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UV spectrum of Enceladus
Icarus 220 (2012) 29–35
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UV spectrum of Enceladus Mark Zastrow a,⇑, John T. Clarke a, Amanda R. Hendrix b, Keith S. Noll c a
Center for Space Physics, Boston University, 725 Commonwealth Ave., Boston, MA 02215, United States Jet Propulsion Laboratory, California Institute of Technology, Mail Stop 230-250, Pasadena, CA 91109, United States c Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, United States b
a r t i c l e
i n f o
Article history: Received 1 December 2011 Revised 1 April 2012 Accepted 3 April 2012 Available online 21 April 2012 Keywords: Saturn, Magnetosphere Saturn, Satellites Enceladus Ices, UV spectroscopy
a b s t r a c t We present a far ultraviolet (FUV) spectrum of Saturn’s moon Enceladus from the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST). We have put upper limits on emission from C, N, and O lines in Enceladus’ atmosphere and column densities for the C lines assuming solar resonance scattering. We find these upper limits to be relatively low—on the order of tens to thousands of Rayleighs and with C column densities on the order of 108–1015 cm2, depending on the assumed source size. We also present a segment of a reflectance spectrum in the FUV from 1900–2130 Å. This region was sensitive to the different ice mixtures in the model spectra reported by Hendrix et al. (Hendrix, A.R., Hansen, C.J., Holsclaw, G.M. [2010]. Icarus, 206, 608). We find the spectrum brightens quickly longward of 1900 Å, constraining the absorption band observed by Hendrix et al. from 170 to 190 nm. We find our data is consistent with the suggestion of Hendrix et al. of the presence of ammonia ice (or ammonia hydrate) to darken that region, and also possibly tholins to darken the mid-UV, as reported by Verbiscer et al. (Verbiscer, A.J., French, R.G., McGhee, C.A. [2005]. Icarus, 173, 66). Ó 2012 Elsevier Inc. All rights reserved.
1. Introduction Seven months after the Cassini spacecraft arrived in the saturnian system in July 2004, its first flyby of Enceladus encountered striking magnetic disturbances. The field lines were draped around the moon as if it were obstructing the plasma flow, suggesting a strong electromagnetic interaction (Dougherty et al., 2006). Closer encounters on later passes revealed spectacular geysers of fine icy particles feeding a large plume spewing from the ‘‘tiger stripes’’— surface fractures in geologically youthful and rugged terrain in the southern polar region of Enceladus (Porco et al., 2006). This finding raised the tantalizing possibility of a subsurface ocean of liquid water. It also shed light on Enceladus’ role within Saturn’s magnetosphere, suggesting it is one of its main sources of plasma. Stellar occultation observations by Cassini’s Ultraviolet Imaging Spectrograph (UVIS) and direct measurements from the Ion Neutral Mass Spectrometer (INMS) confirmed the southerly existence and water ice composition of the plume and suggested a source of 200 kg s1 of H2O (Hansen et al., 2006, 2011; Waite et al., 2006). Burger et al. (2007) used a 3-D Monte Carlo model to simultaneously model the UVIS column density and INMS water gas density results. They found the data were consistent with a twocomponent atmosphere: a thin global H2O atmosphere with a slant column 1013 cm2 neutral density and a source rate of ⇑ Corresponding author. E-mail address: [email protected] (M. Zastrow). 0019-1035/$ - see front matter Ó 2012 Elsevier Inc. All rights reserved. http://dx.doi.org/10.1016/j.icarus.2012.04.002
1026 H2O s1, and a much larger localized plume with a source of 1028 H2O s1 (300 kg s1). From this neutral water source, they estimated a magnetospheric plasma mass loading rate (via charge exchange) of 2–3 kg s1. Khurana et al. (2007), working from magnetic field flyby data, estimated a similar plasma mass loading rate of 0.6–2.8 kg s1, based upon the inference of a pickup source size of 4 REn (where 1 REn = 250 km). However, the plasma deflections at Enceladus were recorded by the Cassini Plasma Spectrometer (CAPS) at distances as large as 30 REn from Enceladus (Dougherty et al., 2006; Tokar et al., 2006), implying an electrodynamic interaction between Saturn’s magnetosphere and Enceladus that requires a much larger mass loading rate. As detailed by Pontius and Hill (2006), depending on the coupling model used, the source rate can be taken to be proportional to either the Pedersen conductivity of Saturn’s ionosphere or the Alfvén conductivity in the Alfvén wings near Enceladus. The plausible range of those parameters implies, at minimum, a plasma mass loading rate of order 100 kg s1. According to this evidence, Enceladus’ role at Saturn appears to be analogous in nature to that of Io’s at Jupiter—a mass loading region comparable in extent, and interacting strongly with Saturn’s corotating magnetic field and plasma. Such magnetospheric planet–satellite interactions have been the subject of intensive research for decades, with particular insight gleaned from the Galilean moons of Jupiter (see Kivelson et al., 2004). The electric potential induced by a satellite’s movement relative to the magnetic field drives ionospheric currents,
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as well as a pickup current as newly ionized particles are incorporated into the plasma flow (Goertz, 1980; Hill et al., 1983). These currents propagate as shear Alfvén waves along magnetic flux tubes to the planet’s ionosphere (Neubauer, 1980). At Jupiter, these produce bright auroral emissions (10–100 kR, where 1 kR = 109 photons cm2 s1 into 4p steradians) located equatorward of the main auroral oval at the magnetic footprint of the field line threading each satellite (Clarke et al., 2004). The currents can also produce auroral emission at the satellites, which has been detected at Io (Roesler et al., 1999) and Ganymede (Hall et al., 1998; Feldman et al., 2000). Auroral evidence for the magnetospheric interaction has also been sought in the case of Enceladus. A study by Wannawichian et al. (2008) put an upper limit on any auroral footprint on Saturn to be a few kR, and the auroral footprint was definitively detected by Pryor et al. (2011) using Cassini UVIS. Three detections of the footprint in the extreme ultraviolet (EUV) and the far ultraviolet (FUV) were reported, with brightnesses varying by a factor of 3, from 450 Rayleighs to 1550 R. They suggest that this range in brightness is driven by variability in the cryovolcanism at the tiger stripes. However, the relative faintness of the footprints suggests that the bulk of the interaction could be localized near Enceladus itself, rather than driving strong currents through the ionosphere of Saturn. With regards to our study, this implies significant airglow/auroral emissions where the currents close near Enceladus. The surface composition of Enceladus has also been long studied. Photometric and spectroscopic observations in the IR have established that Enceladus’ surface is nearly pure water ice (Cruikshank, 1980; Clark et al., 1984; Cruikshank et al., 1998, 2005), but the presence of other compounds has not been ruled out. In particular, ammonia has been suggested as another species and has been definitively detected in the plume (Waite et al., 2009). Emery et al. (2005) reported a possible detection of surface NH3 ice based on spectral features near 2.0 and 2.25 lm. Verbiscer et al. (2006) presented a near-IR spectrum consistent with a mixture of as much as 20% ammonia hydrate (1% NH3H2O). Tentative detections of ammonia in the near-IR have also been made throughout the saturnian system—on Tethys in ground-based observations by Verbiscer et al. (2008) and in dark material seen in Cassini VIMS (Visual and Infrared Mapping Spectrometer) observations of Dione, Phoebe, and Iapetus by Clark et al. (2008). In the UV, Cassini UVIS observations from Hendrix et al. (2010) found that Enceladus’ spectrum in the region 170–190 nm was too dark be made of water ice alone. Hendrix et al. (2010), through modeling of the spectrum, suggested that this characteristic, too, could be explained by the presence of a small amount of ammonia ice (1% in H2O). Hendrix et al. also introduced a small amount of tholins into the mixture in order to darken the NUV portion of the spectrum from 200 nm to 300 nm to be consistent with photometric observations by Verbiscer et al. (2005). We present a segment of a UV spectrum of Enceladus taken with the Cosmic Origins Spectrograph (COS) (Green et al., 2003) on the Hubble Space Telescope (HST). We have measured possible
atomic lines of airglow species and placed upper limits on their column densities. (We consider no molecular features due to their relative weakness.) We also present a segment of a reflectance spectrum from 1900 to 2130 Å, a key region that displays many spectral features of important molecules, including the absorption feature that Hendrix et al. (2010) suggests could be due to the presence of ammonia and/or tholins.
2. Observations The observations were carried out over four HST orbits on May 5–6, 2010 (Table 1). Enceladus’s leading and trailing hemispheres were observed half an Enceladus orbit (16 h) apart, near its western and eastern elongations. The solar phase angle was 4.3°. COS is a highly sensitive UV spectrograph that features a unique slitless design. In lieu of a slit, it has a small aperture, a small field of view, and derives its resolving power from Hubble’s stability and pointing capabilities (Osterman et al., 2011). All exposures were made through COS’ Primary Science Aperture (PSA), a circular aperture 700 lm (2.5 arcsec) in diameter (Shaw et al., 2009). COS has both far ultraviolet (FUV) and near ultraviolet (NUV) channels. Our observations include spectra from each channel, using the G140L and G230L gratings in the FUV and NUV respectively. The FUV channel’s detector is a large format cross delay line (XDL) micro-channel plate consisting of two 16,384 1024 pixel segments (designated A and B). The NUV detector is a single 1024 1024 multi-anode micro-channel array (MAMA) that images three separate, noncontiguous stripes of a spectrum (designated A, B, and C). In the following results, the airglow emission analysis was conducted on segment A of the FUV channel, over a wavelength range of 1267–1800 Å, shown in Fig. 1 smoothed to a 25 pixel (2 Å) boxcar average. Wavelengths longward of 1800 Å and all of segment B (which covers the wavelength range 300–1095 Å) were ignored as the signal-to-noise ratio was too low to be reliable. In Fig. 1, we restrict the wavelengths shortward of 1320 Å to the data from exposures taken while HST was on the nightside of the Earth to avoid contamination of the 1304 Å feature from the geocorona. We believe, however, that much of what is detected is still nightglow from Earth and possibly emission from the greater saturnian magnetosphere; we have no reason to believe that any of the observed emission is from Enceladus or its plume, as detailed below. For the surface composition analysis, a reflectance spectrum was obtained from stripe A of the NUV channel, restricted to the wavelengths longward of 1900 Å up to its limit at 2130 Å (Fig. 2). Shortward of 1900 Å, again, the instrument sensitivity was too low to be reliable. In considering possible hemispheric differences in albedo, we also used the spectrum from stripe B of the NUV channel (2800–3200 Å). Throughout this paper, we will refer to far- and near-UV wavelength ranges according to the COS nomenclature for the NUV and FUV detectors.
Table 1 Log of HST observations of Enceladus. For each grating, the exposures were split between eastern and western elongations. The wavelength ranges listed are the ranges presented here; unreliable data was truncated from the COS spectra and the low wavelength end of the STIS spectrum. The range of resolving power for COS data was interpolated from the published range, where the power increases roughly linearly with wavelength (Dixon and Niemi, 2010). Epoch
Instrument
No. of exposures
Exposure time (s)
Grating
Wavelength range (Å)
Resolving power
Designation
Aperture
Aperture size (arcsec)
May 5–6, 2010
COS
4
1330
G140L
1267–1800
2300–3500
Segment A
PSA
2.5
4
840
G230L
1906–2130 2800–3200
2400–2650 3400–3900
Stripe A Stripe B
2 2
2400 600
G230L
1906–2700
500
52 0.5
0.5
1999–2000
STIS
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M. Zastrow et al. / Icarus 220 (2012) 29–35
Fig. 1. FUV spectrum, used for airglow analysis; disk-averaged, smoothed to 25 pixel (2 Å) boxcar average.
Fig. 2. Segment of NUV spectrum used for surface analysis (disk-averaged, unsmoothed).
In our analysis, we used the solar spectrum for the date of the observation from SOLSTICE on the SORCE satellite at 1 Å resolution (McClintock et al., 2000), convolved with the line spread function (LSF) of COS.
3. Upper limits on airglow emission We present upper limits on airglow emissions of O, C, and N at various lines in Table 2. Enceladus appears essentially as a point
source to COS; with a radius of 250 km it appears with a 0.1 arcsec diameter, close to the instrument’s full width half maximum. The calculated FWHM of the G140L grating is 1 Å across its entire wavelength range. The upper limits on airglow were determined by smoothing over a 2 Å (25-pixel) boxcar average, adding 2r above the noise level from the surrounding 10 Å, and integrating over a 2 Å wide segment for each potential feature. (Empirical LSFs have been released for the G130M and G160M gratings in the FUV mode, but are not yet available for the G140L grating.) Since COS cannot spatially resolve any of the structure of the plume or Enceladus’ atmosphere, and the actual strength of the emission depends on the size of the emitting region, we attempt to give a sense of the plausible range of values for the emission by presenting three columns in Table 2. These correspond to three cases: a global atmosphere with a simple circular cross section of 1.5 REn (375 km); a localized plume source with a circular cross section of diameter 150 km (the width of the plume derived by Hansen et al. (2011) from a Cassini UVIS observation of a solar occultation on May 18, 2010, twelve days after our observations); and the E-ring torus, which would fill the field of view of COS. Although the PSA has a physical radius of 1.25 arcsec, aberrated light up to 2 arcsec from its center can enter the system. However, to derive an upper limit, we assume a field of view simply corresponding to the size of the PSA. Any local emission that is observed could be due to either: atomic or molecular emission, solar resonance fluorescence scattering, or reflected sunlight from the surface. By assuming solar resonance scattering, we can calculate upper limits on the column densities from resonance lines of carbon, also given in Table 2. We note that O I 1356 Å is a forbidden transition and should rather be taken as a measure of electron collisions (discussed in S5). The solar flux was determined for each emission line from the spectrum for the date of the observations from the SOLSTICE instrument. A prime candidate species is atomic oxygen, which has been detected in the saturnian magnetosphere from dissociation of water molecules by Hansen et al. (2006) and Melin et al. (2009). Melin et al. found it to be distributed in a torus centered at Enceladus’ orbit (a distance of 4 Saturn radii.) O emission has been previously detected in the atmospheres of Europa and Ganymede at 1304 and 1356 Å by Hall et al. (1998); in the saturnian system, O2 produced by radiolysis has been detected in the atmosphere of Rhea by Teolis et al. (2010). We note that although O I emission at 1304 Å can be seen in the disk-averaged spectrum (Fig. 1), even though we have restricted that portion of the spectrum to nightside exposures, it is easily explained by nightglow from Earth’s atmosphere and/or background emission from the saturnian magnetosphere. Both sources would fill the field of view of COS; in this case, we derive an O I 1304 Å emission of 1 Rayleigh. This is significantly less than the nightglow of 7 Rayleighs observed by Chakrabarti et al. (1984) at this line in uplooking spectra from a satellite in a polar orbit
Table 2 Upper limits on airglow emission lines obtained from COS, with column densities for C assuming solar resonance fluorescence scattering. 1 Rayleigh = 106 photons cm2 s1 into 4p steradians. Wavelength data from Ralchenko et al. (2011). Species
OI CI CI CI C II NI NI NI
k (Å)
1356 1277 1280 1561 1335 1419 1493 1743
Atmosphere (r = 1.5 RE)
Plume (r = 75 km)
Emission (Rayleighs)
Emission (Rayleighs)
88 17 13 45 33 42 34 360
Column density (atoms cm2) 5.2 1012 1.2 1014 1.2 1012 4.0 1011
2200 420 320 1200 830 1100 870 9200
E-ring torus (aperture radius = 1.25 arcsec) Column density (atoms cm2) 1.3 1014 3.1 1015 3.0 1013 1.0 1013
Emission (Rayleighs) 0.20 0.04 0.03 0.10 0.07 0.10 0.08 0.82
Column density (atoms cm2) 1.2 1010 2.8 1011 2.6 109 8.9 108
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Fig. 4. Composite spectrum of COS and UVIS data with STIS spectrum overplotted (K. Noll).
Fig. 3. Top: Composite disk-averaged spectrum, of Cassini UVIS (Hendrix et al., 2010) and unsmoothed COS data >1900 Å. Bottom: Same, but COS data smoothed over 25 pixel (2 Å) boxcar average.
at local midnight from an altitude of 650 km (higher than HST’s 559 km); it is also less than the saturnian magnetospheric emission at this line detected by Melin et al. (2009) in the vicinity of Enceladus’ orbit at an intensity of 3 Rayleighs. Furthermore, the FWHM of the feature is 7–8 Å, which is consistent with the FWHM of the same feature as measured in previous observations of Earth nightglow contamination in HST (observing programs 11860 and 12414). Based on this evidence, we have no reason to believe that any of our observed 1304 Å emission is from Enceladus. 4. Reflectance spectrum The disk-averaged reflectance spectrum in the wavelength range 1900–2130 Å is obtained by dividing the radiance measured at HST (units of photons cm2 s1) by the radiance expected from a Lambertian surface normally illuminated by the Sun:
I=F ¼
F En =XEn F =p
ð1Þ
where FEn is the measured irradiance of Enceladus, and XEn is the solid angle of Enceladus. F/p is the radiance expected from a Lambertian disk, where F is the solar irradiance (units photons cm2 s1) scaled to 9.503 AU (the heliocentric distance of Saturn on the date of the observation). In Fig. 3, at wavelengths shortward of 1906 Å we include the Enceladus reflectance spectrum from Cassini UVIS taken by Hendrix et al. (2010). We note that these reflectance spectra are given in terms of the canonical I/F with no correction for phase angle; they are equivalent to the product of the geometric albedo times the phase function. In Fig. 4, we overplot a portion of a spectrum from observations from HST Space Telescope Imaging Spectrograph (STIS) carried out over 1999–2000 by Keith Noll. Our phase
angle was 4.3 degrees, while those of Hendrix et al. (2010) and Noll were both 2 degrees. However, we do not expect this difference to factor into our discussion of the broad spectral features. The trailing hemisphere of Enceladus is known to be slightly brighter than the leading hemisphere by 0.06–0.12 mag (approximately 6–12%), both in the visible (Verbiscer et al., 2005) and the near-IR (Buratti et al., 1998; Momary et al., 2000; Verbiscer et al., 2005, 2006). We see this effect in the B stripe of the NUV channel (2800–3200 Å); however, in the A stripe of the NUV channel (1906–2130 Å), we find no significant difference in the continuum reflectivity of the two hemispheres shortward of 2105 Å (Fig. 5). Possible mechanisms for any difference in albedo between hemispheres include micrometeoroid impacts, