11institutetext: 1Department of Physics, Imperial College London, Prince
Consort Road, London SW7 2AZ, United Kingdom
2Department of Physics, Umeå University, 901 87 Umeå, Sweden
3Physikalisches Institut, University of Bern, Sidlerstrasse 5, CH-3012 Bern,
Switzerland
4Center for Space and Habitability, University of Bern, Gesellschaftsstrasse
6, CH-3012 Bern, Switzerland
5LATMOS, 4 Avenue de Neptune, F-94100 Saint-Maur, France
6Department of Climate and Space Sciences and Engineering, University of
Michigan, 2455 Hayward, Ann Arbor, MI 48109, USA
7Royal Belgian Institute for Space Aeronomy (BIRA-IASB), Ringlaan 3, B-1180,
Brussels, Belgium
8Institute of Computer and Network Engineering (IDA), TU Braunschweig, Hans-
Sommer-Straße 66, D-38106 Braunschweig, Germany
9Space Science Division, Southwest Research Institute, 6220 Culebra Road, San
Antonio, TX 78228, USA
10University of Texas at San Antonio, San Antonio, TX

# ROSINA ion zoo at Comet 67P

A. Beth1,2 K. Altwegg3,4 H. Balsiger3 J.-J. Berthelier5 M. R. Combi6 J. De
Keyser7 B. Fiethe8 S. A. Fuselier9,10 M. Galand1 T. I. Gombosi6 M. Rubin3 T.
Sémon3

(Received 24 September 2019; Accepted 7 July 2020)

###### Abstract

Context. The Rosetta spacecraft escorted Comet 67P/Churyumov-Gerasimenko for 2
years along its journey through the Solar System between 3.8 and 1.24 au.
Thanks to the high resolution mass spectrometer on board Rosetta, the detailed
ion composition within a coma has been accurately assessed in situ for the
very first time.

Aims. Previous cometary missions, such as Giotto, did not have the
instrumental capabilities to identify the exact nature of the plasma in a coma
because the mass resolution of the spectrometers onboard was too low to
separate ion species with similar masses. In contrast, the Double Focusing
Mass Spectrometer (DFMS), part of the Rosetta Orbiter Spectrometer for Ion and
Neutral Analysis on board Rosetta (ROSINA), with its high mass resolution
mode, outperformed all of them, revealing the diversity of cometary ions.

Methods. We calibrated and analysed the set of spectra acquired by DFMS in ion
mode from October 2014 to April 2016. In particular, we focused on the range
from 13-39 u·q-1. The high mass resolution of DFMS allows for accurate
identifications of ions with quasi-similar masses, separating 13C+ from CH+,
for instance.

Results. We confirm the presence in situ of predicted cations at comets, such
as CH${}_{m}^{+}$ ($m=1-4$), HnO+ ($n=1-3$), O+, Na+, and several ionised and
protonated molecules. Prior to Rosetta, only a fraction of them had been
confirmed from Earth-based observations. In addition, we report for the first
time the unambiguous presence of a molecular dication in the gas envelope of a
Solar System body, namely CO${}_{2}^{++}$.

###### Key Words.:

Comets: individual: 67P/Churyumov-Gerasimenko; Plasmas; Molecular processes

##  1 Introduction

Relatively small, with a nucleus size of a few kilometres to a few tens of
kilometres, comets are only detectable once they are close enough to the Sun
and display a bright tail. Compared to other planetary bodies and their
atmosphere, the gas envelope of comets, the coma, behaves very differently.
The coma results from the sublimation of ices near the nucleus’ surface, which
then undergoes an acceleration to several hundreds of m·s-1, continuously
replenishing the coma. Mainly made of water, the coma contains a diversity of
neutral species, such as CO2, CO (Krankowsky et al., 1986; Hässig et al.,
2015), and many others (e.g. Le Roy et al., 2015) that have been detected in
situ at 1P/Halley (hereinafter referred to as 1P) and 67P/Churyumov-
Gerasimenko (hereinafter referred to as 67P, Churyumov & Gerasimenko, 1972).
Extreme ultraviolet (EUV) solar radiation penetrates and ionises the neutral
gas envelope, giving birth to the cometary ionosphere. In addition to EUV, an
additional source of ionisation is energetic electrons (Cravens et al., 1987).
Depending on the local neutral number density, newborn cometary ions may
undergo collisions with neutrals, yielding the production of cations which
cannot result from direct ionisation of the neutrals. The diversity of ions is
therefore richer than that of neutrals.

Cometary ions may be observed remotely at ultraviolet and visible wavelengths.
Emissions in these wavelengths arise mainly from the resonant fluorescence of
sunlight. These types of emissions from cometary molecular ions were first
observed at Comet C/1907 L2 (Daniel) (Deslandres & Bernard, 1907; Evershed,
1907). Although the emitting species was unknown at the time of the detection
(Larsson et al., 2012), it was later identified as CO+. The discovery of an
ion tail that is always oriented anti-sunward led to the discovery of the
solar wind (Biermann, 1951; Parker, 1958). Several cometary ions have since
populated the list: N${}_{2}^{+}$ and CH+ (Swings, 1942), CO${}_{2}^{+}$ and
HO+ (Swings & Page, 1950; Swings & Haser, 1956), Ca+ (Preston, 1967), H2O+
(Herzberg & Lew, 1974), CN+ (Lillie, 1976), and H2S+ (Cosmovici & Ortolani,
1984). It is important to note that some ions are detected in cometary
environments through observations in EUV and X-Rays (Lisse et al., 1996).
Nevertheless, they are not cometary as we define here. The emission originates
from the de-excitation of multiply-charged ions (e.g. O6+), produced after a
charge exchange between the high charge state solar wind ions (e.g. O7+) and
the neutral coma (Lisse et al., 2004). In addition, cometary ions may be
detected at radio wavelengths. Emissions at these wavelengths are triggered by
transitions between ro-vibrational or hyperfine levels (Crovisier & Schloerb,
1991). The first ion detected through radio-astronomy observations was HCO+
(Veal et al., 1997), followed by H3O+ and CO+ (Lis et al., 1997); all of them
were at Comet C/1995 O1 (Hale-Bopp). Remote sensing detection of ionised
constituents remains limited in terms of spatial resolution, species, and
number density.

The most efficient way to probe the neutral and ion composition within a coma
is in situ observations such as those performed by the European Space Agency’s
Giotto (Reinhard, 1986) and Rosetta (Glassmeier et al., 2007) missions. At
comet 1P, Giotto carried several instruments able to probe different mass and
energy ranges of the cometary ions. The first measurements were performed by
the Giotto Neutral Mass Spectrometer (NMS, Krankowsky et al., 1986) when
operated in the ion mode, which consisted of two analysers: a double-focusing
mass spectrometer (M-analyser, for the range 1-37 u·q-1) and an electrostatic
energy spectrometer (E-analyser, for the range 1-56 u·q-1 in ion mode).
Krankowsky et al. (1986) reported the ‘quasi-unambiguous’ identification of
C+, CH+, O+, HO+, H2O+, H3O+, Na+, C${}_{2}^{+}$, S+, 34S+, and Fe+. In
addition, peaks were detected at all integer mass positions from 12 to 37
u·q-1, except at 22 u·q-1. A second Giotto instrument was the ion mass
spectrometer (IMS, Balsiger et al., 1986) consisting of two sensors: the high-
energy-range spectrometer (HERS, dedicated to the study of ion composition and
velocity outside the contact surface, covering the mass-per-charge range 1-35
u·q-1) and the high-intensity spectrometer (HIS, dedicated to measurements
inside the contact surface, covering 12-57 u·q-1). Balsiger et al. (1986)
reported the detection of C+, CH+, CH${}_{2}^{+}$+N+, CH${}_{3}^{+}$, O+, HO+,
H2O+, H3O+, CO+, and S+. The third instrument was the Positive Ion Cluster
Composition Analyzer (PICCA, Korth et al., 1987) part of the Rème Plasma
Analyser (RPA, Rème et al., 1987) which measured ions from 10 u·q-1 up to 203
u·q-1. As the E-analyzer in NMS, PICCA is not a true mass spectrometer as it
separates ions in terms of energy-per-charge instead of mass-per-charge.
However, as cometary ions are cold (i.e. their thermal speed is small compared
with the ram speed of the spacecraft during the flyby), ions are collimated in
the ram direction and their energy is roughly proportional to their mass
allowing to deduce a mass spectrum with a limited mass resolution. Korth et
al. (1986) reported the presence of ions belonging to the H2O group (O+, HO+,
H2O+, and H3O+), along with ions from CO+, S+, and CO${}_{2}^{+}$ groups. In
addition, spectra revealed periodic peaks in terms of u·q-1, around 64, 76,
94, and above 100 u·q-1. It was unclear if they were associated with ions from
the Fe group, sulphur compounds, or hydrocarbons. Huebner et al. (1987) and
Mitchell et al. (1987) suggested that these observations are consistent with
the dissociation of polyoxymethylene (CH2O)n. Later on, Mitchell et al. (1992)
showed that these peaks, repeating every $\sim 15$ u·q-1, correlate with the
number of combinations of C, H, O, and N, to form a molecular ion at a given
u·q-1 (see Fig. 2 in Mitchell et al., 1992). Similar patterns are observed at
Titan (Vuitton et al., 2007), though the contribution of O-bearing molecules
is negligible compared with that at comets. In addition, they observed a
predominance of odd mass number ions.

The term ‘unambiguous detection’ or similar formulation should be taken with
great care for Giotto data at 1P since the mass resolution of its instruments
was about $\Delta m\sim 1$ u. The combination and assemblage of the primary
blocks, C, H, O, and N atoms, to build more complex molecules are limited at
low masses (typically below 25 u, Mitchell et al., 1992). At some specific
values of u·q-1, there exists only one combination: C+ (12 u·q-1), CH+ (13
u·q-1), H3O+ (19 u·q-1), C${}_{2}^{+}$ (24 u·q-1), C2H+ (25 u·q-1), if one
disregards isotopes and isotopologues. There is even no candidate between 20
u·q-1 and 23 u·q-1. At other u·q-1 (in particular 18 u·q-1 which corresponds
to H2O+ and NH${}_{4}^{+}$), photo-chemical models are needed to infer the
relative contribution of each ion, or, conversely, constrain the neutral
composition (Haider & Bhardwaj, 2005).

At comet 67P, the Rosetta orbiter carried two instruments which performed a
true mass analysis of the ambient ions: the Ion Composition Analyzer (ICA,
Nilsson et al., 2007), part of the Rosetta Plasma Consortium (RPC, Carr et
al., 2007), and the Double Focusing Mass Spectrometer (DFMS), part of the
Rosetta Orbiter Spectrometer for Ion and Neutral Analysis (ROSINA, Balsiger et
al., 2007). Although we may compare RPC-ICA with its homologue RPA-PICCA, the
former suffers from limitations to probe cold cometary ions. As Rosetta was
moving slowly with respect to the ambient plasma, ions were not collimated
along the spacecraft velocity and RPC-ICA has a wide field of view. In
addition, its minimum energy acceptance is 4-5 eV, such that it only observed
ion species after they were energised either as pick-up ions or accelerated by
the spacecraft potential prior to entering RPC-ICA. Nevertheless, RPC-ICA was
perfectly designed for probing the energetic solar wind ions, such as H+,
He++, and He+, unlike ROSINA-DFMS. ROSINA-DFMS (described in Section 2) has
the ability to probe neutrals as well as ions with two different mass
resolutions either $m/\Delta m\approx 500$ in the ‘Low Resolution’ LR mode or
$m/\Delta m>3000$ in the ‘High Resolution’ HR mode. DFMS is the most powerful
spectrometer in terms of mass resolution ever flown on board a spacecraft so
far. Previous analyses of DFMS ion spectra in high resolution revealed the
unambiguous detection of H2O+, NH${}_{4}^{+}$, and H3O+ (Fuselier et al.,
2016; Beth et al., 2016). Low resolution spectra have been also analysed and
highlighted the presence of other species either at large heliocentric
distances (Fuselier et al., 2015) or near perihelion (Heritier et al., 2017)
with the support of photo-chemical modelling.

In this paper, we present in situ detections of cometary ions at 67P over the
range 13-39 u·q-1 in high resolution and 13-141 u·q-1 in low resolution. In
HR, DFMS pinpointed the mass-per-charge ratio of impinging cometary ions with
such a high accuracy that their composition and identity can be ascertained
without any ambiguity. The DFMS spectrometer and data processing are presented
in Section 2, followed by a review of the mass spectra acquired during the
period Oct 2014-Apr 2016 in Section 3. Section 4 highlights the main results
including the key different ion family behaviours (4.1), the protonated
molecules (4.2), water isotopologues (4.3), and dications (4.4). Discussion
and conclusions are presented in Section 5.

##  2 ROSINA-DFMS

###  2.1 Ion Mode: principle, data processing, and limitations

The description of the instrument and its capabilities have been provided in
Balsiger et al. (2007) and Le Roy et al. (2015). In the ion mode, the ionised
constituents are directly admitted in the ion optics from which they exit
towards the detector. Surrounding the entrance of the instrument, a
negatively-biased grid was designed to attract ions in case of a positive
spacecraft potential as anticipated from simulations based on plasma
conditions in the dense cometary ionosphere encountered by Giotto. However,
the spacecraft potential of Rosetta was very negative during most of the
escort phase (Odelstad et al., 2017) and the grid was permanently set to a
small negative potential, of -5 V. Once inside the instrument, ions are
accelerated significantly by a large negative potential so that their energy
in the ion optics is much higher than their energy at the entrance of the
instrument: they undergo a first deflection in the electrostatic energy
analyser which selects the ion energy before they exit through either the LR
or the HR energy slit, the former being 6.5 times wider that the latter, into
the magnetic analyser where they are deflected according to their mass and
charge. Exiting the magnetic analyser, ions impinge on the detector which
consists of a Micro Channel Plate (MCP) followed by a Linear Electron Detector
Array (LEDA). Since the magnetic field intensity in the magnet varies with the
temperature (see Keyser et al., 2019), the impact position of a given ion on
the detector will depend on the temperature as well. The LEDA is split into
two identical rows (hereafter referred to as channel A and channel B) with 512
pixels, 25-$\mu$m wide and 8-mm long perpendicular to the mean axis of the
row. When an ion hits the MCP, a cascade of electrons is produced and the
total amount of negative charges collected by the LEDA, known as the MCP gain,
depends on the voltage applied to the MCP. The gain is not uniform over the
entire MCP area and, for each pixel, one can define a ‘pixel gain’, which
modulates the average MCP gain, and determine the actual number of electrons
collected by the corresponding LEDA pixel.

