New Upper Limits on Numerous Atmospheric Species in the Native Lunar Atmosphere ⇑ Jason C

Icarus 225 (2013) 681–687

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Icarus

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New upper limits on numerous atmospheric species in the native lunar atmosphere ⇑ Jason C. Cook a, , S. Alan Stern a, Paul D. Feldman b, G. Randall Gladstone c, Kurt D. Retherford c, Constantine C.C. Tsang a a Southwest Research Institute, 1050 Walnut St. Suite 300, Boulder, CO 80302, United States b Johns Hopkins University, Department of Physics and Astronomy, Baltimore, MD 21218, United States c Southwest Research Institute, 6220 Culebra Road, San Antonio, TX 78238, United States article info abstract

Article history: We report on our analysis of twilight observations of the lunar atmosphere as observed by the LAMP Received 6 December 2012 instrument aboard NASA’s Lunar Reconnaissance Orbiter. Using data collected between September Revised 18 March 2013 2009 and March 2013, we have examined over 7.7 million s of integration time obtained when the sur- Accepted 11 April 2013 face was in darkness, but the atmosphere between the spacecraft and the surface was in sunlight. Using Available online 30 April 2013 these data, we have calculated upper limits for 27 species, primarily neutral atomic and molecular species, but also a few atomic ions. All of these species have either been predicted previously or were Keywords: observed by LAMP after the LCROSS impact. Our upper limits are more constraining than previous upper Atmospheres, Composition limits by signiﬁcant factors, providing new constraints on numerous species. Puzzlingly, we did not Abundances, Atmospheres Moon detect the previously detected noble gas species Ne and Ar, a subject we brieﬂy touch on here and plan Spectroscopy a full paper about. Ultraviolet observations Ó 2013 Elsevier Inc. All rights reserved.

1. Introduction of Ar was seen to decrease at sunset from 7–8 103 cm3 to <103 cm3 as it condensed through the night and then rapidly ris- The lunar atmosphere, first detected in 1972, is extremely ten- ing to 3.5 104 cm3 before sunrise, attributed to the nearby sunlit uous and classified as a surface boundary exosphere (SBE), with surface being warmed and the resulting liberation of Ar (see Fig. 1 important similarities to Mercury’s atmosphere (Stern, 1999). of Hoffman and Hodges (1975)). The bound species in this atmosphere follow parabolic trajectories The sensitive LAMP (Lyman Alpha Mapping Project) far- as they bounce across the surface, rarely interacting with one and ultraviolet (FUV) spectrograph (Gladstone et al., 2010b) on NASA’s other, because their mean free path is long compared to their scale Lunar Reconnaissance Orbiter (LRO) was designed to map the height. Moon’s permanently shadowed regions but has also been used to The first detection of the atmosphere was made in situ with the search for gaseous species in the lunar atmosphere. Using LAMP, Lunar Atmospheric Composition Experiment (LACE) mass spec- Stern et al. (2012) detected He for the first time by remote sensing. trometer deployed by Apollo 17 (e.g., Hoffman et al., 1973). The They estimated a surface number density of 7 103 cm3. While instrument, which conducted useful operations primarily during this number density may appear low compared to LACE measure- the lunar nights because of count saturation during the lunar day- ments, the observations in Stern et al. (2012) were conducted at time, functioned for nine lunations (December 1972 to October local dusk and are consistent with the findings of Hoffman et al. 1973) before failing. Among the masses detected by LACE during (1973). Further study of He by Feldman et al. (2012) using LAMP that time were masses 4 and 40 amu, which have been attributed has shown that its number density decreases when the Moon to He and Ar, respectively. The surface density of these species was crosses the Earth’s magnetotail, showing that most lunar He observed to vary strongly over the lunar night. In the case of He, originates from the solar wind, as expected. which is too volatile to condense at night, the surface density in- Besides He, LAMP was also used to detect Mg, Ca, Hg, H2 and CO creased from 1 to 2 104 cm3 at dusk to about 5 104 cm3 by in the plume created by the LCROSS impact (Gladstone et al., midnight before decreasing again just before sunrise (see Fig. 2 2010a; Hurley et al., 2012). These species were released in signifi- of Hoffman et al. (1973)). The behavior of Ar, however, was quite cant quantities from the surface because of the heat generated by different from He because Ar is condensable. The surface density the LCROSS impact. In the absence of impacting meteors the size of LCROSS, of which there are several per year, these species are gen- erally not present in the atmosphere, except for a small fraction of ⇑ Corresponding author. E-mail address: [email protected] (J.C. Cook). native H2 (Stern et al., 2013).

