind moves almost at its asymptotic speed, meaning that the solar wind velocity will not be much different at several AU [34]. The solar wind density at 1 AU ranges from 10 ^-1 to 10 ² cm ^-3 [34]. Due to radial expansion, this density decreases with the heliocentric distance.

Convection of material in the sun gives birth to the solar magnetic field, or interplanetary magnetic field (IMF).

This field modulates the trajectories of the charged particles of the solar wind, which in turn carry the magnetic

field into interplanetary space. The IMF lines do not close before the edges of the heliosphere [35]. Figure 1.4

(left) illustrates the difference between a dipolar magnetic field and the solar magnetic field deformed by the

ions and electrons of the solar wind. As can be seen in this figure, the IMF can be well approximated by a

dipolar field close to the surface of the sun, but deviations occur rapidly due to the interaction with the solar

wind. Because the magnetic field lines are carried away by the solar wind particles, the region between the

Figure 1.3: Relative production rate ranges measured for volatile molecules in cometary atmospheres through remote sensing. Black and cyan bars indicate the lower and upper bounds that were measured. The values on the right indicate

the number of measurements made for each species. Reproduced from Bockelée-Morvan and Biver [23].

field lines in Figure 1.4 (left) marks the inversion of the magnetic field. This region is called the magnetic equa- tor or current sheet [34]. The magnetic equator is not necessarily parallel to the ecliptic plane (see Figure 1.4 (right)). As a result, the current sheet describes a spiral (called Parker’s spiral) as the sun rotates and the solar wind expands. This spiral is represented in Figure 1.5 [33,36].

p The birth of cometary activity and solar wind interaction y The birth of cometary activity occurs when

comets arrive close enough to the sun to trigger ice sublimation (i.e. when they cross the snow line). At

this heliocentric distance, cometary ices begin to sublimate through the influence of solar radiation, creat-

ing the gas coma. The resulting gas drag lifts dust grains up from the surface, creating the dust coma. The

cometary outgassing is not uniform and depends on the surface composition and illumination [38]. Addition-

ally, localised outbursts can happen when e.g. large patches of surface ice are suddenly exposed to direct

solar radiation, resulting in increased expulsion of gas and dust [39,40]. While solar radiation is the driver

of cometary activity, the interaction between the resulting gas and dust atmosphere and the solar wind is

responsible for the characteristic cometary tails that give comets their name ( κ _o µ´ ητη ς – “kometes” meaning

wearing long hair), shown in Figure 1.1.

1.1 A brief introduction to comets 7

Figure 1.4: Left: Solar magnetic field lines (solid lines) compared to dipolar magnetic field lines (dashed lines). The axis units are numbers of solar radii. Reproduced from Russell [34]. Right: Schematic view of the inclination of the magnetic

equator compared to the ecliptic plane. Adapted from Russell [34].

Figure 1.5: Parker spiral described by the current sheet of the solar magnetic field as a result of the rotation of the sun and the influence of the solar wind. Reproduced from NASA’s Cosmicopia [37].

p The plasma tail y The first comet tail, called type I or plasma tail, is composed of ionised species. These species are produced either by photoionisation or by charge exchange reactions between neutrals in the coma and solar wind particles. These ions are then picked up by the IMF and are carried out away from the sun [35].

As soon as a neutral species is ionised through one of these processes, it acquires a gyration motion around the IMF lines that are dragged past the comet by the solar wind flow in the anti-sunward direction. However, the plasma tail does not point exactly radially away from the sun. The angle separating the plasma tail from the sun–comet axis is called the aberration angle. This aberration angle results from the combination of the cometary and the solar wind velocity vectors [33]. Incidentally, it is the observation of these always anti- sunward pointing plasma tails and the analysis of their aberration angles that led Biermann [41] to postulate the existence of a steady flux of particles emanating from the sun, i.e. the solar wind. The blue hue of the plasma tail is due to the emission of the abundant CO ⁺ ion [42].

