The dayside ionosphere of Mars: Comparing a one-dimensional photochemical model with MAVEN Deep Dip campaign observations

The dayside ionosphere of Mars: Comparing a one-dimensional photochemical model with MAVEN Deep Dip campaign observations

Icarus 337 (2020) 113502 Contents lists available at ScienceDirect Icarus journal homepage: www.elsevier.com/locate/icarus The dayside ionosphere o...

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Icarus 337 (2020) 113502

Contents lists available at ScienceDirect

Icarus journal homepage: www.elsevier.com/locate/icarus

The dayside ionosphere of Mars: Comparing a one-dimensional photochemical model with MAVEN Deep Dip campaign observations Vrinda Mukundan a ,∗, Smitha V. Thampi a , Anil Bhardwaj b , C. Krishnaprasad a a b

Space Physics Laboratory, Vikram Sarabhai Space Centre, Thiruvananthapuram 695022, India Physical Research Laboratory, Ahmedabad 380009, India

ABSTRACT A one dimensional photochemical model for the dayside ionosphere of Mars has been developed for calculating the density profiles of ions and electrons under steady state photochemical equilibrium condition. The study focuses on the Deep Dip campaigns of the Mars Atmosphere and Volatile EvolutioN mission (MAVEN) and used the in-situ measurements of neutral density profiles, solar flux and electron temperatures from instruments onboard MAVEN as input to the model. An energy deposition model is employed for calculating the attenuated photon flux and photoelectron flux at different altitudes in the ionosphere. The Analytical Yield Spectrum approach is used for calculating the photoelectron fluxes. Volume production rates of major primary ions, CO+2 , CO+ , O+ , C+ , N+2 , and N+ , due to photon and photoelectron impact are calculated and used as input to the model in which ion-neutral chemistry in the dayside ionosphere is simulated. The modeled ion profiles are compared with the ion mode observations of Neutral Gas Ion Mass Spectrometer (NGIMS) and electron density estimates from Langmuir Probe and Waves (LPW). The model could reproduce the observed structure of the major ion profiles O+2 , CO+2 , and electron density reasonably well. However, when the magnitudes are compared, the modeled values are roughly a factor of 2 larger than observation. By reducing the neutral CO2 density, the model outputs and the observed ion profiles and electron density profile can become closer. The model also calculated the densities of 11 other ions, viz. N+ , C+ , O+ , NO+ , N2 H+ , HCO+ , N+2 , CO+ , OCOH+ , HNO+ , and OH+ . Profiles of each of these ions are compared with the NGIMS observations during the respective orbits and are discussed in detail. Such a comparison for the deep dip periods is reported for the first time, which showcases the level of the current understanding of the ion chemistry in the Martian ionosphere.

1. Introduction The Martian ionosphere is mainly produced by the ionization of the atmospheric neutral gases by solar EUV-photons and EUV-producedphotoelectrons. Collisions of these photons and photoelectrons with the dominant neutral molecules, such as CO2 , CO, and N2 , result in the formation of ionized species, like CO+ , O+ , and N+ . They initiate 2 2 2 a chain of chemical reactions in the upper atmosphere, including dissociation, ionization, charge exchange and recombination reactions and results in an ionosphere with O+ as the most dominant ion (Hanson 2 et al., 1977). The maximum electron density is of the order of 105 cm−3 and is located at ∼125–140 km (Haider et al., 2011). The observational studies of the Martian ionosphere began with the Viking 1 and 2 landers in 1976. The plasma profiles measured by these Viking landers were the only available in-situ measurements of the Martian ionosphere until very recently (Hanson et al., 1977). A very long gap occurred until the arrival of the Mars Global Surveyor spacecraft (MGS) in 1998. Using the Radio Occultation (RO) technique, the Radio Science experiment on board MGS reported more than 5000 electron density profiles for the altitude range 90–200 km (e.g. Hinson et al. (1999), Bougher et al. (2004)). The Mars Express (MEX) mission,

launched in 2003, has two instruments namely, the Mars Advanced Radar for Subsurface and Ionospheric Sounding (MARSIS) and the Mars Express Radio Science Experiment (MaRS), that measures the electron density profiles. The synthetic aperture radar MARSIS measured about 40000 electron density profiles, but only for the region above the main ionospheric peak (Morgan et al., 2008). MaRS uses radio occultation method to obtain vertical profiles of electron densities from the base to the top of the ionosphere (Patzold et al., 2016). However, the RO method is sensitive only to electron densities and cannot provide information on ion composition. Presently, the Mars Atmosphere and Volatile Evolution Mission (MAVEN) (September 2014 — ongoing) is providing in-situ composition measurements and unprecedented data on Martian thermosphere and ionosphere (Jakosky et al., 2015). Even though the typical orbital periapsis altitude is ∼150 km, there are many deep dip (DD) campaigns during which the periapsis gets lowered to 120–130 km so that the measurements near the ionospheric peak could be made. The instruments onboard MAVEN, such as Neutral Gas and Ion Mass Spectrometer (NGIMS) (Mahaffy et al., 2015), Langmuir Probe and Waves (LPW) (Andersson et al., 2015), Solar Wind Electron Analyzer (SWEA) (Mitchell et al., 2016) are making in-situ measurements of the upper atmosphere and plasma environment and have been

∗ Corresponding author. E-mail addresses: [email protected] (Vrinda M.), [email protected] (Smitha V.T.), [email protected] (A. Bhardwaj).

https://doi.org/10.1016/j.icarus.2019.113502 Received 4 May 2019; Received in revised form 11 October 2019; Accepted 17 October 2019 Available online 19 October 2019 0019-1035/© 2019 Elsevier Inc. All rights reserved.

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so as to reproduce not only the profiles of the major ions, i.e., O+ 2 and CO+ , but also the density profiles of 11 other ions, viz. N+ , C+ , 2 + + + + + + + + + O , NO , N2 H , HCO , N2 , CO , OCOH , HNO , and OH as well as the electron density. Comparison of the modeled density profiles of minor ions with NGIMS ion composition observations are done for the first time. The differences and similarities between the model and the observed density profiles of major as well as minor ions are discussed in detail in the subsequent sections. 2. Model description Our methodology is same as that adopted by Mukundan and Bhardwaj (2018b,a), which is briefly described here. The rate of change of number density of an ion species 𝑖 is given by the continuity equation 𝑑𝑛𝑖 + ▽.𝜙 = 𝑃𝑖 − 𝐿𝑖 𝑑𝑡

Fig. 1. MAVEN EUVM Solar flux used as input to the model. Flux available for the day of occurrence of the stated orbit is used in each case.

