Oxygen and carbon isotope ratios in the martian atmosphere

Icarus 192 (2007) 396–403 www.elsevier.com/locate/icarus

Oxygen and carbon isotope ratios in the martian atmosphere Vladimir A. Krasnopolsky a,∗ , Jean Pierre Maillard b , Tobias C. Owen c , Robert A. Toth d , Michael D. Smith e a Department of Physics, Catholic University of America, Washington, DC 20064, USA b Institute d’Astrophysique de Paris, CNRS, 75014 Paris, France c Institute for Astronomy, University of Hawaii, Honolulu, HI 96822, USA d Earth and Space Science Division, Jet Propulsion Laboratory, Pasadena, CA 91109, USA e NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA

Received 25 June 2007; revised 9 August 2007 Available online 12 September 2007

Abstract Oxygen and carbon isotope ratios in the martian CO2 are key values to study evolution of volatiles on Mars. The major problems in spectroscopic determinations of these ratios on Mars are uncertainties associated with: (1) equivalent widths of the observed absorption lines, (2) line strengths in spectroscopic databases, and (3) thermal structure of the martian atmosphere during the observation. We have made special efforts to reduce all these uncertainties. We observed Mars using the Fourier Transform Spectrometer at the Canada–France–Hawaii Telescope. While the oxygen and carbon isotope ratios on Mars were byproducts in the previous observations, our observation was specifically aimed at these isotope ratios. We covered a range of 6022 to 6308 cm−1 with the highest resolving power of ν/δν = 3.5 × 105 and a signal-to-noise ratio of 180 in the middle of the spectrum. The chosen spectral range involves 475 lines of the main isotope, 184 lines of 13 CO2 , 181 lines of CO18 O, and 119 lines of CO17 O. (Lines with strengths exceeding 10−27 cm at 218 K are considered here.) Due to the high spectral resolution, most of the lines are not blended. Uncertainties of retrieved isotope abundances are in inverse proportion to resolving power, signal-to-noise ratio, and square root of the number of lines. Laboratory studies of the CO2 isotope spectra in the range of our observation achieved an accuracy of ∼1% in the line strengths. Detailed observations of temperature profiles using MGS/TES and data on temperature variations with local time from two GCMs are used to simulate each absorption line at various heights in each part of the instrument field of view and then sum up the results. Thermal radiation of Mars’ surface and atmosphere is negligible in the chosen spectral range, and this reduces errors associated with uncertainties in the thermal structure on Mars. Using a combination of all these factors, the highest accuracy has been achieved in measuring the CO2 isotope ratios: 13 C/12 C = 0.978 ± 0.020 and 18 O/16 O = 1.018 ± 0.018 times the terrestrial standards. Heavy isotopes in the atmosphere are enriched by nonthermal escape and sputtering, and depleted by fractionation with solid-phase reservoirs. The retrieved ratios show that isotope fractionation between CO2 and oxygen and carbon reservoirs in the solid phase is almost balanced by nonthermal escape and sputtering of O and C from Mars. © 2007 Elsevier Inc. All rights reserved. Keywords: Mars; Mars, atmosphere; Atmospheres, evolution; Spectroscopy

1. Introduction Isotope ratios are the important clues to study the evolution of volatiles on Mars. Isotope anomalies on Mars were observed for the first time using mass spectrometers onboard the Viking landing probes (Owen et al., 1977). The measured * Corresponding author. Address for correspondence: Department of Physics, 6100 Westchester Park Drive #911, College Park, MD 20740, USA. E-mail address: [email protected] (V.A. Krasnopolsky).

