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# C + 133.5 nm Emission Mechanisms on Mars Revealed by the MAVEN Observations

Shotaro Sakai, Hiromu Nakagawa, Justin Deighan, Sonal Jain, Kei Masunaga, Fuminori Tsuchiya, Naoki Terada, Majd Mayyasi, Nicholas Schneider, David Mitchell, et al.

► To cite this version:

Shotaro Sakai, Hiromu Nakagawa, Justin Deighan, Sonal Jain, Kei Masunaga, et al. C + 133.5 nm Emission Mechanisms on Mars Revealed by the MAVEN Observations. The Astrophysical Journal, 2024, 977 (2), pp.226. 10.3847/1538-4357/ad8e35. hal-04929703

HAL Id: hal-04929703

https://hal.science/hal-04929703v1

Submitted on 6 Feb 2025

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# C⁺ 133.5 nm Emission Mechanisms on Mars Revealed by the MAVEN Observations

Shotaro Sakai<fnref n="1" />,<fnref n="2" /> ID, Hiromu Nakagawa<fnref n="1" /> ID, Justin Deighan<fnref n="3" /> ID, Sonal K. Jain<fnref n="3" /> ID, Kei Masunaga<fnref n="4" /> ID, Fuminori Tsuchiya<fnref n="2" /> ID, Naoki Terada<fnref n="1" /> ID, Majd Mayyasi<fnref n="5" /> ID, Nicholas M. Schneider<fnref n="3" /> ID, David L. Mitchell<fnref n="6" /> ID, Christian Mazelle<fnref n="7" /> ID, Mehdi Benna<fnref n="8" />,<fnref n="9" />, Robert J. Lillis<fnref n="6" /> ID, Go Murakami<fnref n="10" /> ID, Shannon M. Curry<fnref n="3" /> ID, and Kanako Seki<fnref n="11" /> ID

## Abstract

C⁺ emission is generated by electron impact, dissociative ionization, photoionization, and resonant scattering with carbon-related atoms, molecules, and ions in the Martian ionosphere and thermosphere. The contribution of each mechanism to the emission, however, has not been elucidated due to the difficulty of observation and the fact that a part of the emission cross section is unclear. The current paper isolates the C⁺ emission mechanism using remote-sensing and in situ observations on board Mars Atmosphere and Volatile EvolutioN. Both electron impact and dissociative ionization/photoionization contribute to C⁺ emission below 150 km altitude when the CO density is high, but only dissociative ionization/photoionization contributes to the emission for the low CO density case, while only dissociative ionization/photoionization dominates the emission at altitudes between 150 and 165 km for both CO density cases. It is difficult to estimate the total flux of suprathermal electrons in the ionosphere from remote-sensing observations of C⁺ emission because the contribution of electron impact to C⁺ emission is small. In contrast, C-atom remote-sensing observations might provide a better understanding of the total flux of suprathermal electrons in the ionosphere than C⁺ emission, and global ultraviolet observations could be utilized as a tool for monitoring the ionosphere. The total flux of suprathermal electrons estimated from C-atom emission may be utilized to isolate the contribution of each C⁺ emission process to the brightness more accurately. This suggests that the C⁺ and C-atom emissions might be tracers of spatiotemporal variations in the Martian ionosphere and thermosphere.

*Unified Astronomy Thesaurus concepts: Mars (1007); Ultraviolet spectroscopy (2284); Extreme ultraviolet astronomy (2170); Solar system terrestrial planets (797)*

## 1. Introduction

It is believed that the Martian climate was at least episodically warm and humid more than four billion years ago, based on evidence from observations and models, because Mars maintained a thick atmosphere and liquid water was stable on its surface (e.g., J. M. Goldspiel & S. W. Squyres 1991; B. M. Jakosky et al. 1994; M. C. Malin & K. S. Edgett 2003; R. Wordsworth 2016; E. S. Kite 2019). Present-day Mars, on the other hand, has only a thin atmosphere composed mainly of CO₂, and there is no liquid water on the surface. Mars has thus experienced significant escape to space and associated atmospheric evolution over ~4.5 billion years. Hydrogen and oxygen escape associated with water loss have been investigated based on previous observations and simulation studies. Hydrogen escape is primarily caused by Jeans escape (D. M. Hunten & M. B. McElroy 1970; M. B. McElroy & T. M. Donahue 1972) of hydrogen atoms generated by the photodissociation of H₂O (e.g., Chassefière & Leblanc 2004), although nonthermal escape via dissociative recombination of

HCO⁺ can be also significant (B. S. Gregory et al. 2023). Recent observations have found that the hydrogen corona and escape flux have large seasonal variations of more than 1 order of magnitude, peaking around the summer solstice in the southern hemisphere (M. S. Chaffin et al. 2014; J. T. Clarke et al. 2014; J. S. Halekas 2017). Oxygen escape is largely caused by nonthermal processes, including dissociative recombination, sputtering, ion outflow, and ion pickup (e.g., Chassefière & Leblanc 2004; D. A. Brain et al. 2017; R. J. Lillis et al. 2017). Previous studies have reported that oxygen escape by photochemical escape, ion outflow, and ion pickup increases under high X-ray and extreme-ultraviolet (XUV) conditions (B. M. Jakosky et al. 2015; T. E. Cravens et al. 2017; R. Ramstad et al. 2018), and oxygen escape by ion outflow and ion pickup decreases under high solar wind dynamic pressure conditions (R. Ramstad et al. 2018). However, while the escape of hydrogen and oxygen has been investigated, there are still few studies on the escape of carbon atoms and carbon ions originating from carbon dioxide (CO₂) and carbon monoxide (CO), which are major constituents of the Martian atmosphere.

Carbon escape is one of the key processes in understanding Martian atmospheric evolution. Carbon may escape into space not only as molecular ions (S. Inui et al. 2018, 2019) but also as

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Original content from this work may be used under the terms of the [Creative Commons Attribution 4.0 licence](https://creativecommons.org/licenses/by/4.0/). Any further distribution of this work must maintain attribution to the author(s) and the title of the work, journal citation and DOI.