(1)

where ▽.𝜙 represents the change due to transport, and P𝑖 and L𝑖 are the production and loss rate of the ion. In the photochemical equilibrium region where the transport processes are not important, under steady state conditions, the continuity equation reduces to

providing valuable data to improve our understanding of the Martian upper atmosphere and its interaction with the solar wind. Photochemical models have been used for studying the Martian ionosphere in the photochemical equilibrium region (Haider et al., 2011) with the aim of calculating the abundance of the different chemical species and thus to have a better understanding of the upper atmospheric and ionospheric composition. Comparing the model calculations with observations helps to validate our understanding of the physical and chemical processes occurring in the atmosphere. The differences and similarities between model and the observations give insight into whether our current understanding of the atmospheric processes are capable of forming the observed features of the Martian ionosphere. Fox et al. (2015), using the neutral density profile from Viking 1 and modeled electron and ion temperatures, developed a photochemical model for Mars. They predicted the ion density profiles, including water ions, in the Martian ionosphere and compared them with NGIMS observations. The study of Xu et al. (2018) computed the density profiles of the ions O+ and CO+ using MAVEN observations as 2 2 input to the model and compared the modeled ion density profiles with MAVEN observations. But Xu et al. (2018) restricted their discussion to the major ions as the focus of their study was to understand the response of the Martian ionosphere to solar flares. In the present study, we develop a one-dimensional photochemical model for the dayside Martian ionosphere. The observations from the instruments onboard MAVEN, viz. neutral densities by NGIMS, solar fluxes from Extreme Ultraviolet Monitor (EUVM) (Eparvier et al., 2015), and electron temperatures from LPW are used as inputs to the model. The goal of the study is to calculate the ion and electron density profiles and compare them with the plasma measurements of NGIMS and LPW. The study focuses on the photochemical equilibrium region which occurs at altitudes ≤200 km (Schunk and Nagy, 2000). The Deep Dip (DD) campaigns of MAVEN explored this region of the Martian ionosphere–thermosphere as the spacecraft went deeper into the atmosphere making in-situ measurements of the region near the ionospheric peak. Out of the 9 DD campaigns that have happened to date, three campaigns, namely DD2, DD4, and DD8, have occurred in the Martian dayside with solar zenith angle <90◦ . Two orbits from each of these campaigns were selected and the MAVEN measurements of solar UV flux, neutral densities and electron temperatures for those respective orbits were used as model inputs. The modeled photoelectron fluxes are compared with the actual observations from MAVEN. The ion and electron density profiles estimated by the model are compared with the observations of NGIMS and LPW during the respective orbits. The details of these chosen orbits during each DD campaign are given in Table 1. The model is developed with a detailed ion-neutral chemistry

𝑃𝑖 = 𝐿𝑖

(2)

As the constituents of the upper atmosphere are not well mixed, this equation is solved for each ion species at each altitude independently for calculating the densities. The photochemical equilibrium prevails in Martian atmosphere at altitudes below 200 km (Schunk and Nagy, 2000). In this region, the plasma density is controlled mainly by chemistry and hence transport processes can be neglected. An energy deposition model is developed which describes the absorption of solar EUV radiation using Beer Lambert law. The initial inputs to the model are solar EUV radiation reaching the top of the atmosphere, the altitudinal distribution of the neutral constituents and the photoabsorption and photoionization cross sections. The attenuated solar flux, thus obtained is used for calculating photoelectron production rate. This is subsequently used in the calculation of steady state photoelectron flux by employing an Analytical Yield Spectrum (AYS) approach (Bhardwaj and Micheal, 1999a,b). Primary production rate or the volume ionization rate of major primary ions viz. CO+ , CO+ , O+ , C+ , N+ , and N+ , due to photon impact as well as 2 2 due to photoelectron impact is calculated. The ion production rate is used as input to the photochemical model which includes ion-neutral chemistry. The production and loss reactions of various ions are taken in the model to calculate the number density profiles of different ionic species and electrons. 3. Model inputs To estimate the solar ionizing flux at Mars we used the EUVM level 3 data product (data version 11, revision 03) which is a modeled spectral irradiance based on a modified version of the Flare Irradiance Spectral Model (FISM) called the FISM-P (Thiemann et al., 2017; Chamberlin et al., 2007, 2008). The model uses measurements from MAVEN’s EUVM as well as proxies measured at the Earth that incorporates the Earth-Sun-Mars angle corresponding to the time of the MAVEN EUVM data. The modeled irradiances are available for the wavelengths 0–190 nm with a resolution of 1 nm. In the model, we have considered the solar flux in the wavelength range 1–100 nm, which causes ionization of the major neutral gases in the Martian atmosphere. EUVM flux available for the day of occurrence of the respective orbit is used as model input (See Fig. 1). Seven neutral gases are considered in the model, viz., CO2 , N2 , CO, O, O2 , H2 , and NO. Our model uses the level 2 density profiles (version 14, revision 01) for CO2 , N2 , O, and CO measured by NGIMS during the inbound leg of the stated orbits. The NO and O2 densities measured by 2

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Fig. 2. Density profiles of neutral gases used in the model. For CO2 , N2 , O, and CO, we have used MAVEN-NGIMS measurements during the respective orbits. NO and O2 densities are from Mahaffy et al. (2015) (for DD2 and DD8 orbits ) and Vogt et al. (2017) (DD4 orbits). H2 density profile is taken from Krasnopolsky (2002). See text for more details.

NGIMS are currently not available from Planetary Data System (PDS). For the DD2 and DD8 orbits, we have used the NO and O2 density profile as given in Mahaffy et al. (2015) for the orbit 1064. For DD4 orbits, the NO and O2 profile as shown in Vogt et al. (2017) for the orbit 1835 is used. Measured H2 density profiles are not available for the Martian atmosphere. Krasnopolsky (2002) modeled the profiles of several neutral and ionic species for different levels of solar activity.

Since the deep dip campaigns that we have considered in the study have occurred during the medium solar activity condition, we used the modeled H2 profile of Krasnopolsky (2002) for the medium solar activity for all the DD orbits considered in the model. Fig. 2 shows the density profiles of neutral gases that we have used in the model. Electron temperature (T𝑒 ) is an important parameter for calculating the loss rates of various ions via electron recombination reaction. 3

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Fig. 3. Electron temperature profiles measured by the LPW instrument during the deep dip (DD) campaigns included in the study. The yellow lines show the profile measured during individual orbit and the solid black line shows the average profile.