0019-1035/$ – see front matter © 2007 Elsevier Inc. All rights reserved. doi:10.1016/j.icarus.2007.08.013

ratios were 15 N/14 N = 1.7, 40 Ar/36 Ar = 10.3, and 129 Xe/132 Xe = 2.6, all times the appropriate terrestrial ratios. Later a significant enrichment in heavy isotope was found in D/H in the martian water using the ground-based high-resolution spectroscopy (Owen et al., 1988; Bjoraker et al., 1989; Krasnopolsky et al., 1997): D/H ≈ 5.5 times that in the Standard Mean Ocean Water (SMOW), 1.56 × 10−4 . Detections of atomic deuterium using the Hubble Space Telescope (Krasnopolsky et al., 1998) and molecular hydrogen using the Far Ultraviolet Spectroscopic Explorer (Krasnopolsky and Feldman, 2001) made it possible to

Oxygen and carbon isotope ratios on Mars

establish D/H ≈ 2.3 in H2 on Mars (Krasnopolsky, 2002). Loss and isotope fractionation of hydrogen occur by formation of H2 from H2 O in the lower atmosphere with subsequent delivery and decomposition of H2 in the upper atmosphere. Therefore, both the D/H ratios in H2 O and H2 are key values in evolution of water on Mars. The subject of our study is the isotope ratios in CO2 which is the major atmospheric species on Mars. The most restrictive measurements of 18 O/16 O and 13 C/12 C in CO2 were made by the Viking mass spectrometers, and the measured values were within 5% of the terrestrial standards (Owen et al., 1977). Ground-based high-resolution spectroscopy gave 18 O/16 O = 0.96 ± 0.12 and 13 C/12 C = 0.927 ± 0.058 (Schrey et al., 1986), 18 O/17 O = 0.914 ± 0.04 and 13 C/12 C = 0.94 ± 0.15 (Krasnopolsky et al., 1996), and 18 O/17 O = 1.03 ± 0.09 and 13 C/12 C = 1.00 ± 0.11 (Encrenaz et al., 2005). We will discuss these values below. The standard terrestrial isotope ratios are 18 O/16 O = 2.005 × −3 10 , 17 O/16 O = 3.73 × 10−4 , and 13 C/12 C = 1.123 × 10−2 . The oxygen ratios are those in SMOW, and the carbon ratio refers to the Pee Dee belemnite (PDB) standard. Hereafter all isotope ratios will be scaled to these standards. The O and C isotope ratios in CO2 on Mars are very close to the terrestrial values. Therefore the highest accuracy becomes the main requirement to and the main difficulty in measurements of these ratios. The smallest uncertainties are achieved in the measurements of gases trapped in rocks of the SNC meteorites that are believed to be ejected from Mars. The O and C isotope ratios in these meteorites are measured with uncertainties close to 0.1%. However, true uncertainties of the measured isotope ratios may be significantly larger (Krasnopolsky et al., 1996) because of: “(1) possible isotope fractionation of trapped atmospheric gases, (2) contribution of the impactor that ejected the meteoroids from the martian surface, (3) isotope exchange with the terrestrial environment, (4) “during pyrolysis of SNC meteorites, CO2 and other gases may experience considerable isotope exchange with the host rock” (Karlsson et al., 1993), and (5) “the dehydratation process (and other processes which accompany the pyrolysis) evidently has a large kinetic isotope effect favoring light isotopes in the evolved water. Thus the δ 18 O of martian water cannot be determined from SNC data” (Karlsson et al., 1992).” With all these facts taken into account, the isotope ratios of O and C measured directly in the martian atmosphere are of great interest. Below we will describe our observations of the CO2 isotopes on Mars and analysis and extraction of the isotope ratios from the observational data. 2. Observations While all previous ground-based measurements of the CO2 isotopes on Mars were by-products in spectroscopic observations of other species, our observations were specifically aimed at the oxygen and carbon isotope ratios on Mars. There are three major sources of error in the isotope observations by means of high-resolution spectroscopy: (1) equivalent widths of the

397

Table 1 Observational log #

Date

Time at start (UT)

Exposure (min)