The pixel gains vary during the Rosetta escort phase and have been regularly
determined through dedicated in-flight calibrations. Pixel gains degraded
during the mission (Schroeder et al., 2019), especially for pixels located
close to the centre of each row (from pixel 200 to pixel 400), where H2O+ ions
strike in both neutral and ion modes of DFMS. To partially compensate this
degradation and the loss of sensitivity, on the 27th of January 2016, the
post-acceleration was modified in order to move the central pixel $p_{0}$,
such that the position on the detector of the selected mass of each spectrum
was moved forwards, on pixels with a less degraded gain. Pixels at the very
edge of the LEDA rows have a poor gain as well but they are not included in
the analysis.

DFMS has two basic modes of operation. In the ‘neutral’ mode, the neutral
species are ionised and fragmented through electron impact in the ion source,
thanks to a filament emitting electrons at $\sim 45$ eV, before being
accelerated into the ion optics. In the ‘ion’ mode, the filament is not
powered and ions are directly admitted into the ion optics. The ion and
neutral modes are not operated simultaneously but, for both of them, the total
integration time for each individual spectrum is 19.8 s made of 3000 exposures
of 6.6 ms. For both modes, DFMS may operate in Low or High mass-per-charge
Resolution (hereafter referred to as LR and HR, respectively). HR mode, for
which $m/\Delta m>3000$ at the 1% peak height level for 28 u·q-1 (Balsiger et
al., 2007), allows separation of ions that have very close mass-per-charge
ratios (e.g. 13C+ and CH+, H2O+ and NH${}_{4}^{+}$, CO+ and N${}_{2}^{+}$),
which is not possible in LR mode, for which $m/\Delta m\approx 500$. However,
the sensitivity at a given gain step is significantly higher in LR than in HR,
therefore the LR mode was of particular interest during periods of low
outgassing and when fewer ion species can be detected due to limited ion-
neutral chemistry (e.g. Fuselier et al., 2015). By comparison, the HR mode was
of particular interest for periods at high outgassing activity, such as near
perihelion, when ion-neutral chemistry takes place and many new species are
present and need to be separated (e.g. Beth et al., 2016).

A typical sequence of acquisition is as follows. Firstly, the first commanded
(instructed to the instrument when operating) mass-per-charge ratio is 18
u·q-1 both in LR and in HR. Secondly, the second commanded mass-per-charge
ratio is the lowest one which the instrument can perform: 13.65 u·q-1 in LR,
13 u·q-1 in HR. Thirdly, the commanded mass-per-charge ratio is then
incremented: exponentially in LR, linearly in HR, with $m_{0}(i)$ the $i$-th
commanded mass-per-charge ratio $m_{0}$ ($i\geq 2$) defined as:

$m_{0,\text{LR}}(i)\approx 11.27798\times 1.1^{i}\text{ u\textperiodcentered
q}^{-1},$ $m_{0,\text{HR}}(i)=11+i\text{ u\textperiodcentered q}^{-1}.$

Fourthly, the penultimate commanded mass-per-charge ratio: 134.4 u·q-1
($i=26$) in LR, 100 u·q-1 ($i=89$) or 50 u·q-1 ($i=39$) in HR (100 was used as
an upper limit during the first half of the mission but nothing was detected
above 50 u·q-1, this limit was then lowered in July 2015). Finally, the last
commanded mass-per-charge ratio is 18 u·q-1.

The three measurements of 18 u·q-1 during a sequence helped in monitoring the
variability of the ambient plasma conditions and/or the effective DFMS
geometrical factor in the ion mode which depended on the spacecraft potential.
A full sequence lasts between 10 and 20 minutes, depending on the resolution
and the number of commanded mass-per-charge ratios. LR and HR modes differ in
terms of u·q-1 coverage since the mass-per-charge coverage for a given mass-
per-charge ratio $m_{0}$ is roughly $0.1\ m_{0,\text{LR}}$ in LR and $0.016\
m_{0,\text{HR}}$ in HR (see Eq. 1 below). Therefore, successive LR spectra
overlap and cover the full range from 13 to 141 u·q-1. HR successive spectra
may overlap only at high masses from 64 u·q-1 onwards. Finally, less spectra
are required in LR to cover the same mass-per-charge range because several
u·q-1 may be covered in a given spectrum. However, in the latter case, the
peaks fall on different locations on the detector, while, in HR, peaks fall
close to the centre of the detector.

Thanks to the high resolution of DFMS, the ions species presented in this
paper were identified by a detailed and accurate data analysis without the
need to rely on photo-chemical models. The models presented in Section 4 only
aim at understanding the variability of the cations throughout the escort
phase for those confirmed.

###  2.2 Data analysis

The HR mode requires the utmost care for its mass calibration, that is
determining the exact relation between the location of the pixel $p$ on the
detector and the associated mass $m(p)$. This relation is given by (Le Roy et
al., 2015):

$m(p)=m_{0}\exp\left[\dfrac{C}{D}\dfrac{(p-p_{0,m_{0}})}{z_{m_{0}}}\right]\approx
m_{0}\left[1+\dfrac{C}{D}\dfrac{(p-p_{0,m_{0}})}{z_{m_{0}}}\right],$ (1)

where $C=25$ $\mu$m is the centre-to-centre distance between adjacent pixels,
$D=127000\leavevmode\nobreak\ \mu$m the dispersion factor, $p_{0,m_{0}}$ the
location of the commanded $m_{0}$·q-1 on the detector, and $z_{m_{0}}$ the
zoom factor (1 for LR). Eq. 1 is linearisable because the argument inside the
exponential is $\ll 1$. As indicated by their subscript in Eq. 1, both
$p_{0,m_{0}}$ and $z_{m_{0}}$ depend on the commanded mass-per-charge $m_{0}$
and, as aforementioned, on the magnet temperature since the exit location of a
given u·q-1, hence the pixel on the detector, depends on the magnetic field
intensity. If the variation of $p_{0,m_{0}}$ and $z_{m_{0}}$ between two
adjacent u·q-1 (e.g. 18 u·q-1 and 19 u·q-1) are very small, they may be
significant for widely different u·q-1 such as 13 u·q-1 and 40 u·q-1. To
achieve a perfectly accurate spectrum analysis, both parameters should be
reassessed for each sequence of acquisition of DFMS which is possible when two
species, with mass $m_{1}$ and $m_{2}$ located at pixel $p_{1}$ and $p_{2}$
respectively, are present in the same spectrum. The zoom factor $z_{m_{0}}$
can be derived from:

$z_{m_{0}}=\dfrac{p_{2}-p_{1}}{\dfrac{C}{D}\log\left(\dfrac{m_{2}}{m_{1}}\right)},$
(2)

and then $p_{0,m_{0}}$ is inferred from one of the two species from Eq. 1.
Although this procedure may work well in neutral mode, it is seldom applicable
in ion mode since spectra with two well-shaped and separated peaks are only
observed for a few u·q-1 and favourable observation periods such as at 18
u·q-1 and at perihelion. Indeed, in ion mode, the count rates on the detector
are much smaller than in the neutral mode because of the effective geometrical
factor of DFMS for cometary ions, which is lower than that for neutrals due to
several combined factors (e.g. large neutral number density, high ion source
efficiency, ions accelerated by the spacecraft potential). As a matter of
fact, all the ion mode spectra were acquired with the highest gain step to
ensure the maximum sensitivity for the instrument. Following the findings of
De Keyser et al. (2015), we have set the zoom factor $z$ to 5.5 for 13, 14,
and 15 u·q-1 and to 6.4 otherwise. Recent analysis of the spectra in neutral
mode showed that the zoom factor is slightly lower at 13, 14, and 15 u·q-1
confirming that 5.5 is appropriate. For $p_{0,m_{0}}$, we used the value
determined from the most proximate spectrum either at 18 u·q-1 or 19 u·q-1
during the same sequence of acquisition of DFMS, that is either
$p_{0,m_{0}}\approx p_{0,18}$ or $p_{0,m_{0}}\approx p_{0,19}$. Indeed,
spectra at 18 and 19 u·q-1 show strong peaks throughout the escort phase
attributed to H2O+ and H3O+. However, as there is also NH${}_{4}^{+}$ at 18
u·q-1, we preferred to use 19 ($p_{0,m_{0}}\approx p_{0,19}$) to remove any
ambiguity. This approach for deriving $p_{0,m_{0}}$ works well, except for 13,
14, and 15 u·q-1, discussed in Appendix A. $p_{0,m_{0}}$ is less constrained
than $z$ and varies more significantly in comparison. One may evaluate the
uncertainty of the mass $\delta m$ from those of $p_{0}$, $\delta p_{0}$, and
$z$, $\delta z$:

$\dfrac{\delta m}{m}\approx\dfrac{C}{Dz}\left(\dfrac{\delta
z}{z}(p-p_{0})+\delta p_{0}\right).$ (3)

We found that the main source of uncertainty is $\delta p_{0}$. In the dataset
generated by the ROSINA team, the default value for $|\delta p_{0}|$ is set to
10. The reader may find additional information in the ROSINA User Guide. For
the identification in high resolution, we proceeded as follows. Firstly, we
selected a u·q-1-range within which species may be found. As u·q-1 increases,
the range does as well. Secondly, we performed an additional visual inspection
if needed for low counts to remove any suspicious spectrum (e.g. not-flat
spectrum baseline, spurious peak far from any known ion species). Thirdly, we
over-plotted spectra (from a few tens to hundreds, depending on the mass-per-
charge ratio with colour coding which depends on the time of acquisition
through the mission, see Fig. 1). Similar studies may be performed with
different variables (e.g. latitude).

Figure 1: First panel: colour bar used for all spectra in LR and in HR.
Spectra have been acquired between the 30th of October 2014 and the 12th of
April 2016. A separation has been set on 27 January 2016 corresponding to the
time when $p_{0}$ has been voluntarily shifted in DFMS (see text and Appendix
A). Colour bars representing the time coverage of DFMS spectra in LR (second
panel) and HR (third panel) ion mode are also displayed. White means that
sequences of scans are performed on that day and black none. Solstice refers
to the Summer Solstice over the Southern Hemisphere (solar latitude =
$-52^{\circ}$). Fourth panel: heliocentric distance, cometocentric distance,
and local outgassing rate ($\approx n_{n}r^{2}\varv_{n}$) as a function of
time for the period of interest. Black dots correspond to (from left to right)
the inbound Equinox, Perihelion, Solstice, and outbound Equinox. An outgassing
speed $\varv_{n}$ of 1 km·s-1 has been assumed for the outgassing.

In addition to ion identification, one of the main goals is also to assess in
which conditions these ions have been detected: low and high outgassing
activity, close and large heliocentric distance, close and large cometocentric
distance. Because of all of these variables, we decided to colour spectra as a
function of the time of acquisition during the mission. Fig. 1 shows the
colour code used as a function of time for the spectra, the time coverage of
DFMS in LR and HR as well as the heliocentric distance, cometocentric
distance, and outgassing rate with the corresponding colour. Yellow
corresponds to the early phase of the mission, with Rosetta far from the Sun
($>2.2$ au), close to the nucleus ($<50$ km), and 67P with a low outgassing
rate ($Q<10^{27}$ s-1). Orange corresponds to the period before perihelion
with Rosetta close to the Sun ($<2.2$ au), between 100 km and 200 km from the
nucleus, and 67P with an intermediate outgassing rate ($10^{27}<Q<10^{29}$
s-1). Red corresponds to the period after perihelion with Rosetta close to the
Sun ($<2.2$ au), farther than 200 km from the nucleus, and 67P with an
intermediate outgassing rate ($10^{27}\lesssim Q<10^{29}$ s-1). Green and blue
correspond to the period after the pixel shift with Rosetta far from the Sun
($>2.2$ au) and 67P with a low outgassing rate. We strongly advise the reader
to refer to Fig. 1 for the interpretation of the figures in Section 3.

##  3 Ion spectra

###  3.1 Overview

Thanks to its great sensitivity, ROSINA-DFMS allowed probing the ion
composition in LR up to very high masses for the first time. Fig. 2 shows a
series of selected spectra from 13 to 141 u·q-1. Above 72 u·q-1, the mass
calibration is not as good as for lower masses because a different post-
acceleration is applied within the instrument, which explains why peaks are
not centred correctly. The highest counts are recorded at $\sim 18$ u·q-1 and
at $\sim 19$ u·q-1, where H2O+ and H3O+ are found. Other high count regions
are also observed at $\sim 28$ u·q-1 (e.g. CO+) and at $\sim 44$ u·q-1 (e.g.
CO${}_{2}^{+}$). We note that we have gaps, low signals, or non-detections for
instance at 36 u·q-1 and around 51 u·q-1, similar to those showed by Mitchell
et al. (1992), already described in Section 1. However, in contrast, we have
strong peaks at 21 u·q-1 and 22 u·q-1 where no combination of C, H, O, and N
to form a monocation may fit. As the commanded mass-per-charge ratio
increases, the signal-to-noise ratio decreases together with the signal
(physical) and sensitivity (instrumental). Moreover, at high mass-per-charge
ratios, an insidious effect decreases the width of the peak. With a constant
$\Delta m/m$, the mass difference between two successive pixels $m(p+1)-m(p)$
increases with $m_{0}$ such that ions are focused and spread over fewer and
fewer pixels, down to a single pixel in extreme cases. This focusing results
in sharp peaks, with high counts for one pixel (spikes), instead of broad
ones, which may be misinterpreted as ‘ghost’ peaks, that is sharp and spurious
peaks at the location of one pixel with high counts compared with surrounding
pixels. However, over-plotting several spectra reveals that these spikes are
located around each integer mass-per-charge ratio up to 141 u·q-1 and are thus
real. Above 40 u·q-1, the exact species identification cannot be achieved due
to the lack of peaks in HR ion mode as a consequence of the decreased
sensitivity. The following sections are dedicated to the identification of ion
species detected in the range of 13 to 39 u·q-1.

r0.6cmb3.4cmLR r0.6cmb3.4cmLR
r0.6cmb3.4cmLR r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmLR

Figure 2: Concatenation of spectra recorded at each channel in ion low
resolution mode. This covers mass-per-charge ranging from 13 u·q-1 to 141
u·q-1. Dark grey regions represent ranges which are not covered by the
instrument (below 13 u·q-1 and above 141 u·q-1). Light grey regions represent
ranges where two consecutive scans overlap, meaning that these ranges are
covered by the edge of the detector.

###  3.2 Ion mass-per-charge range 13 – 21 u·q-1

r0.6cmb3.4cmLR
r0.6cmb3.4cmHR
r0.6cmb0.33667905cm+3.4cmHR

Figure 3: Stacked individual spectra covering 13 and 14 u·q-1 from the channel
A ($\circ$) and channel B ($\square$) of the MCP, each of them represents the
accumulated counts during 19.8 s at the maximum gain step. The top panel shows
stacked spectra in low resolution covering both ranges (mass-per-charge ratios
are indicated by the long ticks on the upper axis, half-integer ones by the
short ticks). Panels underneath are HR spectra at specific u· q-1: 13 (middle)
and 14 (bottom), corrected from the pixel shift (see Fig. 19). A statistical
vertical error bar of $\pm\sqrt{N}$, where $N$ stands for the number of
counts, is superimposed to the counts for information. The colour coding is
given by the colour bar in Fig. 1 and relies on the time of acquisition during
the escort phase. The mass-per-charge ratios of expected ions from Table 1 are
also indicated and given in App. C, Table 4.

Fig. 3 shows spectra for the range 13-14 u·q-1. In LR, two distinct peaks are
present at each integer. In HR at 13 u·q-1, there are two candidates: 13C+ and
CH+. Once the correction described in Appendix A has been applied, spectra at
mass-per-charge 13 u·q-1 show a very faint signal attributed to CH+ (see Fig.
3, middle). A weak but stronger peak is also visible at 14 u·q-1 (see Fig. 3,
bottom) and attributed to CH${}_{2}^{+}$. There is no evidence for N+. From LR
spectra, it appears relatively difficult to identify the most favourable
periods for the detections of these ions. At 13 u·q-1, even though some peaks
appeared around perihelion, the most favourable conditions seem to be met at
large heliocentric distances (in yellow, blue, and green). One should not be
misguided by the relatively low signal in LR: after the pixel shift on 27
January 2016, the peak 13 u·q-1 shifted to the left edge of the detector such
that the left part of the peak is lost. At 14 u·q-1, there is a similar
behaviour and overall, the highest counts occurred on average at large
heliocentric distances.

r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmHR

Figure 4: Same as Fig. 3, but for 15 u·q-1 only. Stacked spectra in low
resolution (upper) and HR (bottom).

Fig. 4 shows mass-per-charge 15 u·q-1 in LR (top) and HR (bottom). It is one
of the rare commanded mass-per-charge ratio for which LR spectra (with 18
u·q-1 sometimes) only cover one integer in u·q-1. The peak is associated with
CH${}_{3}^{+}$, as seen in HR and its intensity is quite strong compared with
CH+, CH${}_{2}^{+}$ (see Fig. 3), and CH${}_{4}^{+}$ (see Fig. 5). LR and HR
spectra show that the detection is not controlled by cometary conditions, in
particular the outgassing rate. The main reason is that CH${}_{3}^{+}$ is
barely destroyed through ion-neutral reactions with the dominant cometary
neutral species, namely H2O, CO2, and CO, as the corresponding kinetic rates
are $\leq 10^{-11}$ cm3·s-1 (Bates, 1983; Herbst, 1985; Luca et al., 2002).

l0.3115cm+0.6cmb3.4cmLR
l0.3115cm+0.6cmb3.4cmHR
l0.3115cm+0.6cmb0.33667905cm+3.4cmHR

Figure 5: Same as Fig. 3, but for 16-17 u·q-1. Stacked spectra in low
resolution (upper) and HR at 16 u·q-1 (middle) and at 17 u·q-1 (bottom).

Fig. 5 shows spectra for the range 16-17 u·q-1. Three ions are identified at
16 u·q-1 in HR: O+, NH${}_{2}^{+}$, and CH${}_{4}^{+}$. The O+ signal is
stronger prior to spring equinox than near perihelion/winter solstice. At
large heliocentric distances, the source of ions is mainly driven by
ionisation of the neutral molecules by electron-impact (Heritier et al.,
2018). Ion-neutral chemistry is limited or even negligible (Galand et al.,
2016). The major sources of O+ are ionisation of CO2, followed by ionisation
of H2O, based on their respective ionisation rate and volume mixing ratios.
Indeed, although the CO2 abundance is spatially depending on the sub-
spacecraft latitude (Hässig et al., 2015; Gasc et al., 2017), the photo-
ionisation rate yielding O+ is an order of magnitude higher than that of H2O
(Huebner & Mukherjee, 2015). Alongside O+, two other ions are present:
NH${}_{2}^{+}$ and CH${}_{4}^{+}$, while there is no evidence of
13CH${}_{3}^{+}$. As 13CH${}_{3}^{+}$ should slowly react with H2O like
CH${}_{3}^{+}$, if by any chance 13CH${}_{n}^{+}$ ($n=1-4$) would be
detectable, 13CH${}_{3}^{+}$ would be the best candidate. The non-detection of
13CH${}_{3}^{+}$ implies that 13CH+, 13CH${}_{2}^{+}$, and 13CH${}_{4}^{+}$
would not be detected either, which is indeed the case. According to the
isotopic ratio 13C/12C derived by Hässig et al. (2017), 13CH${}_{3}^{+}$
should be at 1% height peak level from 12CH${}_{3}^{+}$, that is about 0.6
counts in the best case, therefore preventing its detection. The
CH${}_{4}^{+}$ signal (Fig. 5) is much weaker than that of CH${}_{3}^{+}$
(Fig. 4), from fivefold to tenfold, and is only detected at large heliocentric
distances. The electron-impact ionisation of CH4 is expected to slightly
favour CH${}_{4}^{+}$ compared to CH${}_{3}^{+}$, as the associated cross
sections are alike (Song et al., 2015), while the ionisation potential is
lower for the production of CH${}_{4}^{+}$ (12.61 eV for CH${}_{4}^{+}$, 14.25
eV for the dissociative ionisation of CH4 into CH${}_{3}^{+}$ at 0 K, Samson
et al., 1989). Assuming that the ionisation of CH4 is the main source of
CH${}_{n}^{+}$ ($n=1-4$), CH${}_{3}^{+}$ however dominates over CH${}_{4}^{+}$
as it almost does not react with H2O, CO, and CO2, as found for 1P/Halley at
0.9 au (Allen et al., 1987). The photo-ionisation rate of CH4 by EUV leading
to CH${}_{4}^{+}$ is roughly twice the corresponding value for CH${}_{3}^{+}$
(Huebner & Mukherjee, 2015). The very high count ratio of CH${}_{3}^{+}$ over
CH${}_{4}^{+}$, especially near perihelion is a clear signature of ion-neutral
chemistry occurring in the coma. Possible other sources of CH${}_{3}^{+}$, and
not of CH${}_{4}^{+}$, are the dissociative ionisation and ionisation
following fragmentation of saturated hydrocarbons (excluding CH4), found at
67P (Schuhmann et al., 2019), or the protonation of CH2 like at 1P (Altwegg et
al., 1994).

In contrast to CH${}_{4}^{+}$, NH${}_{2}^{+}$ is detected near perihelion,
when the photo-ionisation rate and outgassing rates are larger, and not at
large heliocentric distances. NH${}_{2}^{+}$ results from the dissociative
ionisation of NH3, but its yield is fourfold less than that of NH${}_{3}^{+}$
(Huebner & Mukherjee, 2015). In addition, as NH${}_{2}^{+}$ is lost through
ion-neutral chemistry with H2O, this indicates that its detection at
perihelion stems from a higher production rate from NH3. The peaks in LR at 16
u·q-1 show a right ‘shoulder’ and, at times, a double peak. While this
behaviour may be associated with some instrumental effects (De Keyser et al.,
2015), it more likely results from the contribution of two ion species, for
instance O+ and CH${}_{4}^{+}$ ($\sim 11$ pixels apart in LR) and/or O+ and
NH${}_{2}^{+}$ ($\sim 7$ pixels apart in LR). Overall, in view of the HR
spectra, the main contributor at 16 u·q-1 is O+ whose predominance occurred at
large heliocentric distances prior the inbound Equinox, whereas NH${}_{2}^{+}$
appeared at perihelion when photo-ionisation is much stronger. Interestingly,
CH${}_{4}^{+}$ is more abundant near the outbound Equinox as a possible
consequence of the evolution of the neutral composition: Schuhmann et al.
(2019) showed that there was a clear enhancement in the CH4/H2O ratio by a
factor $\sim 20$ between May 2015 and May 2016.

There exists diverse causes of spacecraft pollution specific to Rosetta
(Schläppi et al., 2010) for nitrogen bearing components which cannot be
completely excluded for observations performed after and close to spacecraft
manoeuvres since UV photolysis of hydrazine N2H4 can be a source of N2H3,
N2H2, and, to a lesser extent, of NH3 and NH2 (Biehl & Stuhl, 1991; Vaghjiani,
1993). As it might in turn affect the detection of NH${}_{4}^{+}$ (through the
ion-neutral reaction NH3+H3O+$\longrightarrow$NH${}_{4}^{+}$+H2O), in
particular after manoeuvres, it might affect NH${}_{3}^{+}$ and NH${}_{2}^{+}$
as well (Beth et al., 2016).

At 17 u·q-1 (Fig. 5), two ions, HO+ and NH${}_{3}^{+}$, have been detected.
Both are mainly produced by ionisation of their respective parent neutral
molecules, H2O and NH3 respectively. HO+ follows the same pattern as the water
production with increased intensity as 67P gets closer to the Sun. The
NH${}_{3}^{+}$ signal is quite strong, mainly near perihelion. In addition to
the ionisation of NH3, NH${}_{3}^{+}$ can be produced through charge transfer
between H2O+ and NH3. Although NH${}_{3}^{+}$ may be lost through the reverse
charge exchange reaction with H2O, the reaction is slow (rates of about 10-10
cm3·s-1) and therefore its contribution to the ion composition remains
negligible compared with others ions reacting with H2O (see details in Section
4.1). To summarise, HO+ is seen throughout the escort phase with a maximum in
intensity near perihelion, when outgassing rate and photoionisation are
strong. NH${}_{3}^{+}$ follows the same pattern with high counts near
perihelion but cannot be detected at large heliocentric distances because its
parent molecule NH3 is much less abundant than H2O. For information, we have
indicated the location of 17O+ (see Fig. 5, bottom): even if it might be
present, its closeness to HO+ ($\sim 7$ pixels apart) and the peak deformation
would prevent its detection. In addition, according to the isotopic ratio
17O/16O derived by Schroeder et al. (2019) and the counts for 16O+, 17O+ is
below the background level.

Peaks in HR do not fall at the exact location of a species and may present a
distorted shape. This phenomenon is symptomatic of spectra at 16, 17, and, to
a lesser extent, of 18 u·q-1 in HR (De Keyser et al., 2015). Without
corrections, the peak is not symmetric and the maximum of the peak is slightly
shifted to the right ($\lesssim 5$ pixels) due to the DFMS’ characteristic
double peak structure for this subset of masses (De Keyser et al., 2015). We
did not apply the correction proposed by De Keyser et al. (2015) as it would
not provide further insight on the ion identification.

l0.315cm+0.6cmb3.4cmLR
l0.315cm+0.6cmb3.4cmHR
l0.315cm+0.6cmb0.33667905cm+3.4cmHR

Figure 6: Same as Fig. 3, but for 18-19 u·q-1. Stacked spectra in low
resolution (top panel) and HR at 18 u·q-1 (middle) and 19 u·q-1 (bottom).
Isotopic ions have been indicated for information. As one may see, the DFMS
mass-per-charge resolution cannot separate: neither H17O+ and DO+ from H2O+,
nor H217O+ and HDO+ from H3O+, and the individual contributions have to be
fitted numerically as done in the neutral mode (cf. Altwegg et al., 2015;
Schroeder et al., 2019). In addition, the peak in LR at 19 u·q-1 is at the
edge of the detector such that it is only partially resolved.

Fig. 6 shows spectra for the range 18-19 u·q-1 and two comments must be made.
First, the range around 18 u·q-1 in both LR and HR was scanned threefold more
often than any other mass ranges as a result of the organisation of the DFMS
measurement sequence: while each u·q-1 range was scanned successively and
increasingly, each sequence started and ended by scanning 18 u·q-1. In
addition, in LR, both spectra centred on 18 and 18.16 u·q-1 have the 19 u·q-1
peak at the edge of the detector with its right part often lost out of the
useful pixel range. HR spectra at 18 u·q-1 show clear signatures of H2O+ and
NH${}_{4}^{+}$ which have been already reported by Fuselier et al. (2016) and
Beth et al. (2016). The weak signal of NH${}_{4}^{+}$ seen at large
heliocentric distances is attributed to hydrazine during manoeuvres (Schläppi
et al., 2010). However, near perihelion, it is produced through ion-neutral
chemistry within the coma due to the high proton affinity of NH3, higher than
that of H2O (Beth et al., 2016). Candidate ion isotopologues are also
indicated, such as 18O+, H17O+, and DO+. The H2O+ peak is so strong and
therefore so widely spread that it prevents their detection as their signal is
expected to be weak if present. Indeed, for the strongest recorded signal, the
peak covers almost 0.30 u·q-1 in total, over $\sim 56$ pixels ($\sim 28$
pixels each side). For comparison, 18O+, H17O+, and DO+ are located at 21, 7,
and 3 pixels on the left (shortwards) of H2O+ respectively. Again, assuming
the isotopic ratios D/H, 17O/16O, and 18O/16O from Schroeder et al. (2019),
18O+, H17O+, and DO+ should be less than 1 count, below the noise level in HR.

H2O+ and NH${}_{4}^{+}$ increase as the heliocentric distance decreases. H2O+
results from direct ionisation of H2O, either by EUV or by electron impact
(Galand et al., 2016). However, it reacts with H2O yielding H3O+ (see Fig. 6,
bottom panel). H2O+ is a peculiar ion as its production and loss depend on the
H2O density in the coma. Photochemical equilibrium, for which production
balances chemical losses, is reached at the location of Rosetta or closer to
the nucleus depending on the outgassing activity (see Section 4.1). Under such
a condition, the H2O+ number density is given by: $\nu/k\sim
10\sqrt{T[\text{K}]}/r_{h}[\text{au}]^{2}$ cm-3 (as dissociative recombination
is a negligible loss) where $\nu$ is the H2O ionisation rate, $r_{h}$ is the
heliocentric distance of 67P, $k$ is the reaction rate constant of:
H2O+H2O+$\longrightarrow$ H3O++HO (see Appendix B), $k$ being relatively
constant, and $T$ is the temperature of the gas. Overall, the H2O+ peak
intensity exhibits this trend, increasing with decreasing heliocentric
distance (and hence increasing $\nu$, at least for the photoionisation,
Heritier et al., 2018), though its variability over a day or during a month,
in particular near perihelion is still puzzling (Beth et al., 2016). Possible
reasons, not investigated in this paper, include the variability of the
generally-negative spacecraft potential (Odelstad et al., 2018), the
interaction between corotating interaction regions and coronal mass ejections
(Hajra et al., 2018) with the 67P’s ionosphere, and to the proximity of the
diamagnetic cavity (Goetz et al., 2016a) or plasma boundaries (Mandt et al.,
2019). NH${}_{4}^{+}$ is produced by proton transfer between NH3 and
protonated molecules, mainly H2O+ and H3O+ (Geiss et al., 1991; Beth et al.,
2016). Amongst neutral cometary species, NH3 has the highest proton affinity
and, therefore, the ability to steal a proton from other protonated molecular
ions (Heritier et al., 2017). NH${}_{4}^{+}$ may be only lost through
transport and, to a lesser extent, dissociative recombination, making it quite
stable together with other ions, such as H2O+. Due to the abundance of its
parent molecule H2O, H2O+ is observed during the whole mission with highest
counts near perihelion, whereas NH${}_{4}^{+}$ is mainly detected near
perihelion because the necessary conditions of its yielding, a large NH3
number density and a large ion-neutral collision frequency, are only met
during this period.