Table 1 Information on the MAVEN Deep Dip (DD) campaigns used in our study. DD#

Orbits

a

2 4 8

1059–1086 1801–1838 5909–5950

10 91 25

SZA (◦ )

b

Peri (km)

130 123 128

c

d

L𝑠 (◦ )

328.6 37.5 76.3

Orbits [e SZA (◦ ), Day] used

Lat (◦ )

1064 (10◦ , 18/04/2015) & 1084 (7.8◦ , 22/04/2015) 1815 (83◦ , 04/09/2015) & 1824 (81◦ , 06/09/2015) 5912 (28◦ ,16/10/2017) & 5934 (23◦ , 20/10/2017)

3.8 S 63.9 S 18.9 N

a

SZA: Typical Solar Zenith Angle value at the periapsis of different orbits. Peri: Periapsis altitude. c L𝑠 : Solar longitude. d Lat: Latitude at periapsis. e SZA: Value of SZA used in the model. b

We have used the T𝑒 measurements of LPW instrument (version 14, revision 01) for calculating the ion–electron recombination rates. Large fluctuations are seen in the T𝑒 profiles of individual orbits. Hence we took the average of the T𝑒 profiles measured during the inbound legs of all orbits of a DD campaign. It is this averaged profile for each of the DD campaigns that is used as input in the model (see Fig. 3). Another key input to the model is the cross sections of the major atmospheric constituents for the photon and photoelectron impact processes for calculating the attenuated solar flux and photoelectron flux. The model consider the photoionization of four dominant gases; CO2 , N2 , CO, and O. The photon cross sections for CO2 , CO, and N2 are taken from the database http://amop.space.swri.edu (Huebner et al., 1992). For O2 , photon cross sections are taken from Fennelly and Torr (1992). Inelastic cross sections for the electron impact on CO2 are taken from Bhardwaj and Jain (2009), while those for CO, N2 , and O are taken from the work of Jackman et al. (1977). The list of ion-neutral and dissociative recombination reactions used in the model, their rate coefficients, and references from which they are taken are summarized in Table 2.

where, 𝜏(𝜆, 𝑍) =



𝜎𝑙𝐴 (𝜆)

∫𝑍

𝑛𝑙 (𝑍 ′ )𝑑𝑍 ′ ,

(5)

is the optical depth of wavelength 𝜆 at altitude 𝑍, 𝜎𝑙𝐴 is total photoabsorption cross section of the constituent 𝑙 at 𝜆, 𝐼(∞, 𝜆) is the unattenuated solar flux at 𝜆 and 𝜒 is the solar zenith angle (SZA). The sec(𝜒) approximation, instead of the Chapman grazing function, is valid for small SZAs (<80◦ ). For DD2 and DD8 campaigns, the observations are at a very low SZA and hence the sec(𝜒) approximation can be easily used for both the inbound and outbound legs. For these orbits, we use the value of SZA when the spacecraft was at the altitude of closest approach during the inbound leg. For the DD4 campaign, only the inbound legs of the orbits were located on the dayside and SZA at the orbital periapsis was in the range 89–95◦ . However, the SZA lies in the range 79–85◦ when MAVEN was making observations at the ionospheric peak (∼124 km). For the DD4 orbits 1815 and 1824, we use the SZA values 83.37◦ and 81.35◦ , respectively, for our calculations which are the SZAs at the time of measurement of ionospheric peak altitude during the respective orbits. The value of SZA used in the model for different orbits is listed in Table 1. The primary photoelectrons generated in the atmosphere during photoionization lose their energy by colliding with the background neutrals. Using the Analytical Yield Spectrum (AYS) approach, we calculate the degraded electron spectrum. The photoelectron flux is calculated as:

4.1. Model-observation comparison of photoelectron flux Photoelectron production rate at an altitude 𝑍 is calculated using ∑ ∑ 𝑄(𝑍, 𝜆) = 𝑛𝑙 (𝑍) 𝜎𝑙𝐼 (𝑗, 𝜆)𝐼(𝑍, 𝜆) (3)

100

𝜓(𝑍, 𝐸) =

𝑗,𝜆

∫𝐸

𝑄(𝑍, 𝐸) 𝑈𝑐 (𝐸, 𝐸0 ) 𝑑𝐸0 ∑ 𝑙 𝑛𝑙 (𝑍) 𝜎𝑙 (𝐸)

(6)

Here Q(Z, E) is the photoelectron production rate at altitude Z, which is calculated using Eq. (3). The integral is over the incident electron energy. While integrating, we consider electron energies only up to 100 eV since the production rate of photoelectrons falls down by several orders of magnitude for energies around and beyond 100 eV. 𝑈𝑐 (𝐸, 𝐸0 ) is the composite yield spectrum, which is the yield spectrum of a

Here 𝑛𝑙 (𝑍) is the neutral density of constituent 𝑙 at altitude Z, 𝜎𝑙𝐼 (𝑗, 𝜆) is total photoionization cross section of the 𝑗th ionization state of the 𝑙th constituent at wavelength 𝜆, and I(Z, 𝜆) is the solar flux at altitude Z which is calculated as 𝐼(𝑍, 𝜆) = 𝐼(∞, 𝜆) 𝑒𝑥𝑝[−𝜏]

∑ 𝑙

4. Model results

𝑙

𝑠𝑒𝑐(𝜒)

(4) 4

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Vrinda M. et al. Table 2 The list of ion-neutral reactions used in the model. Reaction

Rates (cm3 s−1 )

References for rate coefficients

CO+2 + O → O+2 + CO CO+2 + O → O+ + CO2 CO+2 + O2 → O+2 + CO2 CO+2 + NO → NO+ + CO2 CO+2 + H2 → OCOH+ + H O+2 + NO → NO+ + O2 N+2 + O2 → O+2 + N2 N+2 + NO → NO+ + N2 N+2 + CO2 → CO+2 + N2 N+2 + O → NO+ + N N+2 + CO → CO+ + N2 N+2 + H2 → N2 H+ + H O+ + CO2 → O+2 + CO O+ + O2 → O+2 + O O+ + NO → NO+ + O O+ + N2 → NO+ + N O+ + H2 → OH+ + H C+ + CO2 → CO+ + CO CO+ + CO2 → CO+2 + CO CO+ + O → O+ + CO CO+ + H2 → HCO+ + H N+ + CO2 → CO+2 + N N2 H+ + CO → HCO+ + N2 N2 H+ + O → OH+ + N2 N2 H+ + NO → HNO+ + N2 N2 H+ + CO2 → OCOH+ + N2 OH+ + CO → HCO+ + O OH+ + O → O+2 + H OH+ + N2 → N2 H+ + O OH+ + CO → CO+ + OH OH+ + NO → HNO+ + O OH+ + CO2 → OCOH+ + O OH+ + NO → NO+ + OH OH+ + O2 → O+2 + OH HNO+ + CO → HCO+ + NO HNO+ + NO → NO+ + HNO HNO+ + CO2 → OCOH+ + NO OCOH+ + O → HCO+ + O2 OCOH+ + CO → HCO+ + CO2 OCOH+ + N2 → N2 H+ + CO2 CO+2 + e− → CO + O

1.64 × 10−10 9.60 × 10−11 5.00 × 10−11 1.20 × 10−10 8.70 × 10−10 4.6 × 10−10 5.0 × 10−11 4.1 × 10−10 3.0 × 10−10 1.3 × 10−10 7.4 × 10−11 2.0 × 10−09 1.1 × 10−09 2.0 × 10−11 8.0 × 10−13 1.2 × 10−12 1.35 × 10−09 1.1 × 10−09 1.0 × 10−09 1.4 × 10−10 7.5 × 10−10 1.00 × 10−09 8.8 × 10−10 1.4 × 10−10 3.4 × 10−10 1 × 10−09 3.55 × 10−10 7.1 × 10−10 2.4 × 10−10 3.55 × 10−10 6.11 × 10−10 1.4 × 10−09 3.59 × 10−10 3.8 × 10−10 8.6 × 10−10 7 × 10−10 9.45 × 10−10 1.0 ×10−09 7.8 ×10−10 1.37×10−09 4.2 × 10−07 × ( 300 )0.75 𝑇

Schunk and Nagy (2000) Schunk and Nagy (2000) Fox and Dalgarno (1979) Schunk and Nagy (2000) Scott et al. (1997) Schunk and Nagy (2000) Schunk and Nagy (2000) Schunk and Nagy (2000) Raina and Haider (1998) Schunk and Nagy (2000) Le-Teuff et al. (1999) Anicich (1993) Schunk and Nagy (2000) Schunk and Nagy (2000) Schunk and Nagy (2000) Schunk and Nagy (2000) Li et al. (1997) Le-Teuff et al. (1999) Le-Teuff et al. (1999) Le-Teuff et al. (1999) Scott et al. (1997) Le-Teuff et al. (1999) Anicich and Huntress (1986) Anicich and Huntress (1986) Anicich (1993) Le-Teuff et al. (1999) Anicich (1993) Le-Teuff et al. (1999) Anicich (1993) Anicich (1993) Jones et al. (1981) Le-Teuff et al. (1999) Fox (2015) Jones et al. (1981) Fox (2015) Fox (2015) Fox (2015) Le-Teuff et al. (1999) Prasad and Huntress (1980) Anicich and Huntress (1986)) Schunk and Nagy (2000)

O+2 + e− → O + O

2.4 × 10−07 × ( 300 )0.7 𝑇

Schunk and Nagy (2000)

NO+ + e− → N + O

4.0 × 10−07 × ( 300 )0.5 𝑇

Schunk and Nagy (2000)

N+2 + e− → N + N

2.2 × 10−07 × ( 300 )0.39 𝑇

Schunk and Nagy (2000)

O+ + e− → O

3.26 × 10−12 × ( 300 )0.7 𝑇

Schunk and Nagy (2000)

CO+ + e− → C + O

2.75 × 10−07 × ( 300 )0.5 𝑇

Schunk and Nagy (2000)

N2 H+ + e− → N2 + H

)0.84 2.325 × 10−07 ×( 300 𝑇

Fox (2015)

OH+ + e− → O + H

3.75 × 10−08 ×( 300 )0.5 𝑇

Matta et al. (2013)

2.00 × 10−07 ( 300 ) 𝑇𝑒 3.00 × 10−07 ×( 300 )0.5 ) 𝑇𝑒 0.5 3.4 × 10−07 ( 300 ) 𝑇𝑒

Fox (2015)

𝑒 𝑒 𝑒 𝑒

𝑒 𝑒

𝑒

𝑒

+



HCO + e → H + CO HNO+ + e− → H + NO OCOH+ + e− → CO2 + H

mixture of gases. It is obtained as: ∑ 𝑈𝑐 (𝐸, 𝐸0 ) = 𝑓𝑙 𝑈𝑙 (𝐸, 𝐸0 )

Fox (2015)

Fig. 4 shows the photoelectron fluxes calculated using the model for different orbits. We validate our modeled electron fluxes by comparing with the observations of Solar Wind Electron Analyzer (SWEA) instrument onboard MAVEN, available on PDS (level 2, version 4, revision 1). We use measurements made during the inbound leg of various orbits for comparing with our calculations. SWEA measures the electron fluxes in the 3 eV–5 keV energy range (Mitchell et al., 2016). The modeled spectra clearly show peaks in the energy range 21–27 eV which is caused by the ionization of CO2 by intense He II 30.4 nm radiation (Mantas and Hanson, 1979). Such fine structures are not seen in the measured flux which may be because of the low energy resolution of the SWEA instrument (𝛥E/E = 17%). The best agreement between the model and the observations take place in the energy range 10–60 eV. Both the modeled and the observed flux shows a depletion in the steady state photoelectron spectrum near 3 eV. This is caused by the large vibrational excitation cross section of CO2 at

(7)

𝑙

where 𝜌𝑛 𝑓 𝑙 = ∑𝑙 𝑙 𝑙 𝑖=1 𝜌𝑖 𝑛𝑖

Fox (2015)

(8)

Here U𝑙 (E, E0 ) is the yield spectrum for individual gases, n𝑙 is the number density of 𝑙th gas, and 𝜌𝑙 is the average value of the total inelastic cross section of the 𝑙th gas between E and E0 . The yield spectrum U(E, E0 ) represents the equilibrium number of electrons per unit energy at an energy E resulting from the local degradation of an incident electron of energy E0 . For CO2 , we used the AYS developed by Bhardwaj and Jain (2009). For CO, N2 , and O, the yield spectrum as given by Singhal et al. (1980) is used. 5

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Fig. 4. Modeled photoelectron (red dashed curves) flux compared with the measurements of SWEA (black solid curves) for different orbits. At energies greater than 60 eV, the modeled flux is a factor of 5 lower than the observations. The green dashed dotted curves indicate the photoelectron flux obtained when the input solar flux at wavelengths less than 17 nm is increased by a factor of 5 so that there is a better agreement with the observed photoelectron flux at higher energies.

this energy which makes the primary photoelectrons to perform more vibrational excitation collisions with CO2 , thus reducing the steady state flux. However, the magnitude of the modeled flux seems to be a factor of 5 higher than the observations at energies less than 10 eV. This is because the AYS of CO2 used for calculating the photoelectron flux is quite approximate at energies less than 10 eV (Bhardwaj and Jain, 2009). Similarly, at energies > 60 eV, the model is underestimating the

photoelectron flux by a factor of 5. The study of Xu et al. (2018), Sakai et al. (2016) also have reported a similar kind of discrepancy between the observed and modeled flux between 60 and 200 eV. The combined uncertainties in the cross sections and solar spectra can contribute towards this disagreement between the modeled photoelectron fluxes and the observations. The exclusion of plasma dynamics also might have contributed partly towards the disagreement. The photoelectrons 6

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Fig. 6. Relative contribution of photon impact and electron impact ion production rates to the total ion production rates at different altitudes. Solid lines shows the ratio of photon impact production rates to the total ion production rates and dashed lines are the ratio of electron impact production rates to the total ion production.

impact can be calculated as: 𝑉𝑖 (𝑍) = 𝑛𝑙 (𝑍)



𝜑(𝑍, 𝐸) 𝜎𝑖𝑙 (𝐸) 𝑑𝐸

(9)

where n𝑙 (Z) is the density of the parent neutral species, 𝜑(Z, E) is the photon/photoelectron flux at the altitude 𝑍 and energy 𝐸, and 𝜎𝑖𝑙 is the photoionization/electron impact ionization cross section for ion species 𝑖 being produced from neutral species 𝑙 at energy 𝐸. Out of the seven neutral gases considered in the study, four gases, viz. CO2 , N2 , CO, and O, are photoionized in our model and results in the generation of primary ions CO+ , CO+ , O+ , C+ , N+ , and N+ . The 2 2 production rate of these ions due to photon and electron impact with the aforesaid gases are calculated. Fig. 5 shows the model calculated photon impact, electron impact, and the total (photon impact + electron impact) ion production rates for different orbits. The DD2 and DD8 orbits had smaller SZAs (∼10◦ and ∼25◦ , respectively, at the periapsis). Hence the ion production rates for these orbits are higher and is of the order of 104 cm−3 s−1 at the periapsis altitude. For DD4 orbits, the inbound orbital segments had SZA ∼78–90◦ . Hence the solar flux received during DD4 orbits were less as compared to DD2 and DD8. The peak ion production rate is of the order of 103 cm−3 s−1 and occurs at higher altitude of ∼135 km. To understand the relative contribution of photon impact and electron impact collisions to the total ion production rates, we calculated the ratio of these two with respect to the total ion production rates (see Fig. 6). For DD2 and DD8 orbits, the ion production rates due to photon impact dominate over electron impact at all altitudes, with photon impact contributing ∼70–80% to the total rates. For DD4 orbits, at 125 km, the photon and electron impact makes equal contribution. Below this altitude the contribution from electron impact (∼60%) dominates over photon impact. Dominance of electron impact production rates over photon impact would be happening at lower altitudes for DD2 and DD8 orbits. Due to smaller SZAs during these campaigns, high energy photons would be able to penetrate deeper into the atmosphere thus generating more photoelectrons in the lower altitudes. As discussed in Section 4.1, the model underestimates the flux of photoelectrons having energy >60 eV, by a factor of 5. To understand the impact of this underestimation in the calculation of ion production we increased the modeled photoelectron flux at energies >60 eV by a factor of 5 and used these enhanced photoelectron flux in the calculation of ion production rates. The maximum impact of these enhancement was seen on the DD4 orbits for which low altitude (<130 km) data is available. In this altitude region, the ion production rates

Fig. 5. Calculated ion production rates due to photon impact (dashed) and electron impact (dotted) for different orbits. The solid curves show the total ion production rates which is the sum of the ion production rates due to photon and electron impact.

in the energy range 60 to 200 eV is produced by photons having wavelength <17 nm. Increasing the solar flux in the wavelengths <17 nm by a factor of 5 makes the model calculations consistent with the observations at energies > 60 eV (see Fig. 4). 4.2. Volume ionization rates The gases in the atmosphere are primarily ionized by photon and photoelectron impact ionization of the neutral molecules. The production rate of the primary ions through photon impact and photoelectron 7

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Fig. 7. Modeled profiles of CO+2 (red lines) and O+2 (blue lines) compared with the NGIMS observations (Solid symbols). The modeled CO+2 (O+2 ) profiles are scaled by a factor of 2.5 (1.5) to make them in agreement with the observations.

due to electron impact increased by a factor of 2 to 4. This indicates

modeled photoelectron flux for our subsequent calculations of plasma density profiles.

that these high energy photoelectrons mainly deposit their energy in the low altitude region by causing ionization of the neutrals. At all

4.3. Plasma density profiles

other altitudes (>130 km), the enhancement in the electron impact ion production rates was within 10%–15%, for all orbits. This shows that

The primary ion production rates, discussed in the previous section, are used as input to a photochemical equilibrium model which simulates the ion-neutral chemistry on the dayside Martian ionosphere. Eq. (2) is solved for each ion species at each altitude independently

the underestimation in the high energy photoelectron flux would have a significant impact on the calculation of plasma density profiles only at altitudes less than 130 km. Therefore, we continue to use our actual 8