Longitude

1 2 3 4 5 6 7 8 9

18/09/99 18/09/99 19/09/99 19/09/99 19/09/99 20/09/99 21/09/99 21/09/99 21/09/99

05:53 06:31 05:04 05:49 06:34 04:59 04:25 05:11 05:58

34 34 40 40 40 47 40 40 40

118◦ 109◦ 101◦ 90◦ 79◦ 112◦ 129◦ 118◦ 107◦

109◦ 101◦ 91◦ 80◦ 69◦ 100◦ 120◦ 108◦ 97◦

Note. East longitudes of the field-of-view center are given at start and end of each exposure.

observed absorption lines, (2) line strengths in spectroscopic databases, and (3) thermal structure of the atmosphere during the observations. We have made special efforts to reduce all these sources of uncertainty. We concluded from the previous spectroscopic study of oxygen and carbon isotopes at 3.7 and 8 µm (Krasnopolsky et al., 1996) that the near infrared region, where reflection of the solar light dominates with a negligible thermal emission, is less sensitive to uncertainties of the temperature profile in the atmosphere. A spectrum we would like to get for extraction of the isotope ratios should involve a great number of the CO2 lines for each isotope. For example, if one observed line results in an isotope abundance with an uncertainty of 10%, then 100 similar lines could provide an uncertainty of 1%. We came to conclusion that the best spectral interval to study is 6020–6320 cm−1 (1.58– 1.65 µm). Our task implies strict requirements to an instrument that should have very high spectral resolution, spectral coverage, and signal-to-noise ratio. Our choice was the Fourier Transform Spectrometer (FTS) at the Canada–France–Hawaii Telescope (CFHT). FTS had the advantage of high spectral resolution that exceeded those of the long-slit grating spectrographs by an order of magnitude. Spectral coverage of FTS was equal to a bandpass of a filter that was designed to cover the range of 6020–6320 cm−1 in full. Spectral coverage of the highresolution grating spectrographs is typically smaller by an order of magnitude. We needed a long exposure time to achieve high signal-tonoise ratio. This was done in nine sessions that were made in four nights in September 1999 (Table 1). The total exposure time was equal to 6 h. Mean observing conditions are given in Table 2. Positions of Mars and the instrument field of view during the observations are shown in Fig. 1. A sum of nine interferograms was processed to get a spectrum which was then divided by a white source spectrum to correct for the filter transmission. The measured spectrum was sampled with an interval of 0.00753 cm−1 . To improve the accuracy of our analysis, we reduced this interval by a factor of 4 using a third degree polynomial fitting for four adjacent points. This fitting is equivalent to the four-term Bessel interpolation formula. The final sampling interval is 0.0018833 cm−1 , and

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Table 2 Mean observing conditions Condition

Value

Heliocentric distance Geocentric distance Mars diameter Diameter of field of view Phase angle Heliocentric velocity Geocentric velocity Doppler shift Areocentric longitude of the Sun Mean surface pressure Mean temperature at 3 mbar

1.41 AU 1.29 AU 7.3 5.0 43◦ −1.52 km s−1 9.85 km s−1 −0.202 cm−1 209◦ 5.96 mbar 217.5 K

Notes. Doppler shift is calculated for ν = 6150 cm−1 . Mean surface pressure is from TES data averaged over the field of view. Mean temperature at 3 mbar (half surface pressure) is from TES and GCM (Forget et al., 1999).

Fig. 1. Position of Mars during the observations. Subsolar point and terminator are shown. The instrument field of view is divided by ten equal parts for proper simulation of the martian conditions during the observations.