<footnote n="1">¹Department of Geophysics, Graduate School of Science, Tohoku University, Sendai, Japan; shotaro@tohoku.ac.jp</footnote>

<footnote n="2">²Planetary Plasma and Atmospheric Research Center, Graduate School of Science, Tohoku University, Sendai, Japan</footnote>

<footnote n="3">³Laboratory for Atmospheric and Space Physics, University of Colorado Boulder, Boulder, CO, USA</footnote>

<footnote n="4">⁴Institute of Arts and Sciences, Yamagata University, Yamagata, Japan</footnote>

<footnote n="5">⁵Center for Space Physics, Boston University, Boston, MA, USA</footnote>

<footnote n="6">⁶Space Sciences Laboratory, University of California, Berkeley, Berkeley, CA, USA</footnote>

<footnote n="7">⁷IRAP, CNRS, University of Toulouse, UPS-OMP, CNES, Toulouse, France</footnote>

<footnote n="8">⁸Center for Space Sciences and Technology, University of Maryland, Baltimore, MD, USA</footnote>

<footnote n="9">⁹Solar System Exploration Division, NASA Goddard Space Flight Center, Greenbelt, MD, USA</footnote>

<footnote n="10">¹⁰Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, Sagamihara, Japan</footnote>

<footnote n="11">¹¹Department of Earth and Planetary Science, Graduate School of Science, University of Tokyo, Tokyo, Japan</footnote>

<footnote>Received 2024 June 21; revised 2024 October 23; accepted 2024 November 01; published 2024 December 16</footnote>
carbon (C) atoms via photodissociation of CO₂ and CO and electron impact on them. In the upper atmosphere, in particular, it escapes as ions via photoionization and dissociative ionization of CO₂ and CO. Understanding the global distribution of C atoms and C⁺ ions could therefore provide insights into the escape of carbon, including CO₂ and CO, to space, an important factor in Martian atmospheric evolution. This is made possible by ultraviolet (UV) observations that include the emission wavelengths of C atoms (156.1 nm and 165.7 nm) and C⁺ (133.5 nm). D. Y. Lo et al. (2022) first characterized the C-atom dayglow in the Martian atmosphere through a combination of photochemical and radiative transfer models and analyses of the Imaging Ultraviolet Spectrograph (IUVS) instrument on board Mars Atmosphere and Volatile Evolution (MAVEN). They found that the C-atom or C I emission is driven by the electron impact of CO₂ at lower altitudes and the resonant scattering at higher altitudes. The C atoms produced by photodissociation and electron impact result in photochemical escape into space (D. Y. Lo et al. 2021, 2022). CO₂ and CO, on the other hand, can also escape into space in the form of C⁺ after dissociative ionization or photoionization (N. B. Pickett et al. 2022). C⁺ was not previously observed in the Martian upper atmosphere because it is difficult to observe due to the closeness of its mass to that of major ions such as O⁺. N. B. Pickett et al. (2022) is the first paper to present the C⁺ observations from MAVEN SupraThermal And Thermal Ion Composition, showing C⁺ fluxes on the order of 10⁵ cm⁻² s⁻¹ in the tail region. However, the number of C⁺ observations is small, and only local values of C⁺ flux have thus been obtained on the nightside, meaning that its global distribution is not well understood.

One method to obtain the C⁺ ion global distribution is the detection of the UV emission described above, referred to as C⁺ or C II emission, at a wavelength of 133.5 nm. C⁺ emission can be driven by the dissociative ionization of CO₂ and CO (e.g., T. E. Cravens & A. E. S. Green 1978), the photoionization of C, the electron impact of CO₂, CO, and C (e.g., T. E. Cravens & A. E. S. Green 1978; S. V. Avakyan et al. 1998), and the resonant scattering of C⁺ (e.g., J. C. Raymond et al. 1997). However, the details of emission mechanisms have not been well understood because the emission cross section is unknown in many cases including dissociative ionization, although the C-atom emission mechanism is somewhat understood from models (e.g., D. Y. Lo et al. 2022). C⁺ emission was observed on the nightside in the aurora observed by the Emirates Mars Ultraviolet Spectrometer instrument (R. J. Lillis et al. 2022), likely from energetic electrons. An accurate understanding of the C⁺ emission mechanism allows us to take a closer look at the C⁺ production mechanisms and the altitudinal distribution of CO₂ and CO. If each contribution of these emission mechanisms can be isolated from UV observations, it might be possible to identify the C⁺ origin and understand whether the time variation of the C⁺ distribution is of internal (caused by, e.g., dust storms and atmospheric waves) or external (e.g., solar extreme-ultraviolet (EUV) radiation and solar wind conditions) origin. Furthermore, temporal variations in the suprathedral electron distribution in the ionosphere might be revealed by understanding both C⁺ and C-atom emission processes accurately, since electron impacts associated with suprathedral electron fluxes (e.g., R. W. Schunk & A. F. Nagy 2009) can be important for both emission processes.

This study, for the first time, isolates the C⁺ emission mechanisms in the ionosphere and thermosphere at each altitude range below 300 km, and examines whether the temporal and spatial variation of the suprathedral electron distribution can be investigated from the combination of C⁺ and C-atom emission processes. Section 2 briefly describes the method and instruments used in this study, and Section 3 presents the results of remote-sensing and in situ observations. Section 4 provides a discussion that includes a comparison of remote-sensing and in situ observation results, and Section 5 summarizes the main results of the paper and discusses future work.

## 2. Method and Instruments

### 2.1. C⁺ Emission Mechanisms

C⁺ emission is excited by dissociative ionization (DI), photoionization with excitation (PI), electron impact ionization and/or excitation (EI), and resonant scattering (RS). In the current study, DI means a dissociation process in which molecules are ionized and excited by solar radiation, while PI means a process by which atoms are ionized and excited by solar radiation. EI denotes the process by which electron impact dissociates molecules with ionization and excitation, ionizes and excites atoms, and excites ions, and RS denotes the process of absorption and reemission of solar photons of specific wavelength by atoms and ions. We can estimate the contribution to C⁺ emission by the RS and EI of CO₂, CO, and C from spacecraft observations. CO₂⁺ and CO⁺ might also contribute to C⁺ emission, but at altitudes below 200 km, they are much smaller than the CO₂ and CO densities (M. Benna et al. 2015; P. R. Mahaffy et al. 2015) and are ignored in the current study.

The brightness of the C⁺ emission caused by EI and RS can be estimated using the EI cross section and the g-factor associated with RS, whereas it is difficult to estimate the brightness for emission by DI/PI because the cross sections are not known. The current study first investigates the dominant emission mechanism and then estimates the contribution of DI/PI to C⁺ emission by a comparison between remote-sensing and in situ observations.

#### 2.1.1. Emission by Electron Impact

The brightness of C⁺ emission by EI of CO₂ and CO is estimated from the electron differential flux, the emission cross sections, and the thermospheric CO₂ and CO densities (e.g., R. W. Schunk & A. F. Nagy 2009), using the following equation with neutral species, s, with energy Eₑ:

$$B_s = \int_{x_0}^{x_1} \left( n_s \int_{E_c}^{\infty} \phi_e \sigma_s^{\text{EI}} dE \right) 10^{-6} dx \quad (1)$$

where B_s, n_s, φₑ, and σ_s EI denote brightness [R], thermospheric density [cm⁻³], electron differential flux [cm⁻² s⁻¹ eV⁻¹], and the emission cross section by electron impact ionization [cm²], respectively; x₀ and x₁ are the interval of the path integral. We can simplify Equation (1) as the following equation by assuming the path integral as a function of the scale height:

$$B_s \sim n_s L_s 10^{-6} \int_{E_c}^{\infty} \phi_e \sigma_s^{\text{EI}} dE \quad (2)$$
where $L_s = 2\sqrt{H_s}r$ is the length scale [cm] with a scale height, $H_s$ [cm], and $r$ [cm] is the planetary radius. C⁺ emission by EI of C atoms can be estimated similarly, but the emission cross section for e-C collision is not available and the C atomic density is small in the thermosphere (D. Y. Lo et al. 2022), and thus the emission by e-C collision is ignored in the current study. It is also discussed in Section 4.