Concerning other lunar atmospheric species, although LACE was sensitive to masses between 1 and 110 amu, reports by Hoffman et al. (1973) and Hoffman and Hodges (1975) show that other than He and Ar, no other detected masses showed clear diurnal variations or enhanced pre-dawn concentrations, such as Ar. Hoffman et al. (1973) showed that cool down tests of LACE were capable of almost purging the instrument of possible hydrogenated halo- gens (HF and HCl) which contaminated measurements at mass 20, 36, and 38. Following these tests, they concluded that residual signal at mass 20 was due to Ne. Repeating the cool down tests four times, they gained confidence in a Ne detection, measured with a surface night time density of 7–9 104 cm3. The same reports also showed that there were indeed some indication of a wide variety of species such as CH4,NH3,H2O, N2, CO, O2 and CO2. Additionally, Fastie et al. (1973) used data taken with the Ultraviolet Spectrometer (UVS) on the Apollo 17 command module to search for other species but were only able to derive upper limits for H, H , O, C, N, and CO, which were later revised by Feldman 2 Fig. 1. Nadir observing geometry used in atmospheric exploration at twilight. The and Morrison (1991) after additional analyses. The Alpha Particle dashed red line represents the orbit of LRO. Fluorescence by species present in the Spectrometer experiment, which was carried into orbit by Apollo illuminated atmospheric column near the terminator will be detectable by LAMP. 15 and 16, detected 222Rn and 210Po (Gorenstein et al., 1974). These Because the surface is dark, reflected sunlight is not an issue. These observations decay products of 238U were seen in greater concentration over a occur twice per orbit, with total integration times between 600 and 3600 s per orbit, depending on the b-angle. (For interpretation of the references to color in this few mare and craters. Since the Apollo missions, ground based figure legend, the reader is referred to the web version of this article.) observations in 1988 showed that Na and K are also present in the lunar atmosphere (Potter and Morgan, 1988; Tyler et al., 1988). In this paper we discuss the search for numerous atomic neu- the atmospheric column below LRO is in sunlight. These ‘‘twilight’’ trals (e.g., H, O, Ne), ions (e.g., O+) and molecules (e.g., CO), includ- observations occur twice per 2-h orbit, about 11–12 times per day. ing other species that were detected in the LCROSS plume In a typical orbit, the duration in this twilight is about 600 s (when (Gladstone et al., 2010a), suggested by the mass spectrometer LACE the spacecraft is at low jbj angle, the angle between the orbit plane (Hoffman and Hodges, 1975) or predicted previously (Flynn and and the vector to the Sun), but these periods are much longer Stern, 1996). In an accompanying paper (Stern et al., 2013) we re- around the solstices, reaching over 3600 s per orbit (i.e., when the orbit jbj90°). The total integration time from September port the detection of native H2 by this same search. The structure of this paper is as follows: in Section 2, we discuss our observations 2009 to March 2013 for all LAMP nadir observations in twilight 6 and how they were obtained. In Section 3, we describe how the is about 18.0 10 s. As a result, these twilight observations are data are extracted. In Section 4, we present our upper limit results. far more sensitive to detecting weak emission features than the Finally, in Section 5, we briefly compare our results with previous near surface tangent observations, because the latter has accumu- measurements and discuss related findings. lated about an order of magnitude less observing time.