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by assuming a constant solar zenith angle within the inbound leg of an orbital pass, whereas in reality the spacecraft solar zenith angle changes with altitude. However, the difference between the SZA value used in the model and the actual SZA value at the respective altitudes is not more than 10 degrees for all the orbits considered in the study. Hence assuming a constant SZA will not cause more than 5% error in the modeled densities. The most abundant ion in the Martian ionosphere is O+ , followed 2 by CO+ . Studies suggest that, at regions near the ionospheric peak, 2 they together constitute ∼95% of the total plasma density (Vogt et al., 2017). Fig. 7 shows the modeled profiles of CO+ and O+ compared with 2 2 the NGIMS observations for different orbits. In general, the modeled ion densities exhibit similar altitude variation as the observations. However, it is seen that the modeled CO+ density is a factor of 2.5 2 larger than the observation. Similar is the case with O+ profiles which 2 are roughly 1.5 times of what NGIMS measured. The overestimation in these modeled ion profiles ultimately causes a similar variation in the electron density, when compared with LPW observations. The model derived electron density has to be scaled down by a factor of 1.5–2 to make the profile consistent with the LPW measured electron densities (see Fig. 8). However, it may be noted that qualitatively, the modeled ion densities exhibit similar altitude variation as the observations, which shows that the photochemistry is correctly represented in the model. , and the electron profiles to have quan, O+ For the modeled CO+ 2 2 titatively better consistency with the MAVEN observations, we tried tweaking the model input parameters. Varying the NGIMS-CO2 profile and LPW-T𝑒 are found to be the most effective ways to make the model results match with the observations. Ergun et al. (2015) reports that electron temperature reported by the LPW instrument yields a higher value of Te than the actual value for temperatures less than 750 K (< 180 km). The theoretical Te reported by Peterson et al. (2018) also suggests a lower value than what LPW has reported for altitudes <180 km. However, calculating the ion density profile using the emeprical T𝑒 profile of Peterson et al. (2018) improved the density of O+ at 2 lower altitudes (<170 km) only by a few percentage. Our study shows that LPW T𝑒 has to be reduced by a factor of 4 to make the modeled O+ agree well with the MAVEN profiles. However, varying electron 2 temperature did not have much impact on the CO+ profile as the main 2 loss channel for the CO+ ion is not dissociative recombination reaction 2 . Reducing CO2 density but the reaction with atomic oxygen to form O+ 2 by a factor of 4 brings good agreement between the modeled and the observed profiles, both for O+ and CO+ (see Fig. 9). 2 2 This exercise suggests that for the model to quantitatively reproduce the observed O+ and CO+ densities, and hence the electron density, 2 2 a significant reduction in the input parameter values, viz. neutral CO2 density or the electron temperature is required. However, since reducing the T𝑒 could only improve the O+ profiles, whereas tweaking 2 the neutral CO2 profile could improve both the O+ , CO+ (and therefore 2 2 the electron density profile), there could be some overestimation in the neutral CO2 density measured by NGIMS. The study of Xu et al. (2018) which modeled the ionospheric response to the solar flare on 10 September 2017 also reported that the modeled O+ and CO+ are higher compared to the NGIMS obser2 2 vations. They made a comparison between the LPW-electron density profile with observed NGIMS-O+ profile and found that the latter is 2 ∼2 times lesser than the former. Hence, they suggested a factor of ∼2 absolute calibration to the NGIMS O+ ion density. They also showed 2 that LPW electron densities are in good agreement with the modeled values when multiplied by a factor of 1.4, which is consistent with our results. Xu et al. (2018) attribute this to several factors, like overestimated modeled solar irradiance, uncertainties in the LPW measurements, uncertainties in the neutral densities measured by NGIMS, and uncertainties in the cross sections and reaction rates. Our study suggests that by scaling down the neutral CO2 density, the discrepancy between the LPW observations and the sum of modeled CO+ and 2

Fig. 8. Modeled electron density profiles (lines) compared with the LPW observations (Solid symbols). The modeled profiles are scaled by a factor of 1.5 to make them in agreement with the observations.

for calculating the density profiles. We included the chemistry of only those ions that have density >1 cm−3 as per the NGIMS measurements reported by Benna et al. (2015). The density profiles of 13 ions, viz., CO+ , O+ , NO+ , N+ , N+ , CO+ , O+ , C+ , OH+ , OCOH+ , HNO+ , HCO+ , 2 2 2 and N2 H+ , are calculated in the model. The ion density profiles are then added together to get the electron density at each altitude. Each of the calculated ion profile is compared with the NGIMS observations during the inbound leg of the respective orbits to validate the model calculations. The model calculated the plasma density profiles 9

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Fig. 9. LPW measured e− (black dots) and NGIMS measured CO+2 (red) and O+2 (violet) compared with modeled profiles obtained by tweaking model input parameters. Dashed (dashed dotted) lines are profiles are obtained by reducing electron temperature (neutral CO2 density) by a factor of 4.

O+ profiles can be reduced. The same is true with the differences 2 between modeled and observed CO+ and O+ profiles. It may also be 2 2 noted that mass spectrometry based neutral density estimates are more prone to uncertainties compared to the ion measurements, owing to the spacecraft background, and therefore, reducing the neutral density profile seems more reasonable.

For the minor ions O+ , C+ , and N+ , our modeled profiles are very much consistent with the NGIMS level 2 profiles for amu 16+ , amu 12+ , and amu 14+ , respectively (see Fig. 10). This indicates that the chemistry of these ions considered in the model is sufficient enough to explain their abundance in the Martian ionosphere. In the case of NO+ ion, the model was not able to reproduce the observed density (see 10

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Fig. 11. The modeled profiles (solid lines) for NO+ compared with the NGIMS profiles (dots) for amu 30+ . The dashed dotted line shows the NO+ profile obtained when NGIMS measured NO density was reduced by 1 one order of magnitude and dotted line is the profile for the case when the average NO profile from Fox and Bakalian (2001) is used. See text for more details.

Fig. 10. The modeled profiles (solid lines) for O+ (red), C+ (black), and N+ (blue) compared with the NGIMS profiles (symbols) for amu 16+ (red), amu 12+ (black), and amu 14+ (blue), respectively.

Fig. 11). For all orbits, the model overestimates the NO+ density by an order of magnitude. NO+ is produced in the Martian ionosphere mainly through the reaction of O+ with NO. This reaction constitutes ∼75–95% 2

Fig. 12. Neutral NO density profiles used in the model. 11

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of the total NO+ production rate. The only possible loss channel for this ion is the dissociative recombination reaction for which almost all available measurements suggest a rate coefficient of 4 × 10−7 cm3 s−1 (see Table 7 of Sheehan and St.-Maurice (2004)). It is also seen that the overestimation in the modeled O+ by a factor of 1.5 is not sufficient to 2 explain the overestimation in the NO+ density. Therefore, the disparity between the modeled and the observed NO+ profile could largely be due to neutral NO density used in the model. As mentioned in Section 2, we have used the NGIMS measured NO profile for orbits 1064 and 1835 as given by Mahaffy et al. (2015), Vogt et al. (2017), which are earlier versions of the neutral data and the measured NO profiles are currently not available in PDS (including data for the DD orbits studied here). There are also indications that NO profiles from NGIMS could be erroneous as the NO signal measured by the mass spectrometer could be largely from the recombination of O and N in the instrument itself (Paul Mahaffy, personal communication). As a case study, we made test run of the model by reducing the NO density by one order of magnitude. In this case we could reduce the difference between the modeled and the observed NO+ density from one order to magnitude to a few factors (see Fig. 11). We also tried using an average of the NO profiles given by Fox and Bakalian (2001) for low and high solar activity conditions which were based on the Viking measurements and Mars Thermosphere General Circulation model (see Fig. 12). In this case also we could largely reduce the difference between the modeled and the observed NO+ density. Thus, our study suggests that to reproduce the NGIMS observed NO+ , the neutral NO density should be having a magnitude between 106 –104 cm−3 in the altitude region 120–200 km. Both the ions N2 H+ and HCO+ have the same mass of 29 amu. Hence, both of them contribute towards amu 29+ profile reported by NGIMS. HCO+ is predicted to be present in the Mars ionosphere relatively recently (Matta et al., 2013). Our model results show that the density of HCO+ ion is about 2 to 3 orders of magnitude larger than N2 H+ (see Fig. 13). This suggests that the major contributor towards amu 29+ in the Martian ionosphere is HCO+ rather than N2 H+ . In case of N+ and CO+ ions, both having mass amu 28, the modeled amu 28+ 2 profile (i.e. the sum of N+ and CO+ profiles) is about a factor of 3 larger 2 than the observations. (see Fig. 14) Similar to NO+ , the current model was not able to reproduce the observed HNO+ density. To the best of our knowledge the model of Fox (2015) is the only model to date which calculated the density of HNO+ in the Martian ionosphere and they predicted a very low abundance (<0.03 cm−3 ) for the ion. However, NGIMS measurements reported a substantially high density, ∼400 cm−3 at 150 km (Benna et al., 2015). Following the same chemistry scheme of Fox (2015), we also incorporated the chemistry of HNO+ in our model and compared the profile with NGIMS measurements for amu 31+ for the respective orbits. Our model also predicts a very low density for the ion, ≤10−2 cm−3 , at different altitudes. As shown in Fig. 15, the NGIMS measured value is ∼6–7 orders of magnitude larger than the modeled profile. As per the present chemistry scheme, the major production process of HNO+ is the reaction of N2 H+ with NO, both of which are minor components in the Martian ionosphere. Increasing the NO density even further by 2 or 3 orders of magnitude also will not increase the HNO+ density to the required magnitudes. On the other hand, this will further worsen the inconsistency between the modeled and the observed NO+ profiles as discussed earlier. Inclusion of some more reactions involving abundant neutrals and ions such as CO2 , N2 , O+ , or N+ may help 2 2 to resolve the inconsistency, or it could also be that the NGIMS ion measurements overestimates the HNO+ . For the ion OCOH+ , the modeled profile do not match well with the NGIMS profile for amu 45+ (see Fig. 16). At lower altitudes (<160 km), the difference is about one order of magnitude and it reduces to a few factors at higher altitudes. For the DD4 orbits the structure of the modeled profiles show a closer resemblance to the observations as compared to other orbits. OCOH+ is mainly produced through the reaction of CO+ with H2 . Thus the structure and magnitude of the neutral 2

Fig. 13. The modeled profiles for N2 H+ (pink solid line) and HCO+ (blue dashed lines) compared with the NGIMS profiles (red dots) for amu 29+ .

H2 profile would have a strong impact on the modeled ion profile. We compared the modeled H2 profile of Krasnopolsky (2002) which we are using in our model with that of two more available modeled H2 profile for dayside Martian ionosphere, viz. Fox and Bakalian (2001), Rodrigoe et al. (1990) (Fig. 17). All the three models predicts a similar structure for the H2 profile. Hence using each of these profiles for calculating the OCOH+ density could change only the magnitude of the ion profile and not the structure. Using the H2 profile from Fox and Bakalian (2001) 12

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Fig. 14. The modeled profiles for N+2 (blue dashed dotted line) and CO+ (pink dashed lines) compared with the NGIMS profiles (dots) for amu 28+ . The solid black line shows the sum of the modeled N+2 and CO+ profiles.

Fig. 15. The modeled profiles (solid lines) for HNO+ compared with the NGIMS profiles (dots) for amu 31+ .

A comparison of the modeled OH+ profile with the NGIMS amu-17+ profiles shows a good match with the observations only at altitudes above ∼150 km. Below this height, the model suggest roughly an exponential decrease in density whereas observations show that density remains more or less the same (see Fig. 18). We could not determine a reason for this inconsistency at lower heights. However, a recent study by Halekas et al. (2015) on the detections of lunar exospheric ions by

increased the ion density by ∼one order of magnitude at all heights. These helped to reduce the difference between the modeled and the observed profiles at altitudes <160 km, but increased the difference at higher heights. The fact that the modeled OCOH+ and CO+ exhibit discrepancies also indicate towards the uncertainties in the input CO2 neutral density profile. 13

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Fig. 17. Modeled H2 profiles available for dayside Martian ionosphere. The profile from Fox and Bakalian (2001) is the average of the H2 profile for low and high solar activity conditions.

cannot explain the agreement between the modeled and the observed values at high altitudes. 5. Summary A photochemical model is developed to compare the MAVEN observed ion density profiles during the deep dip campaigns with the modeled profiles on the dayside Martian ionosphere. The model simulated the conditions that prevailed on Mars during the deep dip campaigns of MAVEN. Two orbits each from three dayside deep dip campaigns viz. DD2, DD4, and DD8 are chosen and the MAVEN measured neutral densities, solar flux and electron temperatures for those respective orbits are used as input to the model. The model computed photoelectron fluxes are compared with the MAVEN-SWEA observations and a good agreement was seen in the energy range 10 eV–60 eV. At energies greater than 60 eV, the input solar UV flux at wavelength <17 nm has to be increased by a factor of 5 to have a better match between the modeled and the observed photoelectron flux. In case of plasma density profiles, the model could correctly re, and , CO+ produce the similar altitude variation as the observed O+ 2 2 the electron density. However, the magnitude of the modeled profiles , seems to be a factor of ∼1.5, 2.5, and ∼1.5 larger than the NGIMS-O+ 2 NGIMS-CO+ , and LPW-electron densities, respectively. To understand 2 this, the model inputs were tweaked and the sensitivity analysis showed that reducing the NGIMS-measured CO2 density or the LPW-measured electron temperature by a factor of 4 can bring good match between the modeled and the observed electron density profiles. Since reducing the NGIMS-measured CO2 density could reproduce the observed O+ , 2 CO+ (and therefore the electron density profile), this seems to be the 2 most viable option. This indicates an overestimation in the neutral CO2 density measured by NGIMS. However, the model could very well reproduce the observed profiles of the ions O+ , C+ , and N+ using the actual NGIMS-CO2 density without applying any scaling factor. Therefore, the simulations cannot conclusively confirm that the neutral CO2 densities are exactly off by a factor of four, but only indicate a possibility of overestimation in the NGIMS-CO2 density. The model overestimates the NO+ by one order of magnitude when the NO density profile reported using the previous version of the NGIMS data is used for the calculation. Using the available modeled profiles of NO as input could reduce the difference between the modeled and the observed NO+ densities. Similar is the case with HNO+ where model predicted densities were much lesser than the observations. The study suggests that either with the currently existing chemistry scheme, the models are not able to reproduce the observed HNO+ density, or the

Fig. 16. The modeled profiles (solid lines) for OCOH+ compared with the NGIMS profiles (dots) for amu 45+ . The dashed lines show the OCOH+ profile obtained when the neutral H2 profile from Fox and Bakalian (2001) is used as model input. See text for more details.