Spectra at 19 u·q-1 (see Fig. 6, bottom panel) exhibit a large and strong peak
associated with H3O+. As H2O has a proton affinity higher than that of HO,
once H2O+ is produced, it readily reacts with H2O to yield H3O+, dominating
the ion composition at short cometocentric distances (Fuselier et al., 2016;
Beth et al., 2016) in the absence of NH3. Nevertheless, when the NH3 number
density is large enough, NH${}_{4}^{+}$ may become the dominant ion species at
distances of a few to tens of kilometres above the nucleus’ surface depending
on the NH3 volume mixing ratio of a few percent. Close to the nucleus, H3O+
dominates, whereas H2O+ becomes the major ion at larger cometocentric
distances. At the location of Rosetta near perihelion ($\sim 150-200$ km),
H3O+ was expected to dominate over NH${}_{4}^{+}$ as is indeed observed. In
Fig. 6 bottom, we have indicated cations containing minor isotopologues,
namely H18O+, H217O+, and HDO+. Like for H2O+, the H3O+ peak is so strong and
wide, covering $\sim 0.30$ u·q-1 ($\sim 52$ pixels slightly lower than at 18
u·q-1) that the small signals corresponding to H18O+, H217O+, and HDO+,
located at 20, 6, and 3 pixels on the left side of H3O+ cannot be
distinguished. Like for NH${}_{4}^{+}$, the H3O+ peak is stronger near
perihelion as expected.

r0.6cmb3.4cmLR
r0.6cmb3.4cm+0.33667905cmHR

Figure 7: Same as Fig 3, but for 20 u·q-1. Stacked spectra in low resolution
(top panel) and in HR at 20 u·q-1 (bottom). The HR spectra at 20 u·q-1 are
already shown in Fig. 6. As one may see, the DFMS mass-per-charge resolution
cannot separate: H317O+ from H2DO+. In addition, the peak in LR at 19 u·q-1 is
at the edge of the detector such that it is only partially resolved.

Fig. 7 shows LR spectra at 19-20 u·q-1 and HR spectra at 20 u·q-1. In HR,
mass-per-charge ratio 19 u·q-1 is at the edge of the detector, on the left
side (shortwards). The peak attributed to H3O+ is sometimes fully resolved
before the pixel shift has been applied, depending on the magnet temperature,
but not afterwards. Indeed, the reference pixel $p_{0}$ has been moved towards
the left by 70 pixels on the 27th of January 2016 such that a peak located at
the pixel $p_{1}$ before the shift was relocated at the pixel $p_{1}-70$
afterwards. Consequently, the lower and upper limits for each spectrum in
terms of mass-per-charge increase by $\sim 1.4\%$ in LR and $\sim 0.23\%$ in
HR after the shift.

A striking difference, when comparing these spectra with LR spectra at lower
u·q-1, is the noise level: it reaches 10 counts, whereas on lower u·q-1, it
never exceeded 3-4 counts. The HR spectra at 20 u·q-1 exhibit two peaks: one
is unambiguously associated with H218O+ and another with either H317O+ or
H2DO+ (or both). The resolving power is not sufficient to separate these
species, only $2\times 10^{-3}$ u·q-1 apart (3-4 pixels). However, we favour
H2DO+ against H317O+: according to Schroeder et al. (2019), in neutral mode,
the H217O signal is too low and buried in the shoulder of the HDO peak and
thus we expect a similar behaviour for the corresponding protonated ions
observed in ion mode. In addition, an estimation might be inferred from the
DFMS-derived isotopic ratios of neutral species (Altwegg et al., 2015;
Schroeder et al., 2019): D/H=$(5.3\pm 0.7)\times 10^{-4}$, 18O/16O=$(2.25\pm
0.18)\times 10^{-3}$, and 17O/16O=$(4.58\pm 0.36)\times 10^{-4}$. Assuming
that the ion-neutral chemistry of H2O+ and H3O+ and their isotopologues relies
on proton transfer and that there is no difference in reaction rate constants,
one may calculate ion abundance ratios. For instance, H2D16O+/H318O+ should be
$\approx$3(D/H)/(18O/16O)=$(0.88\pm 0.12)$ and
H2D16O+/H317O${}^{+}\approx$3(D/H)/(17O/16O)=$(4.3\pm 1.6)$. As the counts
never exceeded $\sim 7$, the peak, which is the sum of both ions, is likely
dominated by H2D16O+ and not H317O+. Moreover, one should note that the peak
20 u·q-1 in LR looks asymmetric for some spectra, with a right shoulder which
might be caused by the simultaneous presence of H218O+ and H2DO+. Overall, a
peak, symmetric or not, is present throughout the mission.

r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmHR

Figure 8: Same as Fig. 3, but for 21-22 u·q-1. Stacked spectra in low
resolution (top panel) and in HR at 20 u·q-1 (bottom). There is no evidence
for a species at 22 u·q-1 in high resolution and the associated spectra are
not shown. In addition, the peak in LR at 23 u·q-1 is at the edge of the
detector such that it is barely caught on its left edge (see also Fig. 2).

Fig. 8 shows spectra in LR and HR for the range 21-23 u·q-1. There is no
ambiguity on the peak at 21 u·q-1, which corresponds to H318O+, detected only
near perihelion. The single chemical path leading to H318O+ is the proton
transfer between H218O and H2O+. The H318O+/21 u·q-1 should follow in
principle the same trend as the H3O+/19 u·q-1 peak. However, it could not be
observed around/after the outbound Equinox (in blue) because of its weakness.

###  3.3 Ion mass-per-charge 22 u·q-1

Fig. 8, top panel, shows a peak at 22 u·q-1 in LR. The associated mass spectra
in HR are not displayed because we do not observe persistent and reliable
signals at 22 u·q-1. The presence of a peak in the LR spectra is at first
surprising since there is no cometary neutral molecule and no candidate for a
protonated one identified at 22 u·q-1. Based on neutral mode observations, we
identify the corresponding ion as CO${}_{2}^{++}$. In neutral mode, the
ionisation of CO2 in the ion source leads to the production of both
CO${}_{2}^{+}$ and CO${}_{2}^{++}$ which is indeed observed in the data,
whereas CO${}_{2}^{++}$ dications detected in the ion mode are naturally
produced in the coma. More details on this finding and a discussion on the
identification of this dication are presented in Section 4.4. It may be
noticed that these cations are only detected before the inbound Equinox at
large heliocentric distances when the H2O density is low which might indicate
that the corresponding ions react with H2O or may only subsist for very low
ion-neutral collision rates. The non detection of CO${}_{2}^{++}$ out of this
period might be also linked to the cometocentric distance of Rosetta and the
CO${}_{2}^{++}$ lifetime. When CO${}_{2}^{++}$ was detected, Rosetta was
around 20-30 km from the nucleus. However, CO${}_{2}^{++}$ is known to have
different lifetimes depending on its electronic states, from $4$ s for its
ground state (Mathur et al., 1995) to $\mu\text{s}$ for excited states (Alagia
et al., 2009; Slattery et al., 2005). Consequently, in order to have
CO${}_{2}^{++}$ be produced in sufficient quantity and detected by DFMS,
$(r-r_{c})/U$, where $r$ is the cometocentric distance of Rosetta, $r_{c}$ is
the nucleus’ radius, and $U$ is the neutral speed, should be of the order of
or lower than the lifetime of CO${}_{2}^{++}$. If $(r-r_{c})/U\gtrsim 4$ s,
only a small fraction of CO${}_{2}^{++}$ ions would have the time to reach the
spacecraft after they have been created. During the periods when DFMS was
operating in ion mode, this condition was the most favourable in December 2014
- January 2015; for this period, $(r-r_{c})/U\approx 20-30$ s (see Fig. 1
assuming $U\sim 1$ km·s-1). This also implies that the CO${}_{2}^{++}$ ions
which are detected are most likely in the ground state (though the latter is
less likely to be produced than the short-lived excited state, Masuoka, 1994;
Alagia et al., 2009).

###  3.4 Ion mass-per-charge range 23 – 40 u·q-1

r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmHR

Figure 9: Same as Fig. 3, but for 23-25 u·q-1. Stacked spectra in low
resolution (top panel) and in HR at 23 u·q-1 (bottom). There is no evidence
for a cation neither at 24 u·q-1 nor at 25 u·q-1 and the associated spectra
are not shown. In addition, the peak in LR at 23 u·q-1 is at the edge of the
detector such that it is barely caught on its right edge (see also Fig. 2).

In this subsection, we present ions which have been detected in this range;
they are not part of the water ion group and are usually minor. Fig. 9 shows
spectra for the range 23-25 u·q-1 in LR and at 23 u·q-1 in HR only. HR spectra
at 24 and 25 u·q-1 do not exhibit any physical and reliable signal and are
thus not displayed. The peak at 23 u·q-1 cannot be fully resolved in LR since,
neither in Fig. 8 nor in Fig. 9, the full extent of the peak is captured by
the detector: in LR for scans centred on 21.98 u·q-1 (Fig. 8), only the left
(lower masses) tail of the peak is captured before the pixel shift, whereas in
LR for scans centred on 24.18 u·q-1 (Fig. 9) only the right (higher masses)
tail of the peak is captured. In both cases, the peak is partially, or not at
all, covered by the detector (see also Fig. 2). However, both indicate a
strong signal near perihelion. This is consistent with the HR spectra and the
detection of Na+ close to perihelion. The associated neutral, Na (sodium), is
a refractory element and has been observed by DFMS early during the escort
phase. Indeed, Wurz et al. (2015) reported its detection in October 2014,
along those of Si (silicium), Ca (calcium), and K (potassium). Their presence
results from solar wind sputtering of dust grains on 67P’s surface. However,
for heliocentric distances below 2 au, the coma of 67P is dense enough to
create the so-called solar wind ion cavity within which the nucleus surface
and dust grains close to Rosetta are shielded from solar wind ions (Behar et
al., 2017). However, the solar wind can still sputter dust grains outside the
cavity and an as-yet-unknown mechanism could therefore transport Na/Na+
inwards, towards Rosetta. Another possibility might be sputtering by ENAs
(Energetic Neutral Atoms), that is, neutralised solar wind ions, which are
expected to be produced in large amount near perihelion, still having access
to the nucleus of 67P unlike the solar wind (Simon Wedlund et al., 2019). Near
perihelion, the source of Na and, therefore, that of Na+ may be different. Na+
was also observed at 21P/Giacobini-Zinner (Geiss et al., 1986; Ogilvie et al.,
1998) and 1P/Halley (Krankowsky et al., 1986). For 21P, Geiss et al. (1986)
ruled out sputtering on cometary grains and favoured the idea that Na was
either trapped in the ice or sorbed in or on carbonaceous grains. Ogilvie et
al. (1998) suggested that Na might come from the evaporation of ice-like
grains, containing Na in ionic form. Another scenario was proposed by Combi et
al. (1997) for 1P/Halley, Benett C/1969 Y1, and Kohoutek C/1973 D for a near
nucleus Na source: the photo-dissociation of the parent molecule, NaOH,
characterised by a high photo-dissociation rate (10-3 s-1 at 1 au, Plane,
1991). We note that Na, as an alkali metal, has a low ionisation energy of
5.139 eV, such that it can be ionised by the intense Lyman-$\alpha$, with a
typical ionisation rate of $7\times 10^{-6}$ s-1 at 1 au (Huebner & Mukherjee,
2015). In addition, it is interesting to point out the peculiar ion-neutral
chemistry of Na and Na+. As an alkali metal, Na is a great electron donor and
reacts with several cations (e.g. H2O+ and NH${}_{3}^{+}$) in the coma through
charge exchange. It also reacts with protonated molecules. Given a high
proton-affinity molecule X (e.g. X = H2O, H2CO, HCN, NH3), and their
protonated version XH+, Na+ results from Na + XH+ $\longrightarrow$ Na+ \+ X +
H. This reaction with protonated molecules, which dominate the ion composition
near perihelion (Heritier et al., 2017), is likely a dominant source of Na+.
Unfortunately, it is impossible to track the evolution of Na+ throughout the
escort phase. In HR, the sensitivity is too low, only allowing the detection
near perihelion. In LR, in spite of the higher sensitivity and the unambiguity
on the ion species at 23 u·q-1, the bad quality of measurements at the edges
of the MCP prevents from obtaining accurate data. However, according to Fig.
8, top panel, one can conclude that the Na+ number density maximises near
perihelion and may reach even higher levels.

Even if peaks are detected in LR at 25 u·q-1, no species have been identified
in HR. As shown in LR, the peaks are lower than 20-30 counts and many spurious
ones contaminate the spectra. From spectra at lower u·q-1, we may define a
rough scaling factor for counts between LR and HR modes. Due to the
differences between slits, the sensitivity in HR is about 10% to 1% of its
level in LR explaining why often no reliable peaks are observed in HR for
C${}_{2}^{+}$ or C2H+. However, the peak at 25 u·q-1 observed in LR at large
heliocentric distances before the inbound Equinox is likely due to C2H+ for
two reasons: there are no other species close to 25 u·q-1 and the peak
disappeared at perihelion, indicating that the corresponding species should
react with H2O, which is the case of C2H+ (Prasad & Huntress, 1980).

r0.6cmb3.4cmLR
r0.6cmb3.4cmHR
r0.6cmb0.33667905cm+3.4cmHR

Figure 10: Same as Fig. 3, but for 26-27 u·q-1. Stacked spectra in low
resolution (top panel), in HR at 26 u·q-1 (middle) and at 27 u·q-1 (bottom).