2. Observations 3. Data reduction

In June 2009, NASA’s LRO entered a polar orbit around the We employ in this paper LAMP twilight observations. Because Moon. As described in detail by Gladstone et al. (2010b), the the twilight atmosphere is illuminated, the species present in the photon-counting LAMP imaging spectrograph is comprised of a atmosphere fluoresce. The brightness of a fluorescing line is a func- telescope and Rowland-circle spectrograph. LAMP has a 40 40 mm2 tion of the column length of illuminated atmosphere, which is the entrance aperture that feeds light to the telescope section of the distance between the top of the lunar shadow and the altitude of instrument. Light entering this aperture is collected and focused LRO. From September 2009 to December 2011 the altitude range by an f/3 off-axis paraboloidal primary mirror at the back end of of LRO was between 30 and 80 km. After December 2011, the orbit LAMP’s telescope section, onto a 0.3° 6.0° spectrograph entrance of LRO was changed to make it more elliptical, with an altitude slit. After passing through the slit, the light falls onto a toroidal range of 30–200 km and apoapsis at the north pole. holographic diffraction grating, which disperses the light onto a From the twilight observation set, we eliminate some of the 2-D (1024 32)-pixel double-delay line (DDL) microchannel plate observations for several reasons. First, we noticed in our analysis detector. This detector is coated by a CsI solar-blind photocathode, of the data, that observations obtained when Solar Particle Events and covers a bandpass of 575–1965 Å. A spectral line that fills the (SPEs) passed the Earth and Moon have greater count rates at all slit has a spectral resolution (FWHM) of 28 Å. Normal LAMP oper- wavelengths compared to data obtained under normal conditions. ations take place in a nadir-staring push-broom scanning fashion The greater count rate, however, does not improve the signal, in- so that the surface is continuously mapped. stead the SPEs adversely affect the signal by being noisier than LAMP observations of the lunar atmosphere were acquired in spectra taken under normal conditions and the detector dark signal one of two ways. In the Stern et al. (2012) He detection paper, measurably increases at these times. We use GOES satellite data to we used observations when LRO is pointed nearly tangent to the determine when proton fluxes triple the background rate of surface. This observing method both eliminated surface albedo 0.5 proton/(cm2 s sr MeV), and remove these data from our atmo- contrast issues with weak atmospheric emission and also greatly spheric search dataset. This eliminates about 0.9 million s of inte- increases the path length observed by LAMP; together these factors gration time, or 5%. make the detection of weak emission lines possible. However, this Second, we exclude data that may be contaminated by sunlight observing mode is rarely used because LRO points toward the sur- reflected from mountain peaks. We assume that when the solar ze- face >95% of the time. Alternatively, Fig. 1 illustrates that there are nith angle is 90°, then the surface is in darkness. This would only be also periods during the orbit when the surface is in darkness, but true if the surface was smooth, which is certainly not the case. We J.C. Cook et al. / Icarus 225 (2013) 681–687 683 noted that including observations that sampled the atmosphere Prior to wavelength correction, we saw that the full range of the below 6 km produces spectra that increase in brightness at longer He I, O I and Lyman-a ghost position was 7 pixels (12.7 Å). After wavelengths, indicative of contamination from sunlight. Of the the correction, the variation is reduced to 2 pixels, or about 3.6 Å. 18.0 million s, another 6.3 million s are contaminated by reflected This variation is similar to that seen with the Lyman-a ghost and sunlight, or 35%. 10% of the 28 Å filled slit. Finally, we do not include observations with a latitude >j75°j.As will be discussed in Section 3.2, we model and remove background 3.2. The removal of reflected celestial, telluric and lunar surface sources. We are only able to accurately model the background backgrounds sources for latitudes where comparable night-time spectra are also obtained. The atmosphere is always seen in sunlight at latitudes All of the LAMP data we analyze contain signals from several >j75°j. We find that interpolating the background signal around sources. The non-atmospheric spectral signals must be accounted the poles was insufficient, and led to a final spectrum that was for because the spectral features we are seeking in the lunar atmo- noisier. By excluding the data poleward of j75°j latitude, we elim- sphere are weak. The sources of background signals include re- inate yet another 3.1 million s of observations. Thus our final spec- flected signals from the IPM (interplanetary medium), UV bright trum is about 43% of the available twilight observations, or stars, Earth shine and changes in the Moon’s reflectivity. Despite 7.7 million s. the low albedo of the Moon (2–6% at these wavelengths; Gladstone et al., 2012), these signals are bright enough for the sensitive LAMP 3.1. Wavelength calibration instrument to detect since these are indeed its primary method for measuring PSR albedos, and thus must be removed. In order to set sensitive upper limits and search for extremely The IPM contributes greatly to the most prominent feature in weak emission lines, we need to be sure that LAMP’s wavelength the LAMP data, i.e., the Lyman-a feature at 1216 Å. LRO’s orbit is calibration is highly