C⁺ emission by EI of C⁺ is estimated using the CHIANTI database (version 10.1; K. P. Dere et al. 1997; G. Del Zanna et al. 2021). CHIANTI provides UV and X-ray spectral lines emitted from atomic ions excited by electron collisions when the suprathermal electron density and temperature are given, and it can figure out the brightness of their lines.

### 2.1.2. Emission by RS

C⁺ emission is also driven by the RS of C⁺ ions. The brightness is estimated from the following equation when assuming the path integral as a function of the scale height as well as the emission by EI:

$$B_i \sim n_i L_i 10^{-6} \frac{g}{r_{\text{Sun}}^2} \quad (3)$$

where $n_i$ is the C⁺ density [cm⁻³], $L_i = 2\sqrt{H_i}r$ is defined as the length scale of C⁺ [cm], g is the g-factor [s⁻¹], $r_{\text{Sun}} = 1.524$ is the distance from the Sun [au], and $H_i$ is the scale height of C⁺. The g-factor is defined as the scattering rate, or the number of photons scattered per second per ion moving with velocity relative to the Sun and with thermal velocity (R. R. Meier 1990), and it is $2.1 \times 10^{-5}$ s⁻¹ for C⁺ 133.5 nm emission at the Martian orbit, which is calculated using Equation (6) in R. R. Meier (1990). This value is calculated using the EUV radiation on 2017 October 20 used in this study, which is mentioned later.

### 2.1.3. Emission by DI and PI with Excitation

The brightness of the C⁺ emission by DI and PI of CO₂, CO, and C with excitation might be important at high altitudes above 200 km, where the atmosphere is optically thin. The emission is approximated in the form

$$B_s \sim n_s L_s 10^{-6} \int_{0}^{\lambda_{\text{si}}} I_{\infty} \exp(-\tau) \sigma_s^{\text{DI/PI}} d\lambda \quad (4)$$

where $I_\infty$ is the photon flux at the top of the atmosphere, $\lambda_{\text{si}}$ is the ionization threshold wavelength for neutral species, s, τ is the optical depth, and $\sigma_s^{\text{DI/PI}}$ is the emission cross section by DI or PI with excitation. These cross sections are not available, however, making it difficult to estimate the contribution of C⁺ emission to DI and PI with excitation.

## 2.2. MAVEN Instruments and Observations

The brightness of the C⁺ emission by EI and RS is estimated using the MAVEN observations. Note that the estimation of emission brightness by DI and PI is difficult for the reasons indicated above. The estimation of emission brightness by EI and RS needs the electron differential flux and thermospheric and ionospheric density.

### 2.2.1. Instruments

The IUVS instrument (W. E. McClintock et al. 2015) measures middle-ultraviolet (MUV) and far-ultraviolet (FUV) emission from the Mars atmosphere. We compare the emission brightness estimated from the MAVEN in situ observations with the brightness measurement from IUVS. IUVS uses a single spherical mirror telescope to image the Martian atmosphere onto the 0°.06 × 11° incident slit of a plane-grating spectrograph. It has two independent fields of regard, one designated as “limb” (24° × 11°) and the other designated as “nadir” (60° × 11°), to allow for both altitude profiles and spatial maps. IUVS can measure the emission line for the atoms and the bands for the molecules using an ultraviolet (UV) spectrograph working in the 115–350 nm wavelength range with a spectral resolution of $\lambda/\Delta\lambda \sim 150–200$ (W. E. McClintock et al. 2015). The UV spectrograph consists of two detectors, an FUV detector (110–190 nm) and an MUV detector (180–340 nm). In the current study, the C⁺ emission line is at 133.5 nm, and thus the data analysis is performed on the FUV detector. We use the IUVS limb scan and periparse data of Level 1C, Version 13, Revision 1 to determine the altitudinal distribution of the C⁺ emission brightness (N. Schneider 2022).

The electron differential flux is obtained from the Solar Wind Electron Analyzer (SWEA) instrument. SWEA measures electron fluxes with electron energies from 3 eV to 4.6 keV with a time resolution of 2 s (D. L. Mitchell et al. 2016). A symmetric hemispheric electrostatic analyzer is designed to measure the energy and angular distributions of solar wind electrons and ionospheric suprathermal electrons on Mars. Since SWEA places a potential difference between two concentric hemispheres, it selects electrons within a specified energy range (Figure 3 in D. L. Mitchell et al. 2016). The center energy of the transmitted electrons is proportional to the potential difference, and the width of the energy response is proportional to the energy ($\Delta E/E = 0.17$). More details can be found in D. L. Mitchell et al. (2016). Our studies utilize SWEA Level 2, Version 5, Revision 1 for the SWEA electron spectrum (D. L. Mitchell 2023).

The neutral and ion densities in the thermosphere and ionosphere are obtained from the Neutral Gas and Ion Mass Spectrometer (NGIMS) instrument. NGIMS is a quadrupole mass spectrometer with a mass range of 2–150 Da and unit mass resolution designed to measure the composition, structure, and variability of the upper atmosphere on Mars (P. R. Mahaffy et al. 2015). It can detect both neutral and ion compositions at altitudes between 150 km (~125 km at “Deep Dip” (DD) campaigns) and 500 km with a vertical resolution below 5 km and an accuracy of less than 25% for most species. More details about the NGIMS instrument can be found in P. R. Mahaffy et al. (2015). The NGIMS data uses Level 2, Version 8, Revision 1 for the neutral and ion densities (M. K. Elrod et al. 2014).

### 2.2.2. Observations

Data obtained from in situ observations (SWEA and NGIMS) during the DD campaigns is used for the current study. Nine MAVEN DD campaigns were performed in total between 2015 February and 2018 April, during which the periparse altitude was lowered from ~150 to ~120 km for approximately one week. We use data obtained from DD8,
<figure><img src="image_5.png" /><figcaption>**Figure 1.** Orbit information for MAVEN in situ observations at (a), (c), (e), (g) orbit 5933 (DD8) and (b), (d), (f), (h) orbit 3586 (DD6). Each panel shows the (a)–(b) altitude, (c)–(d) solar zenith angle, (e)–(f) latitude, or (g)–(h) longitude.</figcaption></figure>

conducted from 2017 October 16 to 2017 October 22, and DD6, conducted from 2016 July 26 to 2016 August 2, to investigate $C^+$ emission on both the dayside (DD8) and nightside (DD6). These two orbits have been selected to facilitate comparisons between the dayside and nightside, due to their periapsis altitude of around 120 km, which is lower than that of other DD campaigns, and to avoid the terminator region. Figure 1 shows orbit information for MAVEN in situ observations at altitudes below 500 km for each orbit during DD8 (orbit 5933; 2017 October 20) and DD6 (orbit 3586; 2016 August 1). Note that the $C^+$ density is not available on the nightside in this case because it is not observed.