Fig. 10 shows the range 26-27 u·q-1 in LR (top) and HR (middle and bottom)
also allowing to distinguish the low u·q-1 part of a peak at 28 u·q-1. In HR,
we barely see a peak for C2H${}_{2}^{+}$ at 26 u·q-1, without evidence of CN+.
The analysis for minor species above 23 u·q-1 in HR is made difficult due to
the increased background level requiring unfortunately a visual inspection of
each spectrum with two main criteria to select reliable peaks: (i) a well-
shaped peak (i.e. the peak can be reasonably well-fitted by a single-Gaussian,
two-Gaussian, Lorentzian, or a Voigt profile and be not too spiky/sharp) and
(ii) the existence of a candidate species with an u·q-1 close enough to the
peak location. Not all spectra at 26 u·q-1 exhibit a peak at C2H${}_{2}^{+}$.
Only three spectra, out of more than hundreds through the escort phase, reach
almost 10 counts at the location of C2H${}_{2}^{+}$, which is however
sufficient to ascertain its presence. There is no evidence for CN+ which can
be explained by a reaction rate of CN+ with H2O about tenfold higher than that
of C2H${}_{2}^{+}$ (Anicich, 2003). Near perihelion, the high cometary
activity leads to the loss of CN+ through chemistry impeding its detection.
Stronger signals are observed on LR spectra at large heliocentric distances
but the counts barely exceed 100 which explain the poor or lack of detection
in HR. This is consistent with the detection of a faint peak at the location
of C2H${}_{3}^{+}$ on the spectra at 27 u·q-1 suggesting that the latter has a
common origin with C2H${}_{2}^{+}$. This might be investigated by correlating
both signals which is beyond the scope of this paper. There is no evidence of
HCN+ in HR. However, at large heliocentric distances, LR spectra at 27 u·q-1
exhibit two overlapping peaks of similar amplitude separated by a few pixels
only. Although the low resolution does not allow to unambiguously separate
HCN+ from C2H${}_{3}^{+}$, the double peak at 27 u·q-1 in LR suggests the
presence of both HCN+ and C2H${}_{3}^{+}$ ions with comparable contributions
at large heliocentric distances.

l0.315cm+0.6cmb3.4cmLR
l0.315cm+0.6cmb3.4cmHR
l0.315cm+0.6cmb0.33667905cm+3.4cmHR

Figure 11: Same as Fig. 3, but for 28-30 u·q-1. Stacked spectra in low
resolution (top panel), in HR at 28 u·q-1 (middle) and at 29 u·q-1 (bottom).
There is no evidence for a cation at 30 u·q-1 and the associated spectra are
not shown.

Fig. 11 shows the range 28-30 u·q-1 in LR and HR. Species at 28 u·q-1 are
detected at large heliocentric distances as well as near perihelion. However,
the contributing species are not the same for both periods. As seen in HR, the
detected species at large heliocentric distances are CO+ and, at times,
C2H${}_{4}^{+}$. CO is one of three major components reported by Hässig et al.
(2015) along with CO2 and H2O. However, CO+ reacts with H2O and CO2 and
disappears or is, at least, strongly attenuated close to perihelion (Heritier
et al., 2017). C2H${}_{4}^{+}$ is seen during both periods. However, its
parent molecule may be different during these two periods. At large
heliocentric distances, C2H${}_{4}^{+}$ can be produced by ionisation of C2H4
or dissociative ionisation of C2H6 since the electron-impact ionisation,
dominant at large heliocentric distances, of C2H6 primarily leads to
C2H${}_{4}^{+}$ (Avakyan et al., 1998; Tian & Vidal, 1998b). On the contrary,
at perihelion, it may be produced either by protonation of C2H3 (C2H3+H3O+) or
by charge exchange (C2H4+H2O+). HCNH+ is detected near perihelion as it is
produced through proton transfer, between HCN (or HNC) and H2O+/H3O+ (Heritier
et al., 2017). One spectrum early in the mission exhibits a peak at Si+ but it
may be a spurious and ghost peak since it is observed on a single spectrum and
on one channel only. In fact, both channels are not equally sensitive (due to
ageing, temperature, tuning, etc.) and hence for low signals this may happen
regularly. No conclusion can be drawn, though Si has been detected by another
mass spectrometer of ROSETTA during the same time interval (Wurz et al.,
2015). Si was assumed to be produced by sputtering of the nucleus’ surface by
solar wind ions, still able to access the surface at low outgassing activity
(Behar et al., 2017). There is no evidence of N${}_{2}^{+}$, which is
consistent with Earth-based observations of other comets (Rubin et al., 2015).
The N${}_{2}^{+}$/CO+ ratio is of interest in radio-astronomy and has been
investigated at several comets (Cochran et al., 2000; Cochran, 2002). Due to
the telluric contamination of N${}_{2}^{+}$ or a lack of detection, an upper
limit of this ratio is usually given of the same order as the ratio of their
parent molecules (Rubin et al., 2015), that is less than $\sim 1\%$, which
explains why a possible N${}_{2}^{+}$ peak would be buried in the background.
Overall, although the peak at 28 u·q-1 is present at any time of the escort
phase, the main contributors may have changed during Rosetta’s escort phase:
CO+ (and maybe C2H${}_{4}^{+}$) at large heliocentric distances,
HCNH++C2H${}_{4}^{+}$ near perihelion.

At 29 u·q-1, two ion species, HCO+ and C2H${}_{5}^{+}$, are clearly visible on
HR spectra near perihelion. Spectra in LR at 29 u·q-1 show a peak in January
2016 (green) with the same amplitude as near perihelion (orange) while, in HR
mode, both species are detected mainly near perihelion. The presence of HCO+
is puzzling as it should be lost through chemistry with H2O. Included in the
photochemical model of Heritier et al. (2017) for 67P at perihelion, HCO+ is
produced through ion-neutral chemistry and its contribution is of the order of
CO+ which is however not observed at that time by DFMS. At 1P/Halley, Haider &
Bhardwaj (2005) had used the same ion-neutral chemical reactions plus another
one: C++H2O. The latter was calculated to contribute up to 10% to the total
amount of HCO+ in 1P’s coma. At 67P, DFMS did not perform detections below 13
u·q-1 such that C+ at 12 u·q-1 cannot be qualitatively assessed. However, we
do not expect C+ to be significantly dense enough near perihelion to yield
HCO+. Rosetta was close to the nucleus between 150 and 200 km and the
potential sources of C+ are limited to carbon-bearing molecules, namely CO2,
CO, and, to a much lesser extent, H2CO. That said, the potential direct
sources of HCO+ are the dissociative ionisation of H2CO (present in the coma,
Heritier et al., 2017), the photo-dissociation of H2CO into HCO followed by
its ionisation, and/or ion-neutral chemistry of H2CO with ions. HCO+ barely
reacts with H2O, like CH${}_{3}^{+}$ (Herbst, 1985): perihelion is hence a
favourable period for its production with a stronger EUV flux and the limited
loss through ion-neutral chemistry. However, it is difficult to assess the
contribution of the different chemical pathways, as the most likely parent
molecule, H2CO, may be a distributed source like at 1P/Halley (Meier et al.,
1993). Due to the poor spatial coverage of Rosetta near perihelion, it is
difficult to determine the neutral number density profile of H2CO and
determine whether or not it departs from a $\sim 1/r^{2}$ dependency. However,
studying distributed sources for neutrals is beyond the scope of this paper
dedicated to ions. No HR spectra at 30 u·q-1 are shown as no peak was
detected, which is consistent with the relatively weak signal in LR ($\lesssim
60$ counts). The LR spectra do not exhibit significant differences between
observations at perihelion and at large heliocentric distances, meaning that
the main contributor may change throughout the mission and/or may not react
with H2O.

l0.315cm+0.6cmb3.4cmLR
l0.315cm+0.6cmb3.4cmHR
l0.315cm+0.6cmb3.4cmHR
l0.315cm+0.6cmb0.33667905cm+3.4cmHR

Figure 12: Same as Fig. 3, but for 31-33 u·q-1. Stacked spectra in low
resolution (top panel) and in HR at 31 u·q-1 (second panel), at 32 u·q-1
(third panel), and at 33 u·q-1 (fourth panel).

Fig. 12 shows spectra for the range 31-33 u·q-1. Peaks at 31, 32, and 33 u·q-1
are similar, showing larger intensities at heliocentric distances from 2 to
2.5 au after perihelion (green), in particular at 33 u·q-1, slightly above the
levels observed near perihelion (orange). At 31 u·q-1, H2COH+ (protonated
formaldehyde) is clearly identified in HR as the major ion species. H2CO has a
proton affinity higher than that of H2O, such that H2COH+ is mainly produced
through H2CO+H2O+ or H2CO+H3O+. There are some weak (¡5-7 counts) peaks at the
location of phosphorus cation P+ near perihelion as seemingly shown by the
observations of phosphorus atoms and amino-acids in neutral mode reported by
Altwegg et al. (2016). The faint level of the P+ signal in HR can be explained
by the loss of P+ which reacts, even slowly, with the most abundant neutral
species such as H2O, CO2, and NH3. For HR spectra at 32 u·q-1, there are
several peaks with low intensity (¡20 counts per 19.8 s) which can be
identified by superimposing a number of spectra. S+ and CH3OH+ are clearly
detected. Amongst the sulphur-bearing molecules detected in the coma of 67P by
DFMS, the most abundant species near perihelion are H2S, of which dissociative
ionisation leads to S+ in part, and neutral S (Calmonte et al., 2016). There
is no significant loss of S+ through chemistry as it does not react with the
dominant neutral species H2O, CO2, and CO. Regarding CH3OH+, the methanol
cation, CH3OH, the parent molecule of the methanol cation CH3OH+, is quite
abundant around perihelion, between 0.5% and 3% with respect to H2O, based on
MIRO (Microwave Instrument for the Rosetta Orbiter, Gulkis et al., 2007) sub-
millimetre radio-telescope observations (Biver et al., 2019). Nevertheless,
these values, derived from measurements of the column density between the
surface of the nucleus and Rosetta, differ from those (0.1-0.3%) measured in
situ by DFMS for the same period (see Fig. 10 from Heritier et al., 2017).
This difference may be attributed to the adiabatic expansion of the gas along
the line of sight (Heritier et al., 2017) and the differences between the bulk
speed of light and heavy species. In addition, there are also disagreements
between H2O local measurements by DFMS and integrated column measurements by
MIRO and VIRTIS with associated impacts on the relative abundances of cometary
species (Hansen et al., 2016; Marshall et al., 2017; Combi et al., 2020).
There is one occurrence for a peak in one channel at the location
O${}_{2}^{+}$, of which the main sources are charge exchange (O2+H2O+) and
ionisation of O2. The latter was not considered in Heritier et al. (2017) such
that O${}_{2}^{+}$ is underestimated in their model. While the conditions were
favourable for its production, its detection cannot be confirmed. Indeed,
another peak of similar intensity is also observed between O${}_{2}^{+}$ and
CH3OH+ locations but cannot be assigned to a given species, it is most likely
a ghost peak, that is an unphysical peak.

Concerned by spacecraft contamination, in Fig. 12 (third panel) we have added
N2H${}_{4}^{+}$, of which the ionisation potential is 8.1 eV (Meot-Ner et al.,
1984). No strong signal is detected at this mass, ruling out its contribution
to 32 u·q-1. Same precaution has been undertaken for 33 u·q-1 with
N2H${}_{5}^{+}$ because N2H4 has a proton affinity just above that of NH3. As
only one peak is observed (see Fig. 12, bottom panel), this is not conclusive.
HR spectra at 33 u·q-1 reveal the presence of a protonated molecule, namely
CH3OH${}_{2}^{+}$ (protonated methanol), produced from CH3OH+H2O+ and
CH3OH+H3O+. CH3OH${}_{2}^{+}$ is more abundant than the ionised methanol and
is the main contributor to mass 33 u·q-1. A faint accumulation of signals is
observed around HS+. Considering the main sulphur-bearing molecules, H2S and S
(see discussion for mass 32 u·q-1), the possible processes to generate HS+ are
rather limited. The most likely source is the photo-dissociative ionisation of
H2S, though this process is not very efficient (Huebner & Mukherjee, 2015).
Regarding its loss, S has a low proton affinity with respect to water and HS+
easily loses its proton for the benefit of H2O, for instance.

r0.6cmb0.33667905cm+3.4cmLR

Figure 13: Same as Fig. 3, but for 34-37 u·q-1. There is no evidence of
cations in high resolution and the associated spectra are not shown.

Fig. 13 shows stacked LR spectra for the range 34-37 u·q-1. None of the
associated HR spectra present clear peaks over this range. The most curious
and surprising absence is that of H3S+, even in LR. H2S is present in the coma
of 67P (Calmonte et al., 2016) and its proton affinity is higher than that of
H2O, though lower than those of H2CO, HCN, CH3OH, and NH3 (Heritier et al.,
2017). It is thus predicted to be detectable by Rosetta near perihelion, a
period favourable for proton transfer (Heritier et al., 2017). However, the
sensitivity of DFMS in the ion mode is affected by (i) the instrument energy
acceptance window and, to a lesser extent, (ii) the decrease of the detector
efficiency at higher energies. As the mass-per-charge ratio u·q-1 increases,
the DFMS energy acceptance window decreases. Indeed, to select ions with
respect to their mass-per-charge, a specific post-acceleration
$V_{\text{acc}}\propto(\text{u\textperiodcentered q}^{-1}$)-1 is applied
within DFMS. The energy acceptance window for the electrostatic analyzer is
$\sim 20\text{ V}\pm 0.1\%|V_{\text{acc}}|$ such that ions are filtered
through a narrower energy range when the mass-per-charge ratio increases.
Further details are given in Schläppi (2011). In views of previous works at
1P/Halley (e.g. Eberhardt et al., 1994), the main contributors are 34S+,
followed by H2S+ and 13CH3OH${}_{2}^{+}$ at 34 u·q-1, and H3S+ at 35 u·q-1.
However, for the latter, it is clear that the signal is not stronger in LR
near perihelion, which is unexpected for a protonated molecule.

r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmHR

Figure 14: Same as Fig. 3, but for 37-40 u·q-1. Stacked spectra in low
resolution (top panel) and in low resolution at 39 u·q-1 (bottom). There is no
evidence of cations in HR at 37 u·q-1 and at 38 u·q-1 and the associated
spectra are not shown.

Fig. 14 shows LR spectra (top panel) for the range 37-40 u·q-1 and HR spectra
(bottom panel) for 39 u·q-1. No cations have been detected except at 39 u·q-1.
Because the signal-to-noise ratio in HR spectra increases with higher masses,
we have only kept spectra with signals above 4 counts and present on both
channels in order to get rid of contamination by spurious and unreliable
signals. Only two species are expected: C3H${}_{3}^{+}$ and K+. Korth et al.
(1989) argued that the peak at 39 u·q-1 at 1P/Halley, detected by PICCA (Korth
et al., 1987) on board Giotto, was C3H${}_{3}^{+}$ and ruled out K+. However,
thrice, on the 8th, the 9th, and 11th of August 2015, K+, and not
C3H${}_{3}^{+}$, was detected. K (potassium) is an alkali metal like Na yet
with a lower electronegativity. Ion-neutral reaction rates of K+ and K of
interest for astrochemistry are practically inexistent; however, as K belongs
to the same group as Na, we can likely assume that K+ will undergo the same
interaction with neutrals as Na+. K was detected early in the mission at large
heliocentric distances by Wurz et al. (2015), along with Na. In spite of the
rather faint K+ and Na+ peaks in the HR mode, we believe that the presence of
these ions in the ionised coma is undoubted, supported, in particular, by the
detection of K and Na neutral atoms by ROSINA instruments. The photo-
ionisation of K and Na appears as a likely production mechanism. However, the
ultimate origin of these neutral species is still debated as already mentioned
for Na.

###  3.5 Ion mass-per-charge range $>40$ u·q-1

r0.6cmb0.33667905cm+3.4cmLR

Figure 15: Same as Fig. 3, but for 41-44 u·q-1. Stacked spectra in low
resolution only.

We have decided not to look at u·q-1 above 40, as the sensitivity of DFMS in
HR becomes too low to unambiguously allow a clear identification of the
detected species . This is illustrated by a comparison of Fig. 15 with Fig. 8.
The most contributing species at 44 u·q-1 is undoubtedly CO${}_{2}^{+}$ and
consequently CO${}_{2}^{++}$ at 22 u·q-1 at least at large heliocentric
distances, especially over the Southern Hemisphere pre-perihelion and over
both hemispheres post-perihelion (although no ion DFMS spectra are available,
see Fig. 1) where the neutral coma is dominated by CO2 (Hässig et al., 2015;
Gasc et al., 2017). As explained in Section 4.4, the production rate of the
CO${}_{2}^{++}$ dication is a hundredfold lower, at least, than that of the
monocation CO${}_{2}^{+}$ but the relative intensity between the peaks at 22
u·q-1 and 44 u·q-1 is only about 10% which means that the DFMS sensitivity at
44 u·q-1 is about tenfold lower than that at 22 u·q-1, a number consistent
with mass-dependency reported by Schläppi (2011). In addition, CO${}_{2}^{+}$
is mainly lost through charge exchange with H2O such that near perihelion, it
is close to photochemical equilibrium (i.e. its loss is through chemistry, not
transport, and its number density barely varies with respect to the
cometocentric distance above tens of kilometres, see Fig. 7 and 8 in Heritier
et al., 2017).

##  4 Discussion

###  4.1 Highlights

HR ion mode dataset of ROSINA-DFMS has allowed us to directly identify for the
first time different ion species and obtain an improved knowledge of the
composition of a cometary plasma (columns 4 and 5 in Table 1). Our analysis
has revealed the complexity of the cometary ionosphere made of ionised neutral
atoms and molecules and also of ionised radicals and protonated molecules (see
Section 4.2). We also confirm the presence of isotopologues (see Section 4.3).
Observations from previous cometary missions, such as Giotto, had to rely on
photochemical models to determine the exact nature of numerous ions detected
at a given u·q-1. From ground-based observations, only a few ions were
directly identified. Table 1 shows in the second and third columns a
compilation of the previous knowledge on cometary plasma composition from the
literature (Delsemme, 1985, 1991; Balsiger et al., 1995; Lis et al., 1997;
Huebner et al., 1991; Haider & Bhardwaj, 2005) from 13 u·q-1 (lower bound for
DFMS) to 40 u·q-1.

We have detected a number of new ion species that have not been predicted
before the Rosetta mission such as the alkali metal ions Na+ and K+ and the
dication CO${}_{2}^{++}$ (see Section 4.4). Some ion species predicted by
photo-chemical models have not been detected by DFMS. Their presence cannot be
ruled out since favourable conditions for their detection may have not been
met, such as a lower outgassing rate of the nucleus compared to 1P during the
Giotto fly-by, lack of DFMS ion mode closer to the nucleus, and, probably of
great importance, the detrimental influence of the spacecraft potential on the
effective acceptance of DFMS in ion mode.

| Previous works (see caption) |  | This work
---|---|---|---
u·q-1 | Identified species | Predicted species |  | Detected species in HR |  Peaks in LR but not in HR
|  | (from modelling) |  |  |  (new candidate)
13 | CH+ |  13C+, CH+ |  | CH+ |
14 |  | N+, 13CH+, CH${}_{2}^{+}$ |  | CH${}_{2}^{+}$ |
15 |  | NH+, 13CH${}_{2}^{+}$, CH${}_{3}^{+}$ |  | CH${}_{3}^{+}$ |
16 |  | O+, NH${}_{2}^{+}$, 13CH${}_{3}^{+}$, CH${}_{4}^{+}$ |  | O+, NH${}_{2}^{+}$, CH${}_{4}^{+}$ |
17 | HO+ | HO+, NH${}_{3}^{+}$, 13CH${}_{4}^{+}$, CH${}_{5}^{+}$ |  | HO+, NH${}_{3}^{+}$ |
18 | H2O+ | H2O+, NH${}_{4}^{+}$ |  | H2O+, NH${}_{4}^{+}$ |
19 | H3O+ | H3O+ |  | H3O+ |
20 |  | H218O+, H317O+, H2DO+ |  | H218O+, H2DO+ |
21 |  | H318O+ |  | H318O+ |
22 |  |  |  |  |  ✓(CO${}_{2}^{++}$)
23 |  |  |  | Na+ |
24 |  | C${}_{2}^{+}$ |  |  |  ✓
25 |  | C2H+ |  |  |  ✓
26 | CN+ | CN+, C2H${}_{2}^{+}$ |  | C2H${}_{2}^{+}$ |
27 |  | HCN+, C2H${}_{3}^{+}$ |  | C2H${}_{3}^{+}$ |
28 | CO+, N${}_{2}^{+}$ | CO+, N${}_{2}^{+}$, HCNH+, C2H${}_{4}^{+}$ |  | Si+*, CO+, HCNH+, C2H${}_{4}^{+}$ |
29 | HCO+ | HCO+, N2H+, C2H${}_{5}^{+}$ |  | HCO+, C2H${}_{5}^{+}$ |
30 |  | NO+, H2CO+, CH2NH${}_{2}^{+}$, C2H${}_{6}^{+}$ |  |  |  ✓
31 |  | HNO+, H2COH+ |  | P+*, H2COH+ |
32 |  | S+, O${}_{2}^{+}$, CH3OH+ |  | S+, O${}_{2}^{+}$*, CH3OH+ |
33 |  | HS+, HO${}_{2}^{+}$, CH3OH${}_{2}^{+}$ |  | CH3OH${}_{2}^{+}$ |
34 | H2S+ |  34S+, H2S+, 13CH3OH${}_{2}^{+}$ |  |  |  ✓
35 |  | H3S+ |  |  |  ✓
36 |  | H234S+, H333S+, C${}_{3}^{+}$ |  |  |  ✗
37 |  | C3H+, H3O·H2O+ |  |  |  ✓
38 |  | C2N+, C3H${}_{2}^{+}$ |  |  |  ✓
39 |  | C3H${}_{3}^{+}$ |  | K+ |
40 | Ca+ | CH${}_{2}^{+}$CN+, C3H${}_{4}^{+}$ |  |  |  ✓
Table 1: Compilation of the ions predicted and detected at comets as a
function of the mass. ‘Identified species’ are those detected by UV, IR,
visible or radio spectroscopic ground-based observations of comets (Delsemme,
1985, 1991; Lis et al., 1997, see Section 1). ‘Predicted species’ are those
included in photochemical models for 1P/Halley (Huebner et al., 1991; Haider &
Bhardwaj, 2005) or considered as dominant for the corresponding mass (Balsiger
et al., 1995). ‘Detected species in HR’ refers to those detected with ROSINA-
DFMS for this study. ‘Peaks in LR but not in HR’ refers to peaks detected in
LR (✓), (✗) otherwise, while no peaks were present in HR with a strong
candidate in parenthesis (see Section 4.4). * means for Si+, P+, and
O${}_{2}^{+}$ that, although a peak is located at the correct mass, it is seen
once or twice in one of the two channels only, which may cast some doubt on
their presence.

Most, yet not all, cometary ion species may be sorted into three main families
(hereinafter F):

  * (F1)

those produced by ionisation of a parent molecule $p$ and lost through
transport, like CH${}_{3}^{+}$ (i.e. reacting slowly or not at all with H2O),

  * (F2)

those produced by ionisation of a parent molecule $p$ and lost through
chemistry with H2O, the dominant neutral species, like HO+, CH${}_{4}^{+}$
(i.e. the ion species $X^{+}$ is reacting with H2O such that
$X^{+}+\text{H}_{2}\text{O}\longrightarrow$ products)

  * (F3)

those produced through ion-neutral chemistry only and lost either by transport
(e.g. NH${}_{4}^{+}$), by chemistry, or both (e.g. molecules with a proton
affinity between those of H2O and NH3, such as H3O+, CH3OH${}_{2}^{+}$, and
H3S+). Ions produced from high proton affinity neutrals are further discussed
in Section 4.2.

However, not all ion species belong to one of these families.

For ions produced by ionisation, there is a possibility to quantify which loss
process, that is transport (F1) or chemical reactions with water (F2),
dominates. Beth et al. (2019) showed that a dimensionless parameter of
interest is:

$\alpha_{X^{+}}=k_{X^{+}+\text{H}_{2}\text{O}}\dfrac{n_{\text{H}_{2}\text{O}}(r)r^{2}}{Ur_{c}}=k_{X^{+}+\text{H}_{2}\text{O}}\dfrac{Q_{\text{H}_{2}\text{O}}}{4\pi
U^{2}r_{c}},$ (4)

where $k_{X^{+}+\text{H}_{2}\text{O}}$ stands for the reaction rate constant
of $X^{+}+\text{H}_{2}\text{O}\longrightarrow$ products (see Appendix B),
$n_{\text{H}_{2}\text{O}}(r)$ is the local water number density close to the
comet (assumed $\propto 1/r^{2}$), $Q_{\text{H}_{2}\text{O}}$ is the
outgassing rate of H2O, $U$ is the ion outward radial speed, assumed to be
that of neutrals and constant with cometocentric distance, and $r_{c}$ is the
radius of the nucleus. $\alpha_{X^{+}}$ gauges which loss process dominates.
For $\alpha_{X^{+}}\gg 1$, $X^{+}$ is mainly lost through chemistry with water
close to the surface, while for $\alpha_{X^{+}}\ll 1$, it is through
transport. At large cometocentric distances, the loss through transport always
dominates such that the ion number decreases in $1/r$ asymptotically. Due to
the range of kinetic rates depending on the species as showed in Table 3 and
the evolution of the water number density, $\alpha_{X^{+}}$ is depending as
well as the heliocentric distances. In order to assess the evolution and the
variability of some detected ions, one should refer to Appendix B for detailed
information on the photo-ionisation and kinetic rates.

Electron-ion dissociative recombination is negligible at the location of
Rosetta (Heritier et al., 2018; Beth et al., 2019); as a result, the ion
number density profile of these ions is given by (adapted from Eq. B.3 in Beth
et al., 2019):

$\displaystyle n_{X^{+}}(r,\alpha_{X^{+}},\tau_{c})=$
$\displaystyle\dfrac{\nu_{p}Q_{p}}{4\pi
U^{2}}\left(\dfrac{r_{c}}{r}\right)^{2}\exp\left[\alpha_{X^{+}}\dfrac{r_{c}}{r}\right]$
$\displaystyle\times\left\\{\dfrac{r}{r_{c}}E_{2}\left[(\alpha_{X^{+}}+\tau_{c})\dfrac{r_{c}}{r}\right]-E_{2}\left[\alpha_{X^{+}}+\tau_{c}\right]\right\\}$
$\displaystyle\phantom{n_{X^{+}}(r,\alpha_{X^{+}},\tau_{c})}=$ $\displaystyle
n_{X^{+}}(r,\alpha_{X^{+}}+\tau_{c},0)\exp\left[-\tau_{c}\dfrac{r_{c}}{r}\right],$

where $E_{2}$ stands for the exponential integral function, $\tau_{c}$ the
mean optical depth at the surface for a water-dominated coma, significant near
perihelion (Heritier et al., 2017; Beth et al., 2019), $\nu_{p}$ the
ionisation frequency of the neutral parent molecule $p$ yielding $X^{+}$, and
$Q_{p}$ the outgassing rate of the parent molecule. For low outgassing
activity ($\alpha_{X^{+}}\ll 1$, i.e. $Q\lesssim 10^{25}$ s-1, and
$\tau_{c}\ll 1$), the main loss for the ion is through transport regardless
the cometocentric distance and its number density profile converges towards
(Galand et al., 2016):

$n_{X^{+}}(r)\approx\dfrac{\nu_{p}Q_{p}}{4\pi U^{2}}\dfrac{r-r_{c}}{r^{2}},$
(5)

For high outgassing activity ($\alpha_{X^{+}}\gg 1$), the loss is dominated by
chemistry with H2O for distances close to the nucleus; as a result, the ion is
in photochemical equilibrium and its density is given by:

$n_{X^{+}}(r)\approx\dfrac{\nu_{p}}{k_{X^{+}+\text{H}_{2}\text{O}}}\dfrac{Q_{p}}{Q_{\text{H}_{2}\text{O}}}\left(\dfrac{\alpha_{X^{+}}}{\alpha_{X^{+}}+\tau_{c}}\right)\exp\left[-\tau_{c}\dfrac{r_{c}}{r}\right],$
(6)

valid up to tens or hundreds of kilometres above the surface, depending how
high $\alpha_{X^{+}}$ is. Interestingly, $\alpha_{X^{+}}/4$ corresponds to the
ratio between the maximum $X^{+}$ number density reached with transport-
dominated loss (i.e. Eq. 5 at $r=2r_{c}$) and that reached with chemistry-
dominated loss, when photo-absorption is ignored (i.e. Eq. 6 with
$\tau_{c}=0$).

Fig. 16 illustrates the effect of increasing the relative importance of
reactions with H2O on the ion number density. As $Q_{\text{H}_{2}\text{O}}$
increases, $\alpha_{X^{+}}$ increases and the ion density profile is damped
and flattened to the photo-chemical value, or lower in presence of photo-
absorption (based on Eq. 6). At large cometocentric distances, the ion number
density profile follows Eq. 5. As 67P got closer to the Sun, $\alpha_{X^{+}}$
increased from $\ll 1$ to $\gg 1$, similarly to $\tau_{c}$. Beth et al. (2019)
assessed the importance of the photo-absorption near perihelion. According to
the average photo-absorption cross-section of H2O that they derived, the
optical depth is $\sim 2--3$ near perihelion at the nucleus’ surface. This
entails a decrease in the effective H2O photo-ionisation rate by 7–20 at the
comet’s surface.

Figure 16: Dimensionless ion number density profile versus scaled
cometocentric distance as a function of $\alpha$ with $\tau_{c}=0$ (left
panel) and $\tau_{c}=3$ (right panel). On the $y$-axis, the values of the ion
number density $n_{i}$ are scaled with respect to the maximum ion number
density reached when only transport is considered for a given ion, i.e.
$\alpha_{X^{+}}=0$. As $\alpha_{X^{+}}$ (coloured curves) increases, more and
more ion-neutral reactions take place especially close to the nucleus. This
results in damping the number density of ions reacting with H2O.

Fig. 17 shows, for different cometary outgassing conditions, the number
density profiles for ions with a number density lower than H3O+ between the
nucleus and the location of Rosetta: H2O+, CO${}_{2}^{+}$, CH${}_{n}^{+}$
($n=1-4$), and NH${}_{m}^{+}$ ($m=0-3$). These ions are produced by
(dissociative) ionisation of parent molecules (H2O, CO2, CH4, and NH3) and
lost through transport or chemistry with mainly H2O. H2O+ may also be produced
by ion-neutral chemistry (charge transfer between cations and H2O; e.g.
H2O+CO${}_{2}^{+}$) but this process has been neglected here as it is
significantly less efficient than photo-ionisation. As H2O is the main EUV
absorber in the coma near perihelion, we have used the same optical depth
(i.e. 3) in Fig. 17 for all photo-ionisation rates. Results shown in Fig. 17
for three outgassing conditions provide a simple insight about the ions of
main interest and their number density profiles, even if obtained using
several simplifying assumptions. Firstly, only photo-ionisation has been
considered, whereas Galand et al. (2016) and Heritier et al. (2018) have shown
that electron-impact ionisation is usually the main ion source, against photo-
ionisation, at large heliocentric distances. Solar-wind charge exchange has
also been neglected although it may contribute to the ionisation of neutrals
(Simon Wedlund et al., 2019) in some cases at large heliocentric distances.
Increasing the ionisation rate would shift profiles to higher number
densities. Then, the ion dynamics and any significant acceleration of the
cometary plasma has been neglected. Galand et al. (2016) and Heritier et al.
(2018) have shown that considering the ion velocity to be that of neutral
species is a good approximation to assess the plasma density at large
heliocentric distances. Near perihelion, if the ion speed was significantly
higher than that of neutrals as suggested by Odelstad et al. (2018),
$\alpha_{X^{+}}$ would be lower, preventing ion-neutral chemistry to happen
and chemically-produced species, such as NH${}_{4}^{+}$ and CH3OH${}_{2}^{+}$,
from being present. Results shown in Fig. 17 are thus more reliable when one
compares various ions originating from the same parent molecule such as, for
example, CH${}_{3}^{+}$ and CH${}_{4}^{+}$. At large cometocentric distances
(see Fig. 17, top panel), transport is always dominating such that
CH${}_{4}^{+}$ number density is slightly higher than that of CH${}_{3}^{+}$
(assuming both only produced from CH4). At the location of Rosetta however,
because CH${}_{4}^{+}$ does react with H2O (F2) and CH${}_{3}^{+}$ does not
(F1), CH${}_{3}^{+}$ is similar to (top panel,
$n_{\text{CH}_{3}^{+}}/n_{\text{CH}_{4}^{+}}\approx 1$) or dominates over
CH${}_{4}^{+}$ (middle and bottom panels,
$n_{\text{CH}_{3}^{+}}/n_{\text{CH}_{4}^{+}}>10$).

Figure 17: Ion density profiles of a few ions as a function of the
cometocentric distance under different conditions: December 2014 ($Q=7.5\times
10^{25}$ s-1, 2.75 au, top panel), end of August 2015 ($Q=5\times 10^{28}$
s-1, 1.24 au, middle panel), end of January 2016 ($Q=2.4\times 10^{27}$ s-1,
2.2 au, bottom panel). We have used the same neutral composition for each
case: 88.7% H2O, 10% CO2, 1% NH3, 0.3% CH4. The photo-ionisation frequencies
$\nu_{p}$ are taken from Huebner & Mukherjee (2015) at low solar activity at 1
au and scaled with respect to the heliocentric distance. The expanding speed
of the gas is set to a constant $U=900$ m·s-1 and the nucleus’ radius $r_{c}$
to 2 km. Ions are produced from photo-ionisation or photo-dissociative
ionisation of the neutral molecules (see Section B). The gas temperature is
assumed constant at $T=100$ K. The grey areas correspond to the cometocentric
distance of Rosetta during these periods. The uncertainties from the kinetic
rates on the ion number density are represented through coloured shades.

Figures 16 and 17 may help defining the most favourable conditions for
detecting these ions. At low outgassing activity, when transport dominates the
ion loss, the ion number density is $\propto\nu_{p}Q_{p}$, both parameters
increase as the comet gets closer to the Sun. The ion number density profile
peaks at $r=2r_{c}$, meaning that the best location for detection is close to
the nucleus. Under high outgassing activity (when photochemical equilibrium is
achieved), the ion number density is approximately $\propto\nu_{p}$ , i.e.
only the photo-ionisation rate and its variations with the heliocentric
distances will control the ion number density. However, the location at which
the ion density profile peaks also depends on the absorption of solar EUV
radiation by neutral species in the coma, essentially H2O. If photo-absorption
is neglected, the ion density peaks closer to the nucleus’ surface and has a
plateau that may extend to distances ranging between a few tens to thousands
of kilometres depending on the cometary activity (see Fig. 16, left panel).
Under optically thick conditions, the EUV radiation cannot penetrate deep
enough into the coma to ionise neutrals such that the ion number density peaks
farther away from the nucleus’ surface (see Fig. 16 and compare profiles of
similar colours).

Depending on the ion species and whether or not it reacts with H2O, the peak
of the ion number density is not located at the same cometocentric distance:

  * (F1)

for ions which do not react with H2O nor with any other major neutral species
and are mainly produced by ionisation of parent molecules (e.g.
CH${}_{3}^{+}$), the maximum number density is reached around $r\approx
2r_{c}$ and the ion density decreases asymptotically in $1/r$,

  * (F2)

for ions which do react with H2O (e.g. HO+ and CH${}_{4}^{+}$), their number
density peaks at different cometocentric distances, depending on the cometary
activity (see Fig. 17). For low activity, the ion number density peaks at
$r\approx 2r_{c}$. For high activity, these species start to reach
photochemical equilibrium such that their number density profile exhibits a
plateau from the nucleus’ surface to tens or hundreds of cometary radii, and
then decreases asymptotically in $1/r$,

  * (F3)

for high proton-affinity molecular ions (e.g. NH${}_{4}^{+}$, H3O+, and
CH3OH${}_{2}^{+}$), their number density peaks close to the comet’s surface,
but decreases asymptotically faster, in $\log(r)/r^{2}$ (Beth et al., 2019).

The aforementioned statements are true when photoabsorption is negligible. For
outgassing rates higher than $10^{28}$-$10^{29}$ s-1, photoabsorption matters
(Beth et al., 2019): the ion number densities peak farther away from the
comet’s surface as EUV solar radiation cannot penetrate deep enough into the
coma to ionise neutrals (see Fig. 16, right panel).

###  4.2 High proton affinity

The detection of protonated high-proton-affinity molecules other than H3O+ at
67P, such as NH${}_{4}^{+}$, confirms the expectations from the models.
Because of the insufficient resolution of the ion spectrometers on board
Giotto, NH${}_{4}^{+}$ was blended with H2O+ such that its contribution to the
peak at 18 u·q-1 has only been assessed from a photochemical modelling based
on the measured neutral composition. Prior to Rosetta’s arrival at 67P, Vigren
& Galand (2013) attempted to assess the total contribution of these water ions
inside the diamagnetic cavity. Though this cavity, detected around perihelion,
was not as extended as expected (Goetz et al., 2016b), that does not prevent
these ions from being produced close to the nucleus, then transported
outwards, outside the diamagnetic cavity, and finally detected by the
instrument. Indeed, during the escort phase and near perihelion, Beth et al.
(2016) have unambiguously detected NH${}_{4}^{+}$ in the coma of 67P, well
separated from H2O+ thanks to the high resolution of the ROSINA-DFMS HR mode.
Fig. 6 clearly shows that near perihelion, at the location of Rosetta, H2O+
was on the average more abundant compared with NH${}_{4}^{+}$. In the present
paper, we have confirmed the presence of three additional protonated molecules
in the coma: HCNH+, H2COH+, and CH3OH${}_{2}^{+}$,that were previously
predicted from a photo-chemical modelling (Heritier et al., 2017). Their
detection attests the importance of ion-neutral chemistry and collisions in a
high activity environment.

Except H3S+ (undetected in HR), HCNH+, H2COH+, and CH3OH${}_{2}^{+}$ are
strong candidates for the peaks observed at 28, 31, and 33 u·q-1 respectively
in the LR mode. By means of a photochemical modelling, Heritier et al. (2017)
investigated the relative contribution at 28 u·q-1. They found that CO+ should
be lower than HCNH+. However, their modelling did not include C2H${}_{4}^{+}$,
another candidate at 28 u·q-1 in LR, and its associated chemistry, although
this species is clearly detected in HR simultaneously with HCNH+ on the same
spectrum (see Fig. 11, middle). The non-detection of H3S+ in HR is still not
clearly understood. It has been pointed out that it may be related to the
energy acceptance of the instrument which decreases at higher masses (Heritier
et al., 2017). It may also come from the origin and the source of H2S. In the
photo-chemical model, the source and background of the neutral species are
supposed to exclusively come from the gas released following the ice
sublimation at the surface, excluding extended sources. However, Calmonte et
al. (2016) have suggested from DFMS neutral measurements that part of H2S is
associated with dust grains.

Finally, CH${}_{3}^{+}$ might be considered to be part of (F3) as CH2 has a
higher proton affinity than HO and H2O (Altwegg et al., 1994). However, Fig. 4
shows no evidence for higher counts of CH${}_{3}^{+}$ at perihelion when
proton transfer reactions are favoured, compared to large heliocentric
distances.

###  4.3 Water isotopologues

Figs. 7 and 8 attest the presence of water ion isotopologues, namely H218O+,
H2DO+, and H318O+, near perihelion. As a consequence, DO+ and H18O+ must also
be present but the mass resolution of ROSINA-DFMS is not high enough to
separate them from H2O+ and H3O+ signals, respectively (see Fig. 6). In view
of the isotopic ratios D/H ($\sim 5.3\times 10^{-4}$, Altwegg et al., 2015)
and 18O/16O ($\sim 1.8\times 10^{-3}$, Schroeder et al., 2019) derived from
the neutral mode of DFMS, the detection and the count rates of H218O+, H2DO+,
and H318O+ ($\lesssim 10$) are consistent with those of H2O+ and H3O+
(103-104) derived in HR near perihelion. In principle, D/H and 18O/16O
isotopic ratios might be determined from the LR ion observations (Balsiger et
al., 1995) using count rates at 19 and 20 u·q-1. However, as already noticed
in Section 3.2, the quality of DFMS measurements at 19 u·q-1 close to the edge
of the detector is not good enough to provide the necessary accuracy for the
data processing. Moreover, ions at 19, 20, and 21 u·q-1 were not probed
exactly at the exact same time (i.e. approximately 30 seconds between
individual spectra) and plasma conditions (and thus the spacecraft potential)
have been shown to vary on short time scales.

On the 16th of September 2015 from 08:59 UT, two successive HR spectra exhibit
reliable peaks in both channels: 1) at 20 u·q-1 associated with H318O+ and the
doublet H317O+ and H2DO+ (not separable) and 2) at 21 u·q-1 with a peak
associated with H318O+. Peaks have been fitted with one Gaussian or double-
Gaussian for the highly abundant ion H3O+ only while the total counts
corresponding to the peaks of its isotopologues with very small amplitudes
were simply obtained by summing the individual pixel counts. In this occasion,
the isotopic ratio may be derived although with a limited accuracy. Firstly,
(3D/H+17O/16O) may be inferred from (H2DO++H317O+)/H3O+ (Eberhardt et al.,
1995). We obtained $\sim 10^{-3}$ which is compared with (2.05$\pm$0.3)$\times
10^{-3}$ from neutral isotopic ratio obtained by DFMS observations in the
neutral mode (Altwegg et al., 2015; Schroeder et al., 2019). Secondly,
(18O/16O) may be inferred from H318O+/H3O+ (Eberhardt et al., 1995). We
obtained $\sim 10^{-3}$ which is compared with (2.25$\pm$0.18)$\times 10^{-3}$
from the DFMS neutral mode (Schroeder et al., 2019).

As aforementioned, the isotopic measurements have only been possible on very
few occasions when the count rates in the ion HR mode were sufficient.
Moreover, the accuracy of their derivation is limited due to several factors
such as the intrinsic smaller effective sensitivity of DFMS and the variable
instrument energy acceptance controlled by the large and varying spacecraft
potential. It is therefore out of scope to perform accurate isotopic
measurements using HR ion data. However, our results show a relatively good
agreement between both the ion-derived and neutral-derived isotopic
abundances.

###  4.4 Dications

As shown in Section 3.3, ROSINA-DFMS provides the first unambiguous detection
of the doubly-charged (or dication) CO${}_{2}^{++}$ in a cometary ionosphere.
Dications have been previously observed in dense planetary ionospheres
throughout the Solar System: O++ at Earth (Hoffman, 1967), Venus (Taylor et
al., 1980), perhaps at Mars (Dubinin et al., 2008) and at Io, along with S++
(Frank et al., 1996) and C++ (Sandel et al., 1979). The recent (and only)
review about ionospheric doubly-charged ions was published by Thissen et al.
(2011). Additional information on organic dications and their structures,
aiming at presenting their chemical properties may be found in Lammertsma et
al. (1989). Thissen et al. (2011) reviewed stable dications stemming from the
most abundant atoms and molecules found in planetary atmospheres, namely: C++,
N++, O++, CO++, N${}_{2}^{++}$, NO++, O${}_{2}^{++}$, Ar++, CO${}_{2}^{++}$.
Notwithstanding a couple of ions indistinguishable by mass spectrometry
(N${}_{2}^{++}$ and N+, O${}_{2}^{+}$ and O+), many of these dications have
u·q-1 very close to monocations from different neutral atoms, molecules, and
radicals (e.g. N++ and Li+, CO++ and N+, NO++ and CH${}_{3}^{+}$, Ar++ and
Ne+) and their detection is difficult or impossible even with high performance
mass spectrometers such as ROSINA-DFMS. Moreover, there were only a few
attempts in assessing their potential role within an ionosphere and exosphere.
For instance, Lilensten et al. (2013) showed that the presence of
CO${}_{2}^{++}$ and its associated chemistry in the upper atmosphere may play
a non-negligible role in the ion escape. Furthermore, Falcinelli et al. (2016)
showed that the Coulomb explosion of CO${}_{2}^{++}$ produces CO+ and O+ at a
few eVs, which is energetic enough to overcome the gravitational attraction of
the planet.