accurate; this is because we must carefully re- approximately inertial, with a normal about the ecliptic longitude move scattered light from the wings of the strong Lyman-a feature 102° and latitude 0°. This means that the same region of sky is re- in all of our data. A wavelength calibration was initially derived by flected from the lunar surface for a period of about 6 months. Witte comparing LAMP spectra to stellar reference spectra. Because the (2004) has shown that the Lyman-a sky from the IPM is 2–3 times instrument shows slight wavelength shifts that vary with instru- brighter in the upstream direction near the ecliptic latitude and ment temperature, we had to solve for the wavelength solution fre- longitude +9° and 252°, respectively. This region of sky illuminates quently in the dataset using three fiducial lines, He I (584 Å), O I the lunar surface from about the winter solstice to summer solstice (1304 Å) and a ghost signal of Lyman-a (near 1000 Å), which are al- (hereafter referred to as the ‘‘spring sky’’). In addition to containing ways in our data. the brightest part of the IPM, the spring sky also contains the The Lyman-a ghost is a known artifact of the LAMP instrument, majority of the UV bright stars from the plane of the Milky Way. caused by internal scattering of the light. We do not use the actual These UV bright stars are mainly found in the southern Lyman-a line for wavelength calibration because the detector has hemisphere. experienced charge depletion (a.k.a. gain sag) at that wavelength, To model the contribution by the IPM and UV bright stars to the which has considerably altered the Lyman-a line shape. This would data, we use the observations of the Moon obtained when LRO and make centering on Lyman-a itself difficult. the surface were in darkness and the phase angle between Earth– The examination of our wavelength fiducial lines show that Moon–LRO was >90°. This excludes data contaminated with re- their positions vary with instrument temperature. As the instru- flected day glow from the Earth. We use these data to (i) measure ment approaches equinox, LRO instrument temperature increases, day-to-day variations in the reflected sky spectrum and (ii) map and conversely near solstice. An increase (decrease) in temperature out changes in the spectrum with respect to the ecliptic latitude causes the spectrum to shift to a greater (lesser) pixel index along and longitude of the instrument boresight. Both of these steps the array. Upon subsequent heating and cooling cycles, the spec- are performed on each wavelength bin. To accomplish the first trum does not necessarily return to the initial position. An addi- step, we examine the data confined to a 20° region of ecliptic lati- tional shift was seen in the data when the high voltage setting tude over a 5 day time span in each of the two sky conditions (i.e., was changed in late-May 2011. fall and spring). This allows us to examine the same region on the To characterize these shifts with temperature we measure the Moon under the same illumination conditions. By measuring the position of the He I, O I and Lyman-a ghost lines. The He I line is average of these data against longitude we see that the primary weakly detected only in twilight spectra when the instrument is source of variation is due to lunar albedo variations. Once these warmest (ingress) and coldest (egress) for each orbit. We average variations are divided out, the resulting ratio is a scale of bright- three consecutive days of He observations and temperature mea- ness changes at each wavelength bin over time. The brightness surements at ingress and egress using the exposure time as variations are then interpolated for all times and latitudes, and weights to build signal. The O I line is due to reflected terrestrial the scale is divided out of the night side observations. We then airglow, as evidenced by being detected only within 4 days of map the sky brightness in ecliptic longitude and latitude. From new Moon (i.e., when the Earth illumination is maximized). The this, we can create high SNR models of the sky brightness for a Lyman-a ghost line is strongly detected every orbit. A Gaussian is range of illuminations. Our technique to map the sky to make fit to each of these lines to measure its position. Using the background model spectra is based on observations from longi- measured position of the He I and O I lines only, we fit a first order tudes 90° to 270° and latitudes 75° to 75°, and therefore average polynomial to the line positions vs. instrument temperature for over variations in the reflectivity of the lunar surface. each heating/cooling cycle. Using the fitted values for He I and O The second most important source we see in the data is due to I positions, we estimate the wavelength dispersion (1.81 Å/pix) reflected telluric dayglow. The appearance of emission lines from and obtain the wavelength of the Lyman-a ghost (988.5 ± 3.2 Å). O I (1304 and 1356 Å), N2 LBH bands (e.g., (2,0) 1384 and (1,0) We then refine our estimates for the wavelength solution, by using 1416 Å) and H I (Lyman-a at the ghost wavelength) (Meier, the fitted He I and O I positions, and actual Lyman-a ghost position. 1991) are apparent in spectra obtained on the Earth facing side We update the wavelength solution for each orbit of LRO and of the Moon. These lines are seen to vary with the synodic period interpolate the spectrum to the wavelength solution obtained of the Moon (29.5 days) and with phase angle from the sub-Earth during preflight measurements. point on the Moon. 684 J.C. Cook et al. / Icarus 225 (2013) 681–687