IUVS data are used for approximately one month, including each DD campaign; however, since IUVS did not observe in DD8 and DD6, we analyze IUVS data for 2017 October and 2016 July excluding the DD campaign period. Figure 2 shows

<figure><img src="image_6.png" /></figure>

geometry information at the tangential altitude of 152.5 km observed by IUVS in 2017 October and 2016 July.

## 3. Results

### 3.1. $C^+$ Emission Estimated from MAVEN IUVS Observations

The total brightness of the $C^+$ emission by all mechanisms of DI, PI, EI, and RS is obtained from the IUVS observations. IUVS data is analyzed during 2017 October and 2016 July excluding the MAVEN DD8 and DD6 campaigns. IUVS performed observations mostly on the dayside in 2017 October and mostly on the nightside in 2016 July in the “periapse” mode, which provides limb scan observations near the periapsis. Figure 3(a) shows the FUV spectrum or radiance as a function of wavelength in an altitude range of ~50–165 km observed by IUVS on board MAVEN during
<figure><img src="image_7.png" /><figcaption>**Figure 2.** Geometry information at the tangential altitude of 152.5 km observed by IUVS in 2017 October (blue) and 2016 July (red). Each panel shows the (a) latitude, (b) solar zenith angle, or (c) local time as a function of longitude.</figcaption></figure>

2017 October (blue line) and 2016 July (green line). The spectrum is integrated over the relevant period and divided by the number of observations. The spectrum is observed on the dayside but not on the nightside because it is likely below the detection threshold then. The magenta line denotes the spectrum after subtracting the detector background using the spectrum on the nightside (green line) and the line-spread function (LSF) coming from the hydrogen Lyα (121.6 nm; gray line). One of the sources of background counts could be produced inside the detector and is probably related to the energetic particle environment, which is corrected by the spectrum on the nightside. Another source is the far wing of the LSF coming from the Lyα (gray line). The far wing is about 1% of the radiance at a wavelength of 133.5 nm (gray line), which is negligible in this case. The calibrated radiance (magenta line) is therefore obtained by subtracting the night-side radiance (green line) from the dayside radiance (blue line). The adjacent oxygen emission lines at 130.4 nm and 135.6 nm are also very bright, which could be other sources of radiance. The C⁺ emission line is blended with the C-atom 132.9 nm line (as discussed later), and the tail of the 130.4 nm line is slightly distorted due to the influence of the C-atom 132.9 nm line. Considering this, if the 130.4 nm line is extrapolated from around 131.3 nm to approximately 133 nm, which is not affected by the C-atom 132.9 nm line, the radiance becomes less than 0.01 kR nm⁻¹, suggesting that the influence on the C⁺ emission line is minimal. Similarly, for the 135.6 nm line, even when extrapolated to around 134 nm, the radiance remains below 0.01 kR nm⁻¹, indicating that its influence on the C⁺ emission line is also minimal. Figures 3(b)–(d) show the FUV spectrum on the dayside at altitudes of 92.5–125 km, 125–150 km, and 150–165 km, respectively, with the brightest spectrum in the 125–150 km altitude range. Note that the trend of increasing radiance at longer wavelengths is due to the CO fourth-positive group bands. The IUVS detector background produces an approximately constant data number at all wavelengths. The Csl photocathode on the FUV detector is progressively less sensitive at longer wavelengths, causing a slope when attempting to “calibrate” the background contribution. It is a coincidence that this slope appears similar to the general trend in spectral structure of the Martian airglow over the wavelength range.

The spectrum needs to be integrated over the wavelength to compare the brightness estimated from the in situ observation. Figure 4 shows the FUV spectra at each altitude fitted for integration restricted to the C⁺ emission line regions. C⁺ emission is represented by a double Gaussian fitting because it is blended with the C-atom 132.9 nm emission line. Spectra between 132.6 and 134.0 nm are utilized for fitting in order to improve the accuracy of the fitting. The integration and fitting methods for C-atom emission are shown in the appendix (Figure A1). Figure 5 shows the brightness of C⁺ emission after the integration observed by IUVS in three altitude ranges (red solid line). The brightness has a peak at altitudes of 125–150 km and the peak value is 110 R. Note that this value might be overestimated because it does not consider the detector background. The detector background at wavelengths around the C⁺ emission is 10–20 R, and thus the actual brightness could be 90–100 R (red dashed line). Figure 5 also shows the brightness of C-atom emission (156.1 nm) during the same period of the C⁺ emission as a reference (orange solid line). The maximum brightness is ~650 R at altitudes of 125–150 km, but it decreases by 50–60 R when assuming a detector background (orange dashed line), which is nearly consistent with the brightness estimated by D. Y. Lo et al. (2022).

## 3.2. C⁺ Emission Estimated from MAVEN In Situ Observations

C⁺ emission by EI is estimated from the electron differential flux, emission cross sections, and CO₂ and CO density profiles. Figures 6(a) and (b) show the altitudinal profile of the suprathermal electron differential flux observed by MAVEN SWEA during orbit 5933 (DD8) and orbit 3586 (DD6). The differential flux basically peaks at ~6 eV at all altitudes on the dayside. The flux then reaches a maximum value of ~2 × 10⁸ cm⁻² s⁻¹ sr⁻¹ eV⁻¹ at 150 km altitude (Figure 6(a)). In contrast, on the nightside, the differential flux peaks at ~3 eV, but it reaches a maximum at altitudes of 225–250 km. The maximum value is 2 × 10⁶ cm⁻² s⁻¹ sr⁻¹ eV⁻¹, which is 2 orders of magnitude smaller than that on the dayside (Figure 6(b)). Figure 6(c) shows the C⁺ emission cross sections from CO₂ and CO dissociated and ionized by EI (S. V. Avakyan et al. 1998). The cross sections are interpolated/extrapolated into the energy bins of SWEA (red and blue dashed lines in Figure 6(c)). Both
C⁺ emission observed by MAVEN/IUVS

<figure><img src="image_8.png" /></figure>

<figure><img src="image_9.png" /></figure>

<figure><img src="image_10.png" /></figure>

<figure><img src="image_11.png" /></figure>

**Figure 3.** FUV spectra as a function of wavelength at (a) all altitudes, (b) 92.5–125 km altitude, (c) 125–150 km altitude, and (d) 150–165 km altitude observed by IUVS during 2017 October (blue line, dayside) and 2016 July (green line, nightside), which are integrated over the relevant period and divided by the number of observations with error bars (cyan lines). The spectrum after eliminating the detector background (magenta line) is shown with the LSF coming from the Lyα (121.6 nm; gray line).

cross sections have a peak around 150 eV and approach zero below ~30 eV, and the emission cross section of e-CO is larger than that of e-CO₂. The CO₂ and CO densities in the thermosphere are obtained from the NGIMS observations. Figure 7(a) shows the CO₂ (red and magenta) and CO (orange and yellow) densities for orbit 5933 on the dayside during DD8 and for orbit 3586 on the nightside during DD6. The CO₂ density is higher than the CO density below ~200 km altitude on the dayside, and lower above ~200 km. The trend in the altitude distribution of density on the nightside is the same as that on the dayside, but the reversal point drops to an altitude of ~150 km because of the small scale height on the nightside.

The brightness of C⁺ emission from CO₂ and CO by EI is estimated from Equation (2) using data from the MAVEN in situ observations as input (Figure 7(b)) as well as the volume emission rate (Figure 7(c)). The estimated brightness of the emission from e-CO collisions has a maximum value of ~100 R at 125 km altitude and decreases with altitude to as low as ~3 R at 150 km altitude. It is less than 1 R at altitudes above

170 km. The brightness of the emission from e-CO₂ collisions is ~30 R at 125 km altitude, fainter than that from e-CO collisions, and less than 1 R at altitudes above 150 km. It means that the emission cross section is more important than the neutral density in determining brightness.

The brightness of C⁺ emission by EI is estimated from the CHIANTI database (e.g., G. Del Zanna et al. 2021). CHIANTI requires the suprathermal electron density and temperature as inputs to estimate the brightness, assuming that the electron differential flux is described by the superposition of multiple Maxwellians. The current study fits three Maxwellians to the SWEA electron differential flux shown in Figure 6(a) (see Figure B1) and evaluates C⁺ emission from e-C⁺ collisions along with densities and temperatures obtained from the fitting. Figure 8 shows the brightness of C⁺ emission from e-C⁺ collisions (blue line) and RS (green line; panel (a)) as well as the volume emission rate (panel (b)). The brightness of C⁺ emission from e-C⁺ collisions has a peak around the 250 km
Fitting for C⁺ emission (133.5 nm)

<figure><img src="image_12.png" /></figure>

<figure><img src="image_13.png" /></figure>

<figure><img src="image_14.png" /></figure>

<figure><img src="image_15.png" /><figcaption>**Figure 4.** C⁺ emission at (a) all altitudes, (b) 92.5–125 km altitudes, (c) 125–150 km altitudes, and (d) 150–165 km altitudes is described by a double Gaussian fitting (orange and green), along with IUVS observations restricted near the C⁺ emission line (blue). The fitted curve is shown as a red line.</figcaption></figure>

altitude and the maximum value is ~0.01 mR, which is much fainter than that for the e-CO and e-CO₂ collision cases.

We can finally estimate the brightness of C⁺ emission due to RS from Equation (3) using the C⁺ density obtained from the NGIMS observation (blue dots in Figure 7(a)). The estimated brightness is very faint, with a maximum value of ~30 mR at an altitude of ~250 km (green line in Figure 8(a)). Note that the estimated RS brightness at altitudes below the peak altitude is not accurate because Equation (3) is only valid for an exponentially decreasing volume emission rate, i.e., the topside.

4. Discussion

IUVS observations have shown a brightness of C⁺ emission of ~90–100 R at altitudes of 125–150 km and ~70–80 R at altitudes of 150–165 km after removing the detector background (Figure 5). This brightness is the total C⁺ emission from all processes (i.e., DI, PI, EI, and RS). In situ observations allow us to estimate the emission by EI and RS as discussed in the previous section. The brightness of emission from e-CO

collisions obtained using Equation (2) is ~100 R at 125 km altitude and ~3 R at 150 km altitude, while the brightness of emission from e-CO₂ collisions is ~30 R at 125 km altitude and less than 1 R at above 150 km altitude (Figure 7(b)). Note that the brightness by other mechanisms estimated from the in situ observations is much fainter than that by EI with CO₂ and CO (Figure 8). The cross section of e-C collision is not available, and thus we cannot estimate the brightness of emission due to e-C collision. D. Y. Lo et al. (2022) estimated the C density in the Martian thermosphere, and the maximum density is 8 × 10⁵ cm⁻³ at ~140 km altitude, which is 3–4 orders of magnitude smaller than the CO₂ and CO density. This suggests that emission due to e-C collision should be ignored.

The contribution of each mechanism to the C⁺ emission can be evaluated by comparing the brightness estimated from remote-sensing and in situ observations. The brightness estimated from IUVS observations is ~90–100 R at around 130 km altitude (Figure 5), while the brightness estimated from the in situ observations is ~10–20 R (Figure 7(b)). The PI/DI contribution is inferred by adding up everything that can be calculated/estimated and subtracting it from the brightness
<figure><img src="image_16.png" /><figcaption>**Figure 5.** The brightness of C⁺ (red) and C-atom (orange) emissions before (solid) and after (dashed) correction for the detector background.</figcaption></figure>

<figure><img src="image_17.png" /><figcaption>**Figure 6.** Electron differential flux at (a) orbit 5933 on the dayside and (b) orbit 3586 on the nightside observed by SWEA. Line colors indicate differences in altitude. (c) The C⁺ 133.5 nm emission cross sections for e-CO₂ (black solid line) and e-CO (gray solid line) are also shown along with interpolated/extrapolated values into the SWEA energy bin (red dashed line for e-CO₂ and blue dashed line for e-CO).</figcaption></figure>
<figure><img src="image_18.png" /></figure>

<figure><img src="image_19.png" /><figcaption>**Figure 7.** (a) Densities of CO₂ (red for the dayside and magenta for the nightside), CO (orange for the dayside and yellow for the nightside), and C⁺ (blue for the dayside) observed by NGIMS. (b) Altitude dependence of the brightness of C⁺ emission by the e-CO₂ collision (blue and cyan) and the e-CO collision (green and light green) on the dayside (solid line) and nightside (dashed line), modeled using Equation (2). (c) Altitude dependence of the volume emission ratio of C⁺ emission by the e-CO₂ collision (blue and cyan) and the e-CO collision (green and light green) on the dayside (solid line) and nightside (dashed line).</figcaption></figure>

<figure><img src="image_20.png" /><figcaption>**Figure 8.** Altitude dependence of (a) C⁺ emission by e-C⁺ collisions (blue) and RS (green), and (b) volume emission rate.</figcaption></figure>
<figure><img src="image_21.png" /><figcaption>**Figure 9.** Altitude dependences of C⁺ emission by e-CO₂ collisions on the dayside (blue), e-CO collisions on the dayside (dark green), e-CO collisions considering the low CO density (light green), e-C⁺ collisions (orange), and RS (magenta) as well as IUVS observations after background correction including DI/PI.</figcaption></figure>

<table><caption>**Table 1.**
Estimated Brightness of C⁺ Emission by Each Process</caption><thead><tr><th>Emission Process</th><th>Estimated Brightness (Dayside)</th><th>Cross Section or Database</th></tr></thead><tbody><tr><td>EI with CO<sub>2</sub></td><td>~4 R at 135 km, ~0.3 R at 160 km</td><td>S. V. Avakyan et al. (1998)</td></tr><tr><td>EI with CO</td><td>~10 R at 135 km, ~1 R at 160 km<br>(NGIMS)</td><td>~1 R at 135 km, ~0.1 R at 160 km (MCD)</td><td>S. V. Avakyan et al. (1998)</td></tr><tr><td>EI with C</td><td>N/A</td><td>N/A</td></tr><tr><td>EI with C<sup>+</sup></td><td>~0.00001 R at 250 km (maximum)</td><td>CHIANTI</td></tr><tr><td>EI with CO<sub>2</sub> and CO<sup>+</sup></td><td>N/A</td><td>N/A</td></tr><tr><td>RS</td><td>~0.01 R at 250 km (maximum)</td><td>g-factor required (R. R. Meier 1990)</td></tr><tr><td>DI and PI of CO<sub>2</sub>, CO, and C</td><td>N/A</td><td>N/A</td></tr><tr><td>IUVS observation (total brightness of all processes)</td><td>~90 R at 135 km, ~70 R at 160 km</td><td>...</td></tr></tbody></table>

observed by IUVS, indicating that both EI and DI/PI contribute to C⁺ emission at this altitude, with EI contributing ~18% and PI/DI ~82% to emission. However, since the uncertainty of CO density observed by NGIMS might be large at lower altitudes, CO density is examined by the Mars Climate Database (MCD-LMD V6.1; F. González-Galindo et al. 2015), and it is found that the density is likely to be 1–2 orders of magnitude smaller at the altitude of 130 km (Appendix C). In that case, the EI contribution to C⁺ emission is less than 10%, suggesting that PI/DI dominates the emission. Note that the CO density from MCD-LMD is examined at a latitude of 10° S, east longitude of 320°, and solar longitude of 76.4°, which is approximately the same position as the NGIMS observation during orbit 5933. At around 160 km altitude, IUVS observations are estimated to be ~70 R (Figure 5) and the in situ

observations ~1 R (Figure 7), and thus EI contributes little to the emission, and DI/PI is more effective. This comparison between remote-sensing and in situ observations provides contributing emission mechanisms for each altitude. It is however difficult to separate each of the reactions of CO₂, CO, and C with the solar UV flux from this method because the IUVS observation provides only the total emission brightness. Table 1 summarizes the brightness for each emission mechanism. Figure 9 also shows the comparison of the brightness of C⁺ emission in each process.

Although the reactions of DI/PI cannot be isolated from the observation, we try to make a rough estimate of the “total emission cross section associated with DI/PI,” which is a cross section integrated over the wavelength in Equation (4), from the case of small CO density. Equation (4) can be transformed
into the following equation once the total emission cross section is assumed:

$$B \sim NL10^{-6}J\Sigma^{\mathrm{DI/PI}}, \quad (5)$$

where B is the brightness; J is the solar flux integrated over wavelengths up to the ionization threshold wavelength, $J = \int_{0}^{\lambda_{\mathrm{si}}} I_{\infty} \exp(-\tau) d\lambda$; $\Sigma^{\mathrm{DI/PI}}$ is the total emission cross section associated with DI/PI, $\Sigma^{\mathrm{DI/PI}} = \int_{0}^{\lambda_{\mathrm{si}}} \sigma_{s}^{\mathrm{DI/PI}} d\lambda$; N is the total density of CO$_2$ and CO (defined in Section 2.1.1); and L is the average length scale of CO$_2$ and CO. The total emission cross section is evaluated from the brightness of C$^+$ emission observed by MAVEN IUVS and Equation (5). The brightness of C$^+$ emission is ~90 R and ~70 R at the altitude of 130 km and 160 km from the IUVS observation, respectively (Figure 5). The solar flux, J, is calculated by the EUVAC model (an EUV flux model for aeronomic calculations; P. G. Richards et al. 1994), which requires F10.7, known as the index of solar activity with a solar radio flux at 10.7 cm, and the optical depth associated with the absorption cross sections (e.g., R. W. Schunk & A. F. Nagy 2009). We here use an F10.7 index of 75, which is archived on the LISIRD website.<fnref n="12" /> The total cross section, $\Sigma^{\mathrm{DI/PI}}$, roughly estimated from Equation (5) including the brightness at 130 km is $7 \times 10^{-21} \mathrm{cm}^2$, and that at 160 km altitude is ~1 $\times 10^{-19} \mathrm{cm}^2$. This total cross section should be independent of altitude, but a cross section that differs by more than an order of magnitude for different altitudes is obtained from this rough estimation. This might be due to the uncertainty of CO$_2$ and CO density and length scale that is utilized for the line-of-sight integration. It is also important to understand the energy dependence of relevant cross sections, and ideally each cross section related to DI/PI could be obtained from the experiment. However, this rough estimation suggests that the total cross section is on the order of ~$10^{-19}-10^{-20} \mathrm{cm}^2$ which might be smaller than the emission cross section of electron impact.

The total flux of suprathelial electrons can be estimated by transforming Equation (2) into the form of Equation (5) if the individual contributions of EI and DI/PI to C$^+$ emission brightness are known. The contribution of EI is however much smaller than the DI/PI contribution, and therefore the estimated total flux of suprathelial electrons is not possible from the C$^+$ emission. On the other hand, this method might be applied for C-atom emission (156.1 nm). C-atom emission mechanisms are suggested by D. Y. Lo et al. (2022) and they are dominated only by electron impact with CO$_2$ at altitudes below 150 km. This method, which uses C-atom emission, provides an estimate of the total flux of suprathelial electrons and is likely to be a good indicator of ionospheric variations and this will be investigated in the future. In addition, the total flux of suprathelial electrons estimated from C-atom emission might be utilized to separate the DI/PI and EI contributions to C$^+$ 133.5 nm emission. If the total flux of suprathelial electrons is known a priori, the contribution of EI to C$^+$ emission can be estimated and thus the DI/PI contribution can be evaluated by subtracting the brightness of C$^+$ emission due to EI from the brightness observed in the UV. Note that this method would only be helpful for the high CO density case.

C$^+$ emission seems in general to not be a tracer of ionospheric variations on Mars because C$^+$ emission mostly originates from DI/PI of CO$_2$ and CO, reflecting only spatiotemporal variations of CO$_2$ and CO in the Martian thermosphere, and the suprathelial electron distribution associated with electron impact is not available due to the low EI contribution to emission. On the other hand, C$^+$ emission is a tracer of CO$_2$ and CO in the Martian thermosphere as well as C-atom emission, and thus observations of C$^+$ emission allow us to investigate the effects of internal (e.g., dust storms) and external (e.g., solar XUV and solar wind conditions) sources on the Martian thermosphere. C-atom emission is a candidate ionospheric tracer on Mars, as discussed above. Global observations of C$^+$ and C-atom emission might allow us to track spatiotemporal changes in the ionosphere and thermosphere on Mars due to dust storms and changing solar conditions.