using a neutral mass spectrometer suggest that there can be issues with the measurements of water group ions using a mass spectrometer. They observed an unusually high count rate for water group mass channels amu 17+ and amu 18+ . They suspected that the source of these ions could be either solar wind and/or solar wind charge exchange with the neutral water molecules degassed from the spacecraft. Thus, in case of the Martian ionosphere also, there is a possibility that high density of amu 17+ reported by NGIMS could have contributions due to the degassed water molecules from the spacecraft. But this conjuncture 14

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Acknowledgments The work at Space Physics Laboratory is supported by the Indian Space Research Organisation. The work at Physical Research Laboratory is supported by Department of Space, Government of India. The MAVEN data used in this work is taken from the Planetary Data System (https://pds.nasa.gov/). We gratefully acknowledge the MAVEN team for the data. Authors would like to thank the two anonymous reviewers for their constructive comments and suggestions which helped in the improvement of the manuscript. References Andersson, L., Ergun, R.E., Delory, G.T., Eriksson, A., Westfall, J., Reed, H., McCauly, J., Summers, D., Meyers, D., 2015. The langmuir probe and waves (LPW) instrument for MAVEN. Space Sci. Rev. 195 (1), 173–198. http://dx.doi.org/10.1007/s11214015-0194-3. Anicich, V.G., 1993. A survey of bimolecular ion-molecule reactions for use in modeling the chemistry of planetary atmospheres, cometary comae, and interstellar clouds. Astrophys. J. Suppl. Ser. 84, 215–315. http://dx.doi.org/10.1086/191752. Anicich, V.G., Huntress, W.T., 1986. A survey of bimolecular ion-molecule reactions for use in modeling the chemistry of planetary atmospheres, cometary comae, and interstellar clouds. Astrophys. J. Suppl. Ser. 62, 553–672. Benna, M., Mahaffy, P.R., Grebowsky, J.M., Fox, J.L., Yelle, R.V., Jakosky, B.M., 2015. First measurements of composition and dynamics of the Martian ionosphere by MAVEN’s neutral gas and ion mass spectrometer. Geophys. Res. Lett. 42, 8958–8965. http://dx.doi.org/10.1002/2015GL06614. Bhardwaj, A., Jain, S.K., 2009. Monte Carlo model of electron energy degradation in a CO2 atmosphere. J. Geophys. Res. 114, A11309. http://dx.doi.org/10.1029/ 2009JA014298. Bhardwaj, A., Micheal, M., 1999a. Monte Carlo model for electron degradation in SO2 gas: Cross sections, yield spectra, and efficiencies. J. Geophys. Res. Space Phys. 104 (A11), 24713–24728. http://dx.doi.org/10.1029/1999JA900283. Bhardwaj, A., Micheal, M., 1999b. On the excitation of IO’s atmosphere by the photoelectrons: Application of the analytical yield spectral model of SO2 . Geophys. Res. Lett. 26, 393. http://dx.doi.org/10.1029/1998GL900320. Bougher, S.W., Engel, S., Hinson, D.P., Murphy, J.R., 2004. MGS radio science electron density profiles: Interannual variability and implications for the Martian neutral atmosphere. J. Geophys. Res. Planets 109 (E3), http://dx.doi.org/10.1029/ 2003JE002154. Chamberlin, P.C., Woods, T.N., Eparvier, F.G., 2007. Flare Irradiance Spectral Model (FISM): Daily component algorithms and results. Space Weather 5, S07005. http: //dx.doi.org/10.1029/2007SW000316. Chamberlin, P.C., Woods, T.N., Eparvier, F.G., 2008. Flare Irradiance Spectral Model (FISM): Flare component algorithms and results. Space Weather 6, S05001. http: //dx.doi.org/10.1029/2007SW000372. Eparvier, F.G., Chamberlin, P.C., Woods, T.N., Thiemann, E.M.B., 2015. The solar extreme ultraviolet monitor for MAVEN. Space Sci. Rev. 1–9. http://dx.doi.org/ 10.1007/s11214-015-0195-2. Ergun, R.E., Morooka, M.W., Andersson, L.A., Fowler, C.M., Delory, G.T., Andrews, D.J., Eriksson, A.I., McEnulty, T., Jakosky, B.M., 2015. Dayside electron temperature and density profiles at Mars: First results from the MAVEN langmuir probe and waves instrument. Geophys. Res. Lett. 42 (21), 8846–8853. http://dx.doi.org/10.1002/ 2015GL065280. Fennelly, J.A., Torr, D.G., 1992. Photoionization and photoabsorption cross sections of O, N2 , 02 , and N for aeronomic calculations. At. Data Nucl. Data Tables 51 (2), 321–363. Fox, J.L., 2015. The chemistry of protonated species in the Martian ionosphere. Icarus 252 (2015), 366–392. http://dx.doi.org/10.1016/j.icarus.2015.01.010. Fox, J.L., Bakalian, F.M., 2001. Photochemical escape of atomic carbon from mars. J. Geophys. Res. 106, 28785–28796. http://dx.doi.org/10.1029/2001JA000108. Fox, J.L., Benna, M., Mahaffy, P.R., Jakosky, B.M., 2015. Water and water ions in the Martian thermosphere/ionosphere. Geophys. Res. Lett. 42 (21), 8977–8985. http://dx.doi.org/10.1002/2015GL065465. Fox, J.L., Dalgarno, A., 1979. Ionization, luminosity, and heating of the upper atmosphere of Mars. J. Geophys. Res. Space Phys. 84 (A12), 7315–7333. http: //dx.doi.org/10.1029/JA084iA12p07315. Haider, S.A., Mahajan, K.K., Kallio, E., 2011. Mars ionosphere: A review of experimental results and modeling studies. Rev. Geophys. 49, RG4001. http://dx.doi.org/10. 1029/2011RG000357. Halekas, J.S., Benna, M., Mahaffy, P.R., Elphic, R.C., Poppe, A.R., Delory, G.T., 2015. Detections of lunar exospheric ions by the LADEE neutral mass spectrometer. Geophys. Res. Lett. 42 (13), 5162–5169. http://dx.doi.org/10.1002/ 2015GL064746. Hanson, W.B., Sanatani, S., Zuccaro, D.R., 1977. The Martian ionosphere as observed by the Viking retarding potential analyzers. J. Geophys. Res. 82 (28), 4351–4363. http://dx.doi.org/10.1029/JS082i028p04351.

Fig. 18. Modeled profiles of OH+ (solid lines) compared the NGIMS amu-17+ profiles (symbols).

NGIMS could be overestimating the HNO+ . The OCOH+ density profile is found to be highly sensitive to the structure of the neutral H2 profile used as input to the model. For the water group ion OH+ , the model calculations are consistent with the observations at higher altitudes (>150 km). This study is the first attempt to compare the minor ion concentrations observed during the MAVEN DD campaigns with the outputs of a photochemical model. 15

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