the spectrum covers a range of 6022–6308 cm−1 . The signal-tonoise ratio is maximal at 180 in the central part of the spectrum. A mean full width at half maximum of the observed CO2 lines is equal to 10.5 points, that is, 0.0198 cm−1 . The expected natural FWHM of the CO2 lines on Mars is (2π)1/2 νD = 0.0096 cm−1 . Here νD = 3.04 × 10−7 ν(T /μ)1/2 is the Doppler line width, T = 218 K is the temperature at a half pressure level (see below), and μ = 44 is the molecular mass. Both the natural line width and the instrument spectral resolution contribute to the observed line width, and the instrument spectral resolution is equal to δν = (0.01982 − 0.00962 )1/2 = 0.0173 cm−1 . This resolution corresponds to resolving power of ν/δν = 350,000. Small parts of the observed spectrum are shown in Fig. 2. Numerous lines of all CO2 isotopes are clearly seen and well resolved in the spectrum. Taking into account the two-way absorption in the martian atmosphere, the telluric CO2 is less

abundant than that on Mars by a factor of 65. Therefore the telluric CO2 lines are much weaker than the martian lines. Furthermore, they are broader because of the high pressure in the Earth’s atmosphere and Doppler-shifted by 105–110 points to the red. The telluric lines do not present a significant problem in our analysis. 3. Atmospheric conditions within the field of view The FTS field of view covers a large part of the martian disk (Fig. 1). Observing and atmospheric conditions vary significantly within the field of view. To reduce error associated with thermal structure of the atmosphere, we will study these variations and then properly average the results. Our field of view is divided in ten equal parts as it is shown in Fig. 1, and we will obtain mean parameters for each part. Cosines of solar zenith angle ϕ and of viewing angle ψ are the basic observing parameters. Contribution of each part is proportional to its brightness which is a product of albedo and cos ϕ in the Lambert approximation. Albedos of various regions on Mars in the near infrared were measured using ground-based observations (McCord and Westphal, 1971), ISM/Phobos (Erard and Calvin, 1997), and OMEGA/MEX (Bibring et al., 2006). Our long-term highresolution spectroscopy of Mars O2 (1 ) dayglow at 1.27 µm and CO at 1.57 µm using CSHELL/IRTF (Krasnopolsky, 2007) is another source of the data. According to Erard and Calvin (1997), Mars albedo at 1.27 µm is equal to that at 1.6 µm in all spectra of the dark and bright regions. We retrieve the albedo values from our observations (Krasnopolsky, 2007) in the continuum; they vary from 0.21 to 0.38. McCord and Westphal (1971) measured the albedo variations from 0.15 to 0.4 at 1.6 µm. We need mean temperature profiles for each part of the field of view for our analysis. These profiles are taken from simultaneous MGS/TES observations (Smith, 2004). The TES temperature profiles measured at nadir and limb extend from the surface to ∼60 km with a step of a quarter of scale height. We averaged the data observed within a few days from our observations at similar latitudes and longitudes for each part of the field of view. However, the TES profiles refer to the local times of 2:00 and 14:00 while local time (LT) varies within the field of view from 7:20 to 13:20. Therefore we apply data of general circulation models (GCM) for temperature corrections for the local time variability. These corrections for LS = 210◦ at 30◦ N and 30◦ S from the GCM by Forget et al. (1999) are shown in Fig. 3. Results of this GCM are presented as the Mars Climate Database at http://www-mars.lmd.jussieu.fr/. For comparison, we also use similar data from the GCM by Hartogh et al. (2005). The corrected temperature profiles have been obtained for all parts of the field of view. The weighted-mean atmospheric pressure at the surface is equal to 5.96 mbar using the TES data. The weighted-mean temperatures at half surface pressure are derived at 217.5 and 211.0 K using the TES data and the temperature corrections from Forget et al. (1999) and Hartogh et al. (2005), respectively. This level corresponds to the Curtis–

Oxygen and carbon isotope ratios on Mars

399

Fig. 2. Parts of the observed spectrum of Mars. Numerous lines of all CO2 isotopes are seen. Broad lines are either telluric or solar. Upper panel: 7% of the total spectrum. Lower panel: 1.2% of the total spectrum.

4. Retrieval of the isotope abundances

Fig. 3. Differences between temperatures at local times of 8:00, 10:00, and 12:00 and those at 14:00. The data are taken from the general circulation model by Forget et al. (1999).