The Life-environmentology, Astronomy, and PlanetarY Ultra-violet Telescope Assembly (LAPYUTA) mission (F. Tsuchiya et al. 2024), which is recognized as the successor to the HISAKI satellite (K. Yoshioka et al. 2013; I. Yoshikawa et al. 2014), was selected as one of the candidates for the sixth Japan Aerospace Exploration Agency M-class mission. LAPYUTA is a UV space telescope with a diameter of 60 cm orbiting Earth scheduled for launch in the 2030s to probe UV spectroscopy in the spectral range between 110 and 190 nm with its spectral resolution of 0.02 nm, and it can thus measure C-atom and C$^+$ emissions. LAPYUTA can observe global C-atom and C$^+$ emissions, indicating that it is possible to investigate global ionospheric variation on Mars temporally and spatially.

## 5. Conclusions

The C$^+$ emission mechanism from the Martian upper atmosphere has been investigated using remote-sensing and in situ observations by MAVEN. The IUVS instrument on board MAVEN can measure the total brightness of C$^+$ emission from all mechanisms, including electron impact, DI, PI, and RS, where the brightness of the corrected C$^+$ emission is 90–100 R at altitudes of 125–150 km and the brightness of C-atom emission is ~600 R. C$^+$ emission could not be observed at altitudes above 165 km by IUVS. The contribution of electron impact and RS to C$^+$ emission is estimated using in situ observations using SWEA and NGIMS on board MAVEN. The brightness of the emission from e–CO collisions has a maximum value of ~100 R at 125 km altitude and decreases with altitude to less than 1 R at altitudes above 170 km. The estimated brightness of the emission from e–CO$_2$ collisions is fainter than that from e–CO collisions. It means that the emission cross section is more important than the neutral density in determining brightness. The brightness of C$^+$ emission by e–C$^+$ collision and RS is at most less than 1 R and thus can be neglected. We cannot evaluate the brightness of the emission by e–C collision because the emission cross section is not available, but since the C density is at least 3 orders of magnitude smaller than the CO$_2$ and CO densities (D. Y. Lo et al. 2022), it must also be a very minor contribution.

The contribution of DI and PI to C$^+$ emission is estimated from both remote-sensing and in situ observations. The brightness estimated from IUVS observations is ~90–100 R at around 130 km altitude, while the brightness estimated from the in situ observations is ~10–20 R, indicating that both electron impact and DI/PI contribute to C$^+$ emission at this

<footnote n="12">¹² https://lasp.colorado.edu/lisird/data/noaaanelflux</footnote>
altitude. However, since the uncertainty in the CO density
observed by MAVEN NGIMS appears to be large, particularly
at lower altitudes, the CO density is examined by MCD-LMD
and found to be likely smaller than the density observed by
NGIMS at 130 km altitude. The contribution of electron impact
to C⁺ emission is less than 10%, suggesting that DI/PI
dominates the emission. In contrast, IUVS observations are
estimated to be ~25 R and the in situ observations ~1 R at
around 160 km altitude, and thus DI/PI is more effective. It is
however difficult to separate each of the reactions of CO₂, CO,
and C with the solar UV flux from this method.

An attempt is considered to determine the total flux of
suprathermal electrons in the Martian ionosphere from only C⁺
remote-sensing observations using ultraviolet radiation. However,
the total flux could not be estimated because it is not possible to
separate the emission contribution from electron impact and from
DI/PI in the lower ionosphere and thermosphere. This method
might be applied using observations of C-atom emission because
the C-atom emission mechanism is better understood than the C⁺
emission mechanism and the electron impact of CO₂ is dominant
in the lower ionosphere. If the total flux of suprathermal electrons
can be determined from global UV observations, it will be
possible to study the temporal and spatial variation in the
ionosphere globally using the UV telescopes proposed in the
LAPYUTA project and other projects.

In summary, observations of both C⁺ and C-atom emissions
might allow us to separate the DI/PI and EI contributions to C⁺
emission, and the total flux of suprathermal electrons can be
estimated from C-atom emission, which suggests that C⁺ and
C-atom emissions might be tracers of the Martian ionospheric
and thermospheric variation associated with the CO₂ and CO
atmospheres.

Acknowledgments

The MAVEN mission is supported by NASA through the
Mars Exploration Program in association with the University of
Colorado and NASA's Goddard Space Flight Center. The
MAVEN IUVS processed (Level 1c) periapse limb data are
publicly available in FITS format on the NASA Planetary Data
System (N. Schneider 2022), identified by “periapse” with
version/revision tag v13_r01. The MAVEN SWEA and
NGIMS data are publicly released in D. L. Mitchell (2023)
and M. K. Elrod et al. (2014), respectively. Parts of this work
for the observations obtained with the SWEA instrument are
supported by the French space agency CNES (National Centre
for Space Studies). This work was supported by JSPS
KAKENHI grant Nos. JP22KK0044, JP22K03695,
JP22K03708, and JP22H00164. The work of M.B. is supported
by NASA under award No. 80GSFC21M0002.

**Appendix A**

FUV Spectrum Fitted to Calculate Brightness of C-atom
Emission

Integration in the wavelength direction is necessary to
calculate the emission brightness from the FUV spectra
observed by MAVEN IUVS. Figure A1 shows the FUV
spectra at each altitude fitted for integration restricted to the
C-atom emission line regions. C-atom emission is represented
by a single Gaussian fitting, in contrast to blended C+
emission. Spectra between 155.8 and 156.5 nm are utilized
for fitting in order to improve the accuracy of the fitting. The
“fit function” in MATLAB is used for the fitting.
Fitting for C-atom emission (156.1 nm)

<figure><img src="image_22.png" /></figure>

<figure><img src="image_23.png" /></figure>

<figure><img src="image_24.png" /></figure>

<figure><img src="image_25.png" /><figcaption>Figure A1. Same as Figure 4, but the fitting is described by a single Gaussian.</figcaption></figure>

Appendix B

Differential Fluxes of Suprathermal Electrons Fitted to
Estimate the Brightness of C⁺ Emission by e-C⁺ Collisions

The brightness of C⁺ emission by e–C⁺ collisions is estimated
from the CHIANTI database (e.g., G. Del Zanna et al. 2021).
CHIANTI requires the suprathermal electron density and

temperature as inputs to estimate the brightness, assuming that
the electron differential flux is described by the superposition of
multiple Maxwellians. Figure B1 shows the differential flux of
suprathermal electrons fitted by a triple Maxwellian at each altitude.
This fitting is conducted by visual inspection, but there is no
significant deviation from the flux observed by MAVEN SWEA.
<figure><img src="image_26.png" /><figcaption>**Figure B1.** The differential flux of suprathermal electrons fitted by a triple Maxwellian (red, green, and orange lines) is shown with the SWEA observations (blue line) at each altitude. The magenta line denotes the sum of the three Maxwellians.</figcaption></figure>

## Appendix C
Uncertainty in the CO Density

The uncertainty in the CO density observed by MAVEN NGIMS appears to be large. It is compared to the CO density obtained by the MCD-LMD model. Figure C1 compares the CO densities obtained with MAVEN NGIMS (orange) and