ROSINA-DFMS has two advantages for the detection of these peculiar and scarce
ions compared with the ion spectrometer on Giotto: a higher sensitivity and a
wider mass coverage. By covering several integers in terms of u·q-1, DFMS may
probe those stemming from double ionisation of odd-mass-number parent
molecules.

It is likely that even-mass-number dications are also present but buried into
monocations’ signals (e.g. O${}_{2}^{++}$ with O+ or N${}_{2}^{++}$ with N+),
but, fortunately, this is not the case for CO${}_{2}^{++}$ at 22 u·q-1. In
neutral mode, the neutrals are ionised by electron impact with energies of 45
eV. The recorded signal at 22 u·q-1 stems from CO2 whose double ionisation
threshold is $\sim 37$ eV. In the ion mode, ions have not been ionised within
the instrument and originate from the cometary plasma. In order to firmly
establish the nature of ions detected at 22 u·q-1, we have considered the
following possibilities. We first check for candidate monocations: 22Ne+,
D218O+, HD217O+, and H2D18O+ (Balsiger et al., 1995). Ne (Neon) has never been
detected (Rubin et al., 2018) and, regarding the isotopic composition (Altwegg
et al., 2015; Schroeder et al., 2019), hydronium ion isotopologues have a too
low abundance to be detected. In a second step, we check for candidate
dications and consider other neutral candidates than CO2 at 44 u: CS, CH2ON,
C2H4O, C2H6N, and C3H8 (Altwegg et al., 2017). Within this list, two arguments
are in favour of the CO${}_{2}^{++}$ dication: CO2 is dominant at 44 u and the
peak at 22 u·q-1 correlates with latitude, as CO2 does. Fig. 18 shows LR
spectra at 22 u·q-1 above both hemispheres. To exclude seasonal variations, we
selected the time period from October 2014 to February 2015 (pre-equinox).
ROSINA-DFMS probed frequent and strong signals at 22 u·q-1 above the Southern
Hemisphere, where CO2 is abundant (Hässig et al., 2015). Over the pre-equinox
Northern Hemisphere where CO2 is minor, no peak is present at 22 u·q-1.
ROSINA-DFMS cannot simultaneously probe both ion and neutral compositions such
that a ‘direct’ correlation cannot be assessed. However, we find that, pre-
equinox, both the neutral CO2 abundance and the signal at 22 u·q-1 in ion mode
correlate with the spacecraft latitude. Finally, to strengthen our
argumentation, one should compare the pre-equinox count rates at 22 u·q-1 and
at 44 u·q-1. At 22 u·q-1, the count rates were almost 100 while at 44 u·q-1,
they were almost 1000. However, as aforementioned, the instrument sensitivity
is a tenfold higher at 22 u·q-1 than at 44 u·q-1 such that the 44 u·q-1/22
u·q-1 ratio is $\approx 100$, consistent with the ratio of photo-ionisation
cross-sections of CO2 leading to CO${}_{2}^{+}$ and CO${}_{2}^{++}$ (Masuoka,
1994; Tian & Vidal, 1998b).

r0.6cmb3.4cmLR
r0.6cmb3.4cmLR
r0.6cmb0.33667905cm+3.4cmLR

Figure 18: Spectra in LR at 22 u·q-1 over three latitudinal regions prior to
the 1st of March 2015: above +30∘ (upper panel, $\sim 170$ stacked individual
spectra), between -30∘ and +30∘ (middle panel, $\sim 470$ stacked individual
spectra), and below -30∘ (bottom panel, $\sim 290$ stacked individual
spectra). Due to a uneven latitudinal coverage when DFMS was operating in ion
mode, more spectra were acquired above the Southern Hemisphere which might be
a source of observational bias.

Unfortunately, no detection of dications except at 22 u·q-1 could be achieved
from DFMS data, even if the design of ROSINA covers half-odd-integer masses.
Indeed, in neutral mode, when neutrals are ionised in the gas chamber and then
travel within the instrument, dications are observed at 13.5 u·q-1 for
example. There are some hints to explain the lack of detection of stable
doubly-charged ions other than CO${}_{2}^{++}$. Dications are primarily
produced from direct ionisation of neutrals. As ROSINA-DFMS detection starts
at 13 u·q-1, only dications associated with neutral species above 26 u could
be detected. Considering in first instance CO, H2CO, CH3OH, H2S, and
hydrocarbons (saturated or not), all of them have an even-integer mass number
and the corresponding u·q-1 of the doubly-charged ion falls on a similar value
of existing singly-charged ions except at 22 u·q-1 (see Table 1). The peak at
22 u·q-1 never exceeded 100 counts. Let’s assume that double-ionisation of
neutral species follows a similar pattern than that of CO2, i.e. it occurs
hundredfold less than a single ionisation (Masuoka, 1994; Tian & Vidal,
1998a). A neutral parent species with an odd-integer mass number (which offers
no overlap of the dication peak with other species) should exceed $\sim 1$% of
the CO2 volume mixing ratio in order for its daughter dication to generate 1
count, and even more to have the peak above the noise level. No cometary
neutral species fulfil these two requirements.

The detection of stable dications also raises questions about their effects
upon plasma dynamics and behaviour. Unlike terrestrial ionospheres, gravity is
not at play at comets such that any neutral or ionised species will leave
independently of their mass. The dynamics of an ion will only depend on its
mass-per-charge ratio, as the dynamics is usually dominated by electromagnetic
forces. In the case of CO${}_{2}^{++}$, its mass-per-charge is not so
different from the water-group ions and, therefore, its effect is marginal on
the dynamics. However, it might play a role in the ion-chemistry and in the
production of ‘energetic’ CO+ and O+, of about a few eVs. The term energetic
may mislead the reader though. Within a terrestrial ionosphere like those of
Mars, Earth, or Titan, the ions are (almost) thermalised with the ambient
neutral species such that their temperature is equal to the neutral
temperature which never exceeds 1000 to 2000 K at most. In that case, an ion
of 1 eV or more is classified as energetic compared with the neutral species.
Within a cometary ionosphere, ions already have a significant amount of
energy. Considering ions travelling between 1 (that of neutral species) and 8
km·s-1, as observed by and at Rosetta (Odelstad et al., 2018), the ion kinetic
energy is between 0.09 and 6.0 eV for H2O+, between 0.15 and 9.3 eV for CO+,
and between 0.083 and 5.3 eV for O+. In comparison, the kinetic energy
released by the Coulomb explosion of CO${}_{2}^{++}$ into CO+ and O+ lies
within these values (Falcinelli et al., 2016) such that this process is not a
significant additional source of energy. Aside this energetic aspect, the
double photoionisation of CO2 might be a marginal source of CO+ and O+ ions.
As suggested in Section 3.3, DFMS likely detected the ‘stable’ CO${}_{2}^{++}$
which is produced in lower quantities than its metastable/unstable
counterpart, quickly dissociating into CO++O+ (Masuoka, 1994; Slattery et al.,
2005; Falcinelli et al., 2016).

Lastly, our analysis shows that not only the low outgassing rate
($Q<10^{25}-10^{26}$ s-1) but also the proximity of Rosetta to 67P (i.e. a few
to tens of kilometres) played a key role in the detection of CO${}_{2}^{++}$
because of the low lifetime of CO${}_{2}^{++}$ ($\leq 4$ s, Mathur et al.,
1995). Hence, we anticipate its detection is possible at a comet only if these
conditions are met.

###  4.5 Future missions

Given the instrumental, operational, and orbital constraints, we showed that
DFMS had a remarkable capability to assess the ion composition of the cometary
plasma, though this was not its primary goal. Unlike Giotto, Rosetta was a
non-spinning and quasi-stationary spacecraft with respect to 67P. DFMS was
often operating in ion mode when it was pointing towards the comet, though
manoeuvres were ongoing. These different characteristics may greatly influence
the ion detection. In the case of a flyby for a future mission, the spacecraft
will fly at speeds of tens of km·s-1 with respect to the target. There are two
advantages of such a trajectory. Firstly, it gives a radial coverage of the
plasma number density and composition over thousands of kilometres over a
short period of time (typically a few hours) limiting time-dependent effects
(e.g. changes in outgassing and neutral composition). Secondly, most of the
cometary ions will be collimated in a limited region of velocity phase space
in the spacecraft frame of reference (since the ion thermal speed is below or
of the order of their mean speed) such that the instrument will capture the
bulk of the cometary ions with a limited field of view. In addition, having a
fast spacecraft minimises the troublesome effect of the spacecraft potential,
very negative most of the time and varying for Rosetta. Indeed, considering a
cometary ion, its energy, mainly being kinetic, will be
$\tfrac{1}{2}m_{i}(\@vec{v}_{i}-\@vec{v}_{\text{SC}})^{2}\approx\tfrac{1}{2}m_{i}\@vec{v}^{2}_{\text{SC}}$
with respect to the spacecraft. Typically, for H2O+, its kinetic energy in the
spacecraft frame spans from 9 eV (${v}_{\text{SC}}=10$ km·s-1) to 330 eV
(${v}_{\text{SC}}=60$ km·s-1) to be compared with the largest negative
spacecraft potential of Rosetta, around -30 V.

In the case of an escorting spacecraft like Rosetta, being closer (a few tens
of kilometres instead of hundreds of kilometres, or performing the dayside
excursion) to the nucleus near perihelion would have improved the signal to
noise and the detection of ions resulting from chemistry and dissociative
ionisation of neutrals. However, being too close (typically $\sim 10$ km) may
not be the best location for very active comets and is not safe for the
spacecraft. Indeed, the maximum in the plasma number density is not located at
the same cometocentric distance depending on the outgassing activity (see
Section 4.1).

Although 67P has a low-to-intermediate outgassing activity, many cations have
been detected even at $150-200$ km from the nucleus near perihelion. This
result indicates that a very active comet is not a requirement regarding the
current instrumental capabilities for mass spectrometers like DFMS. For weakly
active comets (¡1027-1028 s-1), an escorting spacecraft, like Rosetta, is the
best option as it allows to measure the ion composition over a long period, at
different stages of its outgassing activity. However, additional aspects
should be considered for future similar missions: limiting the manoeuvres
during the spectrometer’s scans, allocating operational time for the ion mode
close to the nucleus, and a more uniform time coverage throughout the mission.
As shown in Fig. 1, DFMS ion dataset is relatively sparse, excluding safe mode
and excursions. Regular ion scans would have helped to track the evolution of
the ion composition through the mission. Running alternatively LR and HR modes
should be also considered. For very active comets (¿1029 s-1), a flyby is the
best option from an ion composition perspective, though a different
instrumentation is required such as a time-of-flight (TOF) mass spectrometer.
A fly-by requires a high time resolution and spectrometers like DFMS do not
suit (one DFMS spectrum at a given u·q-1 is acquired during 20 s every 10-15
min). TOF spectrometers have a much higher time resolution acquiring several
u·q-1 all at once, at the expense of a lower sensitivity. Beside fostering the
ion-neutral chemistry, very active comets exhibit several boundaries (Mandt et
al., 2019) and regions including the diamagnetic cavity (Cravens, 1989). Ion
composition may differ inside and outside such a cavity (Balsiger et al.,
1986). Although it was detected at 67P (Goetz et al., 2016a, b), diamagnetic
crossings were on average less than 30 minutes, which corresponds to less than
3 scans for a specific mass-per-charge ratio by DFMS, limiting its ability to
probe the composition inside the cavity. More time should be spent inside this
region and that can be best achieved at very active comets. In addition, one
flyby will allow to assess the composition inside and outside over a short
time period.

##  5 Conclusion

The mass spectrometer DFMS with its HR mode outperforms any in situ
measurements made at comets so far. In particular, its HR ion mode has been
extremely valuable for identifying ions present within a coma at close range.
Although the time coverage in ion mode is more restricted than that of the
neutral mode by far, DFMS has produced invaluable results, which are presented
in this paper. Amongst all of them, a very new and interesting result is the
first detection of the CO${}_{2}^{++}$ dication. For future studies, to make
the most out of the ion ROSINA-DFMS dataset, cross-analysis should be
performed with the set of instruments from the RPC Consortium.

###### Acknowledgements.

The authors very warmly thank Cyril Simon Wedlund for having provided
invaluable feedback on the manuscript. The dataset analysed during the current
study together with a user manual for data analysis are available in the ESA-
PSA archive (ftp://psa.esac.esa.int/pub/mirror/INTERNATIONAL-ROSETTA-
MISSION/ROSINA/) and on request to the authors. Rosetta is an ESA mission with
contributions from its Member States and the U.S. National Aeronautics and
Space Administration (NASA). Work at Imperial College London is supported by
STFC of UK under grant ST/N000692/1 and ESA under contract
No.4000119035/16/ES/JD. AB and MG thank the support from Imperial College
London through the European Partners Fund. AB thanks also the support from the
Swedish National Space Agency (grant 108/18). JDK was supported by the Belgian
Science Policy Office via PRODEX/ROSINA PEA 90020. We would like to
acknowledge ISSI for the great opportunity it offered us for very valuable
discussions on this topic as part of the International Team “Plasma
Environment at 67P After Rosetta”.

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##  Appendix A Correction for 13, 14, and 15 u·q-1

Additional corrections are required for the mass calibration at 13, 14, and 15
u·q-1. First, as mentioned in Section 2.2, the zoom factor $z$ is slightly
lower than at higher mass-per-charge ratios, at 5.5. Moreover, by applying the
procedure described in Section 2.2, Fig. 19 shows spectra with two distinct
peaks, each of them associated with a different period: warm-coloured before
the 27th of January 2016 and cold-coloured afterwards. On the 27th of January
2016, the location of the impinging beam was shifted such that the ions would
hit the detector on less-deteriorated parts (see Fig. A.6 in Schroeder et al.
2019).

r0.6cmb3.4cmHR
r0.6cmb3.4cmHR
r0.6cmb0.33667905cm+3.4cmHR

Figure 19: Original spectra when applying the procedure described in Section