Godson approximation (Chamberlain and Hunten, 1987); pressure and temperature at this level are close to the effective mean values. According to the TES data, the mean dust opacity was 0.19 at 9.3 µm within the field of view during our observations. Dust optical depth in the visible should be greater by a factor of ∼2, and the expected dust opacity at 1.6 µm is ∼0.3.

Extraction of the isotope abundances from the observed spectra requires exact line strengths. A special laboratory study is being made on this subject by one of us (R.A. Toth). One of the goals of this study is to achieve uncertainty of ∼1% for the CO2 isotope line strengths in the spectral range of our observations. This work has been completed for the main isotope, CO18 O, 13 CO2 (Toth et al., 2007a, 2007b) and is in progress for CO17 O. Therefore we do not consider 17 O/16 O in this paper. The isotope ratios of CO2 in the spectroscopic databases and in the files of Toth et al. (2007a, 2007b) are those of SMOW and PDB and not equal to ratios in the terrestrial CO2 . For example, there is fractionation of the oxygen isotopes between liquid water and water vapor and between liquid water and CO2 , and the oxygen isotope ratios in gaseous H2 O and CO2 in the Earth’s atmosphere are not those in SMOW. Mean quantities of CO2 at each part of our field of view are known. Our task is to derive isotope abundance from each observed absorption line. Hereafter we define the isotope abundance as a ratio of the derived isotope quantity to that in the field of view assuming the isotope ratios in SMOW and PDB, respectively. We use the equivalent width technique (Chamberlain and Hunten, 1987; Krasnopolsky, 1986) to obtain absorber abundance from its absorption line. Equivalent width of an absorp-

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Fig. 4. Measurement of a line equivalent width. The continuum intensity (dashed line) is obtained by a second order polynomial fitting to spectral intervals i1–i2 and i3–i4.

tion line is ∞ W= 0

I0ν − Iν dν. I0ν

Here Iν is the spectral intensity in the line and I0ν is the continuum intensity estimated from the spectrum beyond the line (Fig. 4). In this formulation the line equivalent width does not depend on the instrument properties if the line is not contaminated. We approximate the continuum by a second order polynomial fitting to the spectral intervals i1–i2 and i3–i4 (Fig. 4). We use i1 = −70, i2 = −20, i3 = 20, i4 = 70 as a standard case when the interval of ±70 points centered on the line is not contaminated. 102 and 41 points are assigned to the continuum and to the line, respectively, in this case. If the standard interval of ±70 points is contaminated, then i1–i4 are properly changed. Effects of dust opacity on equivalent widths of absorption lines in the martian atmosphere were studied by Hunten et al. (2000) and Fedorova et al. (2004). Using their results for the dust opacity of 0.3 and the mean airmass of 2.45 in our observations, we may neglect the effects of dust absorption and scattering. Absorption of the solar photons by the CO2 molecules in our spectral range results in vibrational excitation of CO2 . The vibrational quanta exchange in the CO2 –CO2 collisions proceeds with a rate coefficient of 10−12 –10−10 cm3 s−1 . The density at the half pressure level (see Section 3) is ∼1017 cm−3 , and the lifetime of the excited vibrational levels is [1017 × (10−12 –10−10 )]−1 = 10−7 –10−5 s. Radiative lifetimes are very much longer even for the strongest lines (∼104 s in our spectral range), and resonant scattering may be neglected in our analysis. Next task is to calculate the expected equivalent width of a chosen line assuming the terrestrial isotope ratio. These calculations for each line are made for 10 parts of the field of view, in ∼30 atmospheric layers, and in 300 wavenumber intervals with a step of 10−3 cm−1 , that is, in ∼105 points for each line. The adopted step is a tenth of the FWHM for the natural line.

Fig. 5. Relative sensitivity of line strength to variation of effective temperature. This temperature is 217.5 K in our observations. The calculated dependence is approximated by a linear function (dashed line).

Absorption in each layer at each wavenumber is calculated using the Voigt formulation. Then the calculated spectrum of the line for each part of the field of view is convolved with the instrument response function and the line equivalent width is obtained. The final equivalent width is a weighted-mean of the equivalent widths for 10 parts of the field of view. This procedure is repeated for the isotope abundance varying from 0.5 to 2.5. The calculated curve of growth for each line is fitted by a fourth degree polynomial and then compared with the observed equivalent width to yield the isotope abundance that corresponds to the observed line. Retrieved isotope abundances are sensitive to errors in the thermal structure of the atmosphere because the relative populations of the rotational levels depend on the local temperature. Line strength as a function of temperature may be given as S(T ) T0 [1 + (1.715 + T0 /640)e−aω/T0 ] 1 − e−aν/T = × S0 T [1 + (1.715 + T /640)e−aω/T ] 1 − e−aν/T0 ×e

−aElow ( T1 − T1 ) 0

.

Here the subscript 0 refers to the room conditions (T0 = 296 K), a = hc/k = 1.43883 cm K, ν is wavenumber in cm−1 , Elow is the lower state energy in cm−1 , and ω is the wavenumber of the bending vibrational mode: ω = 667.3799, 648.4780, 662.3734, and 664.7291 cm−1 for CO2 , 13 CO2 , CO18 O, and CO17 O, respectively. The first ratio in the right hand side is a rather accurate form of the partition function for the temperature range of 150 to 300 K. While the range of ν is small in our spectrum (6022 to 6308 cm−1 ) and ω is almost constant as well, Elow varies from 0 to 1700 cm−1 in our data. Relative sensitivity of line strength to variation of effective temperature as a function of the lower state energy is calculated in Fig. 5. It is calculated for the effective temperature of 217.5 K and may be approximated by a linear function d ln S 1 dS = ≈ −0.0038 + 2.63 × 10−5 Elow . dT S dT

Oxygen and carbon isotope ratios on Mars

Fig. 6. Abundance of the primary CO2 isotope for each of 229 lines. The mean CO2 abundance and the mean lower state energy are given. Dashed line shows a least square fitting to the data.

401

Fig. 8. Abundance of CO18 O for each of 80 lines. The mean abundance and mean lower state energy are also given. Table 3 Isotope abundances and ratios for LT corrections from GCMs by Forget et al. (1999) and Hartogh et al. (2005) LT corrections

Forget et al. (1999)

Hartogh et al. (2005)

Effective temperature Temperature correction CO2

217.5 K −1.0 K 1.2404 ± 0.0124

211.0 K 3.6 K 1.2408 ± 0.0122

13 CO

2 13 C/12 C

1.2154 ± 0.0198

1.2071 ± 0.0195

0.980 ± 0.020

0.973 ± 0.020

CO18 O

1.2639 ± 0.0156

1.2589 ± 0.0154

18 O/16 O

1.019 ± 0.018

1.015 ± 0.018

Notes. C and O isotope ratios are times PDB and SMOW, respectively. Their uncertainties include the uncertainties of line strengths.

Fig. 7. Abundance of 13 CO2 for each of 76 lines. The mean abundance and mean lower state energy are also given.

For example, if an error in the effective temperature is 5 K, then relative errors in the line strengths are 0.94%, 2.25%, and 4.9% for Elow = 500, 1000, and 2000 cm−1 , respectively. These errors in line strengths are directly converted into errors of the isotope abundances because the SN products are used in the Voigt line profile. (N is the column abundance in cm−2 .) A CO2 line with strength of 10−27 cm2 cm−1 (in short, cm) at 217.5 K (Table 2) results in a peak absorption of ∼3%. Our spectrum includes 475 lines of CO2 , 184 lines of 13 CO2 , 181 lines of CO18 O, and 119 lines of CO17 O exceeding 10−27 cm, that is, ∼1000 lines total in the spectrum of ∼150,000 sampling points. We made three independent evaluations of equivalent widths for each line with time intervals of a few days. Then we choose lines with differences in the equivalent widths between three values smaller than 5%. The numbers of these lines are 229 for CO2 , 76 for 13 CO2 , and 80 for CO18 O. The retrieved isotope abundances for these lines are shown in Figs. 6–8. A least-square fitting to the data for the main isotope revealed a weak negative slope. This slope may be compensated