MCD-LMD (black) in two time periods at the same latitude (10° S), east longitude (320°), and solar longitude (76.4°) as the MAVEN in situ observation during orbit 5933 (solid line), and at the same latitude, east longitude, and perihelion (251°; dashed line). The two model densities are 1–2 orders of magnitude smaller than the NGIMS density at the altitude of 130 km.
<figure><img src="image_27.png" /><figcaption>**Figure C1.** The CO density of the MAVEN NGIMS observations (orange) is compared to the MCD-LMD calculations at the same location as the MAVEN in situ observations (black solid line) and the perihelion (black dashed line).</figcaption></figure>

## ORCID iDs

Shotaro Sakai https://orcid.org/0000-0001-9135-2076

Hiromu Nakagawa https://orcid.org/0000-0001-8166-0171

Justin Deighan https://orcid.org/0000-0003-3667-902X

Sonal K. Jain https://orcid.org/0000-0002-1722-9392

Kei Masunaga https://orcid.org/0000-0001-9704-6993

Fuminori Tsuchiya https://orcid.org/0000-0003-3386-6794

Naoki Terada https://orcid.org/0000-0001-5685-9736

Majd Mayyasi https://orcid.org/0000-0002-5663-602X

Nicholas M. Schneider https://orcid.org/0000-0001-6720-5519

David L. Mitchell https://orcid.org/0000-0001-9154-7236

Christian Mazelle https://orcid.org/0000-0001-5332-9561

Robert J. Lillis https://orcid.org/0000-0003-0578-517X

Go Murakami https://orcid.org/0000-0002-2759-7682

Shannon M. Curry https://orcid.org/0000-0002-7463-9419

Kanako Seki https://orcid.org/0000-0001-5557-9062

## References

Avakyan, S. V., Il'in, R. N., Lavrov, V. M., & Ogurtsov, G. N. 1998, Collison Processes and Excitation of UV Emission from Planetary Atmospheric Gases: A Handbook of Cross Sections (Amsterdam: Gordon and Breach)

Benna, M., Mahaffy, P. R., Grebowsky, J. M., et al. 2015, *GeoRL*, 42, 8958

Brain, D. A., Barabash, S., Bougher, S. W., et al. 2017, in The Atmosphere and Climate of Mars, ed. R. M. Harberle et al. (Cambridge: Cambridge Univ. Press), 464

Chaffin, M. S., Chaufray, J.-Y., Stewart, I., Montmessin, F., Schneider, N. M., & Bertaux, J.-L. 2014, *GeoRL*, 41, 314

Chassefière, E., & Leblanc, F. 2004, *P&SS*, 52, 1039

Clarke, J. T., Bertaux, J.-L., Chaufray, J.-Y., et al. 2014, *GeoRL*, 41, 8013

Cravens, T. E., & Green, A. E. S. 1978, *Icar*, 33, 612

Cravens, T. E., Rahmati, A., Fox, J. L., et al. 2017, *JGRA*, 122, 1102

Del Zanna, G., Dere, K. P., Young, P. R., & Landi, E. 2021, *ApJ*, 909, 38

Dere, K. P., Landi, E., Mason, H. E., Monsignori Fossi, B. C., & Young, P. R. 1997, *A&AS*, 125, 149

Elrod, M. K., Benna, M., & Navas, T. 2014, MAVEN Neutral Gas and Ion Mass Spectrometer Data, NASA Planetary Data System, doi:10.17189/1518931

Goldspiel, J. M., & Squyres, S. W. 1991, *Icar*, 89, 392

González-Galindo, F., Chaufray, J.-Y., López-Valverde, M. A., et al. 2015, *JGRE*, 118, 2105

Gregory, B. S., Elliott, R. D., Deighan, J., Gröller, H., & Chaffin, M. S. 2023, *JGRE*, 128, e2022JE007576

Halekas, J. S. 2017, *JGRE*, 122, 901

Hunten, D. M., & McElroy, M. B. 1970, *JGR*, 75, 5989

Inui, S., Seki, K., Namekawa, T., et al. 2018, *GeoRL*, 45, 5283

Inui, S., Seki, K., Sakai, S., et al. 2019, *JGRA*, 124, 5482

Jakosky, B. M., Grebowsky, J. M., Luhmann, J. G., et al. 2015, *Sci*, 350, aad0210

Jakosky, B. M., Pepin, R. O., Johnson, R. E., & Fox, J. L. 1994, *Icar*, 111, 271

Kite, E. S. 2019, *SSRv*, 215, 10

Lillis, R. J., Deighan, J., Brain, D., et al. 2022, *GeoRL*, 49, e2022GL099820

Lillis, R. J., Deighan, J., Fox, J. L., et al. 2017, *JGRA*, 122, 3815

Lo, D. Y., Yelle, R. V., Deighan, J. I., et al. 2022, *Icar*, 371, 114664

Lo, D. Y., Yelle, R. V., Lillis, R. J., & Deighan, J. I. 2021, *Icar*, 360, 114371

Mahaffy, P. R., Benna, M., Elrod, M., et al. 2015, *GeoRL*, 42, 8951

Mahaffy, P. R., Benna, M., King, T., et al. 2015, *SSRv*, 195, 49

Malin, M. C., & Edgett, K. S. 2003, *Sci*, 302, 1931

McClintock, W. E., Schneider, N. M., Holsclaw, G. M., et al. 2015, *SSRv*, 195, 75

McElroy, M. B., & Donahue, T. M. 1972, *Sci*, 177, 986

Meier, R. R. 1990, *GeoRL*, 17, 1613

Mitchell, D. L. 2024, MAVEN SWEA Survey Rate Omni-Directional Electron Energy Spectra Data Collection, NASA Planetary Data System, doi:10.17189/3khv-zp66

Mitchell, D. L., Mazelle, C., Sauvaud, J.-A., et al. 2016, *SSRv*, 200, 495

Pickett, N. B., McFadden, J. P., Fowler, C. M., Hanley, H. G., & Benna, M. 2022, *JGRA*, 127, e2021JA029635

Ramstad, R., Barabash, S., Futaana, Y., Nilsson, H., & Holmström, M. 2018, *JGRE*, 123, 3051

Raymond, J. C., Blair, W. P., Long, K. S., et al. 1997, *ApJ*, 482, 881

Richards, P. G., Fennelly, J. A., & Torr, D. G. 1994, *JGR*, 99, 8981

Schneider, N. 2022, Mars Atmosphere and Volatile Evolution (MAVEN) Imaging Ultraviolet Spectrograph (IUVS) Processed-level Data Product Bundle, NASA Planetary Data System, doi:10.17189/1518964

Schunk, R. W., & Nagy, A. F. 2009, Ionospheres: Physics, Plasma Physics, and Chemistry (2nd ed.; Cambridge: Cambridge Univ. Press)

Tsuchiya, F., Murakami, G., Yamazaki, A., et al. 2024, *Proc. SPIE*, 13093, 130930I

Wordsworth, R. 2016, *AREPS*, 44, 381

Yoshikawa, I., Yoshioka, K., Murakami, G., et al. 2014, *SSRv*, 184, 237

Yoshioka, K., Murakami, G., Yamazaki, A., et al. 2013, *P&SS*, 85, 250