p The dust tail y The second comet tail, called type II or dust tail, is composed of the dust grains that are

released upon the sublimation of ices in the nucleus. The initial gas drag responsible for the ejection of the

grains in the coma quickly fades away as the coma expands. The grains are then left to drift away from the nu-

cleus because of the lack of substantial cometary gravity and are left behind in the elliptical trail of the comet

(although it has been shown that a small fraction of the larger grains and pebbles could remain gravitation-

Figure 1.6: Gradual draping ((a) to (d)) of the interplanetary magnetic field lines around a comet as a result of the dia- magnetic cavity. Reproduced from Ip [2].

ally bound to the nucleus [43]). The dust grains are then only subjected to the solar gravity and the radiation pressure of the solar photons, i.e. these photons transfer part of their momentum to the grains. The mass spread of the grains means they will be more or less affected by these forces, only the smaller grains being effectively deviated by the solar radiation pressure. This deviation is responsible for the broadening of the dust tail compared to the plasma tail visible in Figure 1.1. Its yellow-white colour is due to the scattering of the solar radiation by the grains [35,44]. Both types of cometary tails can extend millions of kilometres into space.

p The diamagnetic cavity y As more and more cometary ions produced are picked up by the IMF, the solar wind velocity decreases because the speed of cometary neutrals and ions is much smaller than the speed of the solar wind (< 1 km · _s ^-1 compared to several hundreds) [44]. This process is known as the mass loading of the solar wind since the cometary ions are for the most part much heavier than the solar wind ions. The added mass is accompanied by a velocity decrease per principle of conservation of momentum. If the coma is suffi- ciently dense, this mass loading process results in the solar wind being stopped altogether at some point. The magnetic field lines end up draping around the shielded inner coma since the mass loading is more important closer to the sun–comet axis than near the edges of the nucleus [45] (Figure 1.6). Though the nuclei of comets are not inherently magnetised [46], the electric current resulting from the gyration motion of cometary ions and electrons generates an induced cometary magnetic field that locally opposes the IMF, resulting in a region of the inner coma where the magnetic field amplitude is essentially zero, called the diamagnetic cavity [47].

The size of this cavity depends on the intensity of the IMF, which depends on the heliocentric distance, as well as on the density of the coma that varies from one comet to another. As a result, its size is very variable: a dia- magnetic cavity with a radius of several thousands of kilometres has been observed around comet 1P/Halley during the Giotto flyby [48,49], while a comparatively very small diamagnetic cavity with a radius of 40 to 400 kilometres has been sporadically observed around the much more weakly outgassing comet 67P/C-G [50,51].

p Distributed sources in the coma y A distributed source in the coma has been defined by Cottin and Fray [52]

to be “an additional source of a gaseous species being produced in the coma from the grains”. ³ Distributed sources have first been proposed to explain the radial distributions of CO and H 2 CO in the coma of 1P/Halley that could not be explained by nucleus sources only [53,54]. Both molecules have also been determined to originate from a distributed source in comets C/1995 O1 (Hale-Bopp), C/2012 F6 (Lemmon) and C/2012 S1 (ISON) [55–57], and CO also exhibited a distributed source in comet C/1996 O2 (Hyakutake) [58]. Besides these two molecules, other species presenting distributed sources in one or several comets have been detected, such as HNC [59], the CN and C 2 radicals [60], several sulphur-bearing compounds [52] and the hydrogen halides HF and HCl [61]. The origin of the distributed source of formaldehyde and HNC in comets has been attributed to

3 In previous literature, the term “extended source” has also widely been used to describe this process.

1.1 A brief introduction to comets 9

the degradation of polyoxymethylene (POM (CH 2 O)n), a polymer composed of formaldehyde monomers, and hexamethylenetetramine (HMT) respectively [52,62]. The same mechanism has been proposed to explain the origin of distributed sources of C 2 or CN [52,59,60]. A potential discovery of such large molecular compounds in cometary grains would therefore strengthen this hypothesis.

The analysis of grains captured by the COSIMA ⁴ instrument during the Rosetta mission has revealed they con- tain high molecular weight C- H-, N- and O-bearing compounds that could fit polymers capable of generating a distributed source of H 2 CO or HNC [64]. Additionally, the mass spectra measured at the surface of comet 67P/C-G by the Ptolemy instrument aboard the Philae lander [65] have been reported to be compatible with the presence of POM, though the observed mass peaks probably originate from a mixture of CHO-bearing organic compounds [66]. More recently, Altwegg et al. [67] have reanalysed the Ptolemy mass spectra in light of the fortuitous capture of a dust grain in the ionisation chamber of the high resolution mass spectrometer ROSI- NA/DFMS [68] and have concluded that the mass peaks assigned to the formaldehyde trimer by Wright et al.

[66] were more likely due to toluene which was identified in the ROSINA/DFMS spectra. Regardless, mixtures of CHNO-bearing materials in the refractory fraction of comet 67P/C-G were detected in the data retrieved by 4 different mass spectrometers, COSIMA, Ptolemy, ROSINA/DFMS and COSAC ⁵ [64,66,67,70]. It has been shown that any mixture of such compounds could probably produce distributed sources of formaldehyde, CO or HNC in cometary comae, no matter their precise composition [67,71]. This being said, Rubin et al. [72] have shown through modelling that the radial distribution of CH 2 O in comet 1P/Halley can be reproduced without any contribution from distributed sources in the coma by assuming temporal variations in the outgassing rate of the comet, which raises the question of whether these observed distributed sources are a result of activity changes rather than the products of degradation of large molecules in the grains.

In the case of the observed distributed source of hydrogen halides, it is more likely they originate from re- maining icy mantles around the dust grains lifted off by the gas outflow of comet 67P/C-G that sublimate further as the grains move in the coma, as is discussed by De Keyser et al. [61]. This hypothesis has been sparked by the observation of icy dust aggregates that continue to sublimate in the coma 67P/C-G by Gicquel et al. [73] and Agarwal et al. [74]. This is also consistent with the tentative detection of a distributed source of water accompanying that of the hydrogen halides, though it is much more difficult to isolate accurately given its small contribution to the overall water content in the coma [61].

p Remote sensing and in situ observations y Remote sensing of comets passing near the sun is a powerful tool to study both the composition [23] and dynamics of comets [75]. Remote sensing measurements can be taken without the decades necessary to plan a space mission, and are also essential as a support to such missions [76,77]. This being said, in situ missions are indispensable in order to acquire additional information.

For instance, the ROSINA/DFMS mass spectrometer aboard Rosetta was able to detect homonuclear diatomic molecules that can difficultly be measured remotely such as N 2 [27] or O 2 [28]. These molecules furthermore have huge background signals from the terrestrial atmosphere when using Earth-based telescopes. Other examples are the determination of the internal structure of the nucleus that was studied using the CONSERT radars on the Rosetta orbiter and the Philae lander [78] or the return of cometary refractory samples to Earth by the Stardust spacecraft [8]. A brief overview of the recent ESA Rosetta mission to comet 67P/C-G is given in the following subsection.

4 COSIMA is the high resolution time-of-flight secondary ion mass spectrometer for the analysis of cometary dust particles flown aboard the ESA Rosetta spacecraft [63].

5 COSAC is the Cometary Sampling and Composition Experiment flown aboard the Philae lander [69].

1.2 The Rosetta mission

The Rosetta mission is one of the European Space Agency Horizon 2000 cornerstone missions that hardly needs introducing [79]. After a delayed takeoff, initially planned for 2003 with comet 46P/Wirtanen as its destination, the Rosetta satellite was launched in 2004. Arriving at its new target, comet 67P/Churyumov- Gerasimenko (67P/C-G), implied a long series of orbital manoeuvres, necessitating three gravitational assists from the Earth and one from Mars in order to reach close to the orbit of Jupiter in 2011. At this distance from the sun, the energy received by the solar panels of the satellite was no longer sufficient to sustain nominal activity and it was therefore decided to place the orbiter in a hibernation state until solar conditions would be more favourable as Rosetta would be pulled back closer to the sun. The spacecraft woke up from its hi- bernation as planned in January, 2014. Figure 1.7 illustrates the long journey undertaken by Rosetta to arrive at its target that will briefly be described below.

Figure 1.7: Schematic overview of the journey of Rosetta in the inner solar system from 2004 to 2016 [80].

p Asteroid flybys y En-route to comet 67P/C-G, Rosetta flew by two selected asteroid targets, 2867 Steins in

2008 and 21 Lutetia in 2010 [81]. The unique observation opportunities these flybys represented by themselves

will not be discussed as they are not the focus of this thesis. A recent review of the scientific results obtained

at asteroids Steins and Lutetia can be found in Barucci and Fulchignoni [82]. These flybys furthermore pre-

sented the perfect opportunity to use the orbiter instrument suite for the first time in space. In the case of

ROSINA/DFMS, the mass spectrometer was not sensitive enough to detect the exosphere of Lutetia [83], but

both flybys were excellent opportunities to assess the influence of the spacecraft background even after sev-

eral years in space [84,85].

1.3 A brief introduction to coma modelling 11

p _Philae y Rosetta entered the cometary atmosphere (or coma) during the month of August 2014 after a se- ries of complex manoeuvres to approach the comet within 100 km. The first priority was then to map out the cometary surface in order to be able to select the ideal landing spot for the Philae lander [86]. The lander was successfully deposited on the cometary surface on November 12 ^th , 2014, though it ended up in a different location than was planned due to a harder than predicted surface crust and the failure of the Philae harpoon to stabilise the lander on the surface. As a result, the position of the lander was unclear and its orientation on the surface prevented its antenna to communicate effectively with the orbiter and its solar panels to receive sufficient energy. Consequently, only few results were ever able to be transferred back to Earth. Nonetheless, this landing was unprecedented and delivered the first in situ measurements taken on a cometary surface.

The lander was eventually spotted on an OSIRIS ⁶ image less than a month before the end of the mission.

Philae was overall a scientific, engineering and flight dynamics success and the results obtained by Philae are summarised by Bibring et al. [88] and Boehnhardt et al. [89].

p Comet escort to perihelion and beyond y After the delivery of Philae to the cometary surface, the Rosetta orbiter followed comet 67P/C-G as activity increased closer to perihelion on August 13 ^th , 2015 and beyond as the mission was extended from its nominal end of mission to September 30 ^th , 2016, at which time it was softly crashed into the comet. During this unprecedented long period of continued in situ observation of a comet, Rosetta followed 67P/C-G along its entire range of cometary activity and executed various orbital manoeuvres to accomplish its scientific objectives. Figure 1.8 shows both the heliocentric and cometocentric distances of Rosetta throughout the mission. In particular, the cometocentric distance time profile reflects the different mission phases, e.g. the close orbits for mapping purposes in October 2014, close flybys in February and March 2015, far excursions in September 2015 and March 2016 or the close elliptical orbits near end of mission. An overall review of the many successes of the orbiter is again beyond the scope of this thesis and is given by Taylor et al. [90].

1.5 2 2.5 3 3.5

Heliocentric distance (AU)

10 100 1000

0 100 200 300 400 500 600 700

Cometocentric distance (km)

Days since 01/08/2014

Figure 1.8: Top panel: heliocentric distance of comet 67P/C-G in astronomical units (AU) from August 1 ^st , 2014 to September 30 ^th , 2016. Bottom panel: cometocentric distance of Rosetta in km for the same time period. The x axis represents the number of days passed since August 1 ^st , 2014. Vertical gray lines mark the limits between months and thick lines indicate

the passing of a year.

6 OSIRIS is the Optical, Spectroscopic, and Infrared Remote Imaging System aboard the Rosetta orbiter [87].