using additional corrections of −1.0 K and 3.6 K to all temperature profiles for the LT corrections from Forget et al. (1999) and Hartogh et al. (2005), respectively. The corrected data in Fig. 6 result in the mean abundance of the main isotope of 1.240 with standard deviation of 0.187. Then uncertainty of the mean value is 0.187/(229 − 1)1/2 = 0.0124. Similarly, the abundance of 13 CO2 is 1.215 ± 0.0198 (Fig. 7) and 13 C/12 C = 0.980 ± 0.0187. The abundance of CO18 O is 1.264 ± 0.0156 (Fig. 8) and 18 O/16 O = 1.019 ± 0.0162. The retrieved isotope abundances exceed one. [Encrenaz et al. (2005) met a similar deviation of the opposite sign.] This may be explained by a shift of the field of view from the position in Fig. 1 to the limb during a significant part of the exposure time. The airmass is large on the limb, and this could increase the retrieved abundances. Fortunately, lines of all isotopes are affected similarly by this effect, and it cancels out in the isotope ratios. Possible random errors in the determination of the line strengths are already involved in the derived isotope ratios. We assume that a systematic error for each band is ∼1%, and these errors are statistically independent. The abundances of CO2 , 13 CO , and CO18 O have been determined using the lines of six, 2 three, and two bands, respectively. This changes only slightly the final uncertainties of the isotope ratios (Table 3).

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Table 4 Comparison of the isotope ratios Reference

13 C/12 C

18 O/16 O

Owen et al. (1977) Schrey et al. (1986) Krasnopolsky et al. (1996) Encrenaz et al. (2005) This work

1.00 ± 0.05 0.927 ± 0.058 0.94 ± 0.15 1.00 ± 0.11 0.978 ± 0.020

1.00 ± 0.05 0.96 ± 0.12 − − 1.018 ± 0.018

Note. C and O isotope ratios are times PDB and SMOW, respectively.

We repeat this retrieval for the LT corrections from the GCM by Hartogh et al. (2005). The results are compared with those with the LT corrections from Forget et al. (1999) in Table 3. The differences are much smaller than the achieved uncertainties. Our final isotope ratios are 13 C/12 C

= 0.978 ± 0.020 times PDB,

18 O/16 O

= 1.018 ± 0.018 times SMOW.

These ratios are weighted-mean with the weights of 1/|−1.0| = 1.0 and 1/3.6 = 0.28. Here −1.0 and 3.6 are the additional temperature corrections in Table 3. 5. Discussion We have made special efforts to acquire the best spectrum for determination of the CO2 isotopes on Mars with the best available instrument for this purpose. The laboratory study of the absorption line strengths for the CO2 isotopes in the spectral range chosen by us has reduced the line strength uncertainty to ∼1%. The simultaneous MGS/TES measurements of temperature profiles and the GCM simulations of thermal structure of the martian atmosphere helped to reduce this type of uncertainties as well. We also applied special means to reduce these uncertainties in the isotope ratios. Therefore we succeed to retrieve the isotope ratios that are more accurate than the ratios in the previous studies. Table 4 gives a summary of all measurements of 13 C/12 C and 18 O/16 O in the martian CO2 . All values agree within their uncertainties. Nonthermal escape and sputtering of O and C result in enrichment of the remaining gas in heavy isotopes. Fractionation between gaseous CO2 and solid reservoirs of O and C favors depletion of heavy isotopes in the gas phase. Evidently this type of fractionation is almost balanced by the escape processes on Mars. The achieved uncertainties of the isotope ratios make them possible to use for history of volatiles on Mars and other geophysical applications with a greater confidence. 6. Conclusions We made spectroscopic observations of Mars to measure the oxygen and carbon isotope ratios in CO2 . The chosen spectral range of 6022–6308 cm−1 includes ∼1000 lines of four main CO2 isotopes with strengths exceeding 10−27 cm at 217.5 K which is the mean temperature at a half pressure level in the instrument field of view. We chose the Fourier Transform Spectrometer at the Canada–France–Hawaii Telescope that had

spectral resolution and spectral coverage better than those of high-resolution grating spectrographs by an order of magnitude. The achieved resolving power and signal-to-noise ratio are 350,000 and 180, respectively, in the measured spectrum. Incidence and observing angles and albedo variations within the field of view are reproduced using the geometry of the observations and the previous mapping observations with IRTF/CSHELL. Temperature profiles within the field of view are taken from the simultaneous measurements by MGS/TES. Corrections to the temperature profiles for variations with local time are taken from the GCM. All these means help to reproduce thermal structure of the atmosphere during the observations. Laboratory studies have been made to measure line strengths of the CO2 isotopes with uncertainty of ∼1% in our spectral range. Equivalent widths of each line were calculated in different parts of the field of view using the Voigt line shape and then averaged to fit the measured equivalent width. Using a combination of all these factors, the highest accuracy has been achieved in measuring the CO2 isotope ratios: 13 C/12 C = 0.978 ± 0.020 and 18 O/16 O = 1.018 ± 0.018. The retrieved ratios show that isotope fractionation between CO2 and oxygen and carbon reservoirs in the solid phase is almost balanced by nonthermal escape and sputtering of O and C on Mars. These isotope ratios may be used for history of volatiles on Mars and other geophysical applications with a greater confidence. Acknowledgments We are grateful to François Forget and Alexander Medvedev for some data on temperature variations with local time from their general circulation models. References Bjoraker, G.L., Mumma, M.J., Larson, H.P., 1989. Isotopic abundance ratios for hydrogen and oxygen in the martian atmosphere. Bull. Am. Astron. Soc. 21, 991. Bibring, J.P., Langevin, Y., Mustard, J.F., Poulet, F., Arvidson, R., Gendrin, A., Gondet, B., Mangold, N., Pinet, P., Forget, F., 2006. Global mineralogical and aqueous Mars history derived from OMEGA/Mars Express data. Science 312, 400–404. Chamberlain, J.W., Hunten, D.M., 1987. Theory of Planetary Atmospheres. Academic Press, San Diego, 481 pp. Encrenaz, Th., and 10 colleagues, 2005. Infrared imaging spectroscopy of Mars: H2 O mapping and determination of CO2 isotopic ratios. Icarus 179, 43–54. Erard, S., Calvin, W., 1997. New composite spectra of Mars, 0.4–5.7 µm. Icarus 130, 449–460. Fedorova, A.A., Rodin, A.V., Baklanova, I.V., 2004. MAWD observations revisited: Seasonal behavior of water vapor in the martian atmosphere. Icarus 171, 54–67. Forget, F., Hourdin, F., Fournier, R., Hourdin, C., Talagrand, O., Collins, M., Lewis, S.R., Read, P., Huot, J.P., 1999. Improved general circulation models of the martian atmosphere from the surface and above 80 km. J. Geophys. Res. 104, 24155–24176. Hartogh, P., Medvedev, A.S., Kurola, T., Saito, R., Villanueva, G., Feofilov, A.G., Kutepov, A.A., Berger, U., 2005. Description of a new general circulation model of the martian atmosphere. J. Geophys. Res. 110, doi:10.1029/2005JE002498. E11008.

Oxygen and carbon isotope ratios on Mars

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