Fig. 2. UV map of the Moon. We ratio the night-side data to the modeled reﬂected IPM/bright UV stars and telluric signals and sum the data from wavelengths 570 to 1900, excluding 1175–1275 Å. The map is overlayed with a visible image of the Moon. This shows that much of the near-side mare appear brighter than the surrounding terrain at LAMP wavelengths, but the correlation does not always hold. There are regions near the north pole which appear bright at these wavelengths, while there are mare that appear near neutral (green) around 10° longitude, and 20° latitude. (For interpretation of the references to color in this ﬁgure legend, the reader is referred to the web version of this article.)

Fig. 3. The average spectrum for the lunar atmosphere as seen by LRO/LAMP. The average spectrum is shown in black, and its 3r error range is shown in gray. Observations obtained two hours before dawn, and two hours after dusk as shown in red and blue, respectively. We indicate the location of He I (584 Å), its 2nd (1168 Å) and 3rd order

(1752 Å) reflections, H2 (1025–1607 Å; Stern et al., 2013), an instrument artifact (783 Å), and the location of the Lyman-a ghost. We indicated the expected location of the Ar and Ne lines. In green, is a model H2 spectrum assuming only resonance fluorescence. The strongest lines, indicated here by the vertical black lines, arrise from Lyman b pumping to the upper J =0,m = 6 state and falling into the lower J = 1 and m states 0–13. (For interpretation of the references to color in this figure legend, the reader is referred to the web version of this article.)

To derive a model of the telluric contribution, we subtract an For illustrative purposes, we show in Fig. 2 our UV map of the IPM and bright UV star spectral model from observations obtained Moon with a superimposed visible wavelength image. This figure at night when the Earth–Moon–LRO angle was <90°. The residual shows that many, but not all, of the mare appear brighter than spectrum contains the remaining telluric signal. As with deriving the average (colored yellow and red in Fig. 2) surface (green). We the IPM and bright UV star spectral model, we use observations note that because there are uncertainties disentangling the IPM, over a variety of surface reflectivity and thus average out their ef- telluric and albedo variations poleward of j75°j, we do not directly fect. At each wavelength bin, we map the count rate to the phase of compare this map with those shown in Gladstone et al. (2012). the Moon and the phase angle between the Earth–Moon–LRO an- However, we can state that our observations are in agreement with gle. As expected, the lines detected are brightest near the sub-Earth Lucke et al. (1974) and Henry et al. (1976) who also concluded that point during new Moon (i.e., full Earth). The brightness of each line the mare are brighter than the highlands at these wavelengths. decreases with distance from the sub-Earth location and as the lunar phase nears full Moon (new Earth). 3.3. Mission averaged spectrum The final important background source is the lunar surface itself. As stated earlier, at these wavelengths, the lunar surface has We show in Fig. 3 our mission averaged spectrum after all of the known albedo variations from 2% to 6% (Gladstone et al., 2012). background signals are removed. Apparent in this spectrum are Our procedure so far has averaged out variations in lunar reflectiv- several features, including the He I line at 584 Å which has been re- ity. These variations (±5%) must be accounted for to avoid over or ported earlier (Stern et al., 2012; Feldman et al., 2012). This He line under estimating the signal from reflected IPM, UV bright stars and also appears as a third order reflection at 1752 Å and likely as a the Earth. In addition, by accounting for the lunar reflectivity, we second order reflection at 1168 Å, near the blocked out Lyman-a can correct any wavelength to wavelength variations which may area, and blended with the (6,3) line of H2. For the first time, we be the case when examining spectra from different terrains (mare detect H2 in the lunar atmosphere as reported in Stern et al. vs. highlands). We estimate the differences in surface reflectivity (2013). This spectrum also shows a known instrumental artifact by taking the ratio between the spectra obtained at night with around 783 Å. The spectrum shows very little evidence for the the summation of the IPM/UV star signal and the telluric dayglow Lyman-a ghost. Other than the afore mentioned features, there is signal. The ratio shows the differences due the surface reflectivity. little evidence for anything else in the spectrum. J.C. Cook et al. / Icarus 225 (2013) 681–687 685

4. Calculation of species upper limits The next step is to convert the column density to surface density so we can compare our results with previous upper limits. To calculate the surface density upper limits we first calculate To calculate the surface density, we apply a Chamberlain Model the upper limit on the brightness of each line. This is measured (see Chapter 7 of Chamberlain and Hunten (1987)). Z from the summation of the average spectrum and its 3r error. z1 The average spectrum is shown in Fig. 3 with the shaded points N ðcm2Þ¼ NðzÞdz; ð2Þ representing the 3r error range. The spectrum is converted from z0 counts/s to Rayleighs/Å using the effective area (Fig. 4) and solid where z0 is the shadow height and z1 is the altitude of the spacecraft angle of the instrument. We integrate the signal at the wavelength and N(z) is defined as of interest over a full width of 28 Å. This is the measured filled slit spectral resolution (Gladstone et al., 2010a). We list the upper lim- NðzÞ¼N ðcm3Þfeðkkc Þ: ð3Þ its on the line brightness in the fourth column of Table 1. We use the measured line-of-sight columns for N (cm2), f is the We next calculate the upper limit on the line-of-sight column partition function and k ¼ GMm where G is the universal constant, number density using, c kTr M is the mass of the Moon (7.36 1025 g), m is the mass of the mol- B ¼ gN; ð1Þ ecule, k is the Boltzman constant, T is the temperature and r is the radius of the Moon (1.7374 108 cm). k replaces the r with r + z and solving for N. In Eq. (1), B is the line brightness in Rayleighs, g is where z is the integration variable. For our analysis we use a static the fluorescence efficiency factor and N is the line-of-sight column temperature of 120 K, which is appropriate for our twilight observa- number density in cm2. When possible, we estimate the g-factor of tions. We note that only f and k are functions of z, which allows us 3 a line at the time of the observations. Many of the initial values for to take N (cm ) out of the integral in Eq. (2) when using Eq. (3) for the g-factors are taken from Gladstone et al. (2010a), when LCROSS N(z). Since the spectrum that we present in Fig. 3 is an average spec- impacted the Moon on October 9, 2009. For species not examined trum weighted by the exposure time over many different shadow by Gladstone et al. (2010a), we calculate the initial g-factors using heights (z0) and altitudes (z1), we calculate values for the integrand the BASS2000 solar spectrum (Paletou et al., 2009). Those g-factors in Eq. (2) for each shadow height/altitude combination and average are scaled using the ratio of the Lyman-a g-factor derived from the them using the exposure time as weights. BASS2000 solar spectrum and the Lyman-a g-factor listed in Glad- The partition function, f, is a sum of several partition functions stone et al. (2010a). For Ne I, we are unable to calculate the g-factor describing particles on ballistic orbits, satellite orbits, and those because the BASS2000 solar spectrum does not cover 630 Å. Instead, which escape. These partition functions are given by we use a g-factor from Krasnopolsky and Mumma (2001) and leave ! 2 3 k2 k2 1=2 3 it constant. ð c Þ w fbal ¼ 1=2 c ; k e c ; k w ; ð4Þ Using data from the Solar Dynamics Observatory’s EVE (Ex- p 2 kc 2 treme Ultraviolet Variability Experiment) instrument, we first cal- ! 1=2 culate the Lyman-a g-factor and note that it has a 30% variability. 2 k2 k2 3 Using Lyman-a as a proxy for the variation in each of the lines, we f ¼ c ewc ; k w ; ð5Þ sat p1=2 k 2 scale the g-factors of each line in proportion to the Lyman-a c g-factor we calculate for each day of the mission. We can apply this 0 1 1=2 scale to the LAMP data obtained after April 2010, shortly after the 2 2 1 B 3 3 kc k 3 3 C launch of EVE. For the LAMP observations obtained prior to April f ¼ @C c ;k ew C c ;k w A; esc p1=2 2 2 k 2 2 2010, we assume a constant value scaled with the median c Lyman-a g-factor between April 2010 and January 2011, when ð6Þ solar activity was relatively stable. We list the wavelength at line 2 center and our adopted g-factor in the second and third columns where w = k /(k + kc), f = fbal + fsat + fesc and c and C are the incom- of Table 1, respectively. Many of the species we searched for using plete and complete c functions. In column 5 of Table 1, we list our these data have multiple lines in our wavelength range. We chose 3r upper limits for each species following from the above only to list the results of the brightest line because these give the algorithm. most stringent upper limit. 5. Discussion

Table 1 depicts the upper limits we derive, and compares them to the most stringent previously derived values for surface number density. Table 1 is arranged by species mass. As noted above, our upper limits are computed for a common reference temperature (which sets the scale height conversion of column density to surface density) of T = 120 K. It should also be noted that the previously derived upper limits listed in Table 1 are not strictly at the same local time as our measurements. Number densities can vary for a variety of reasons, for example if a species is condensable during the lunar night. This, however, turns out to be a minor issue with our observations because over 75% of the total integration time occurs two hours before dawn, and two hours after dusk. At these times, condensable species, such as Ar, are most likely to be detected because the atmosphere is cooling and the scale height of many species is below LRO. Finally, Fig. 3 shows that there is very little difference between the dawn spectrum (red curve) and Fig. 4. LAMP effective area as a function of wavelength. dusk spectrum (blue curve). 686 J.C. Cook et al. / Icarus 225 (2013) 681–687

Table 1 The measured upper limits for 27 species.

Species Wavelength (Å) g-Factor (s1) LAMP upper limits Previous upper limits Ratio (previous/current) Brightness (lRayleighs) Surface density (cm3) Surface density (cm3) H 1025.7 4.3 106 460 24 <17a 0.71 B 1825.7 1.0 104 86 0.46 None known 1 C 1656.9 2.2 105 56 1.6 <200a 125 C+ 1335.7 4.4 105 47 0.63 None known 1 N 1135.0 1.4 107 66 340 <600 (2r)b 1.8 N+ 1085.7 1.2 106 220 130 None known 1 O 1304.9 8.9 106 58 5.4 <500a 93 O+ 834.5 5.8 107 480 680 None known 1 Nec 630.5 2.5 107 1100 4400 8 104d 18 Mg 1827.9 3.5 105 87 3.4 <6000 (5r)e 1800 Al 1766.4 2.2 104 150 1.1 <55 (5r)f 50 Si 1845.5 1.8 104 91 0.90 <48 (5r)a 53 P 1775.0 5.6 105 110 4.2 None known 1 S 1807.3 7.3 105 80 2.3 <150a 65 Cl 1335.7 8.2 106 47 15 None known 1 Ar 1048.2 5.3 108 420 2.3 104 3.5 104 1.5 Ca 1883.2 1.8 105 89 16 <1 (5r)a 0.06 Sc 1744.7 2.5 105 150 24 None known 1 Mn 1785.3 5.3 106 68 78 None known 1 Fe 1851.7 1.4 105 98 45 <380 (5r)a 8.4 Co 1822.4 1.5 105 87 41 None known 1 Zn 1589.6 4.2 106 103 220 None known 1 As 1890.4 9.6 104 103 1.4 None known 1 Xe 1469.6 2.1 106 100 3000 <3000a 1.0 Au 1879.8 2.4 105 110 1100 None known 1 Hg 1849.5 7.0 104 97 39 None known 1 CO 1510 1.9 107 77 710 <1.4 104a 20

a Feldman and Morrison (1991). b Fastie et al. (1973). c Constant g-factor. d Hoffman et al. (1973). e Stern et al. (1997). f Flynn and Stern (1996).

We clearly detect the presence of He and H2. The detection of tively. Compared to the upper limits derived from the average He was previously reported in Stern et al. (2012). Based on those spectrum, the upper limits for the spectrum at dawn and dusk limb observations, we measured a surface density of are lower by 5% and 38%, respectively. In both cases, our upper lim- 7 103 cm3. In these nadir observations, we report a number its are lower than the LACE upper limits. In the case of Ar, the upper density of 6.5 103 cm3 at dusk and 10.6 103 cm3 at dawn limit may be slightly inflated since the spectrum in Fig. 3 shows with errors of 50 cm3. These numbers are in agreement with Stern there residual continuum flux. In the case of Ne, our upper limit et al. (2012) and the original LACE detection (Hoffman et al., 1973) is over a magnitude smaller than the measured LACE value. The at dusk, but lower at dawn. We also show the first spectroscopic reason for our non-detection of Ne and Ar is uncertain. Several pos- detection of native H2 in the lunar atmosphere, reported elsewhere sible reasons include: (i) the lunar atmosphere has changed since in Stern et al. (2013). The presence of H2 had been predicted at the LACE measurements, (ii) both the Ne and Ar sources were local- 1.2 104 cm3 at night (Hodges, 1973). LACE had measured a sur- ized near the Apollo 17 site but their presence is averaged out in face density of 6.5 104 cm3 at night (Hoffman et al., 1973). our spectrum or (iii) the detection or identification of Ar and Ne These measurements were conducted after several cool down in the LACE data was incorrect. We plan to explore this mystery tests, and Hoffman et al. (1973) considered this to be an upper in greater detail by searching for Ar and Ne at different local times limit. Our observations show that H2 is more than an order of and locations on the Moon. This will also be the subject of a future magnitude lower than either of these values, with a surface paper (Cook et al., 2014). density of 1.7 ± 0.4 103 cm3 at dusk and increasing to 2.1 ± Our upper limits on several species (O, Mg, Al, and Ca) are near 0.3 103 cm3 at dawn. the predicted levels for an atmosphere where sputtering is a signif- In total, 27 species upper limits are presented in Table 1.Of icant component. Wurz et al. (2007) used Monte Carlo models to these, 13 species have never before published upper limits, the estimate the surface density of several species such as O (6– other 14 having been previously published and many with higher 10 cm3), Mg (1–1.5 cm3), Al (0.5–1.5 cm3) and Ca (1–2 cm3). upper limits than we derived. In almost every case for the 14 spe- We find that our upper limit on O, Al is in the predicted cies with previously published upper limits, our revised upper lim- range. However, our results for Mg and Ca are slightly greater than it is one to two orders of magnitude more stringent that the the predictions. Our results are unable to constrain the significance previous value, and sometimes more so. of the sputtering process in maintaining the lunar atmosphere. Absent from our discussion thus far are Ar and Ne. In our wave- Additional data that will be collected as the mission continues length range, each species have two lines. In the case of Ar, they are may help to answer this question. 1048 and 1066 Å and for Ne they are 630 and 735 Å. As discussed in the introduction, both Ar and Ne were certainly anticipated 6. Conclusion based on LACE results. However, as can be seen in Fig. 3, neither species are detected at the level of our noise. Our measured upper Using data collected by the LAMP instrument aboard NASAs 4 3 3 3 limits for Ar and Ne are 2.3 10 cm and 4.4 10 cm , respec- LRO spacecraft in lunar orbit since 2009, we analyzed 7.7 million s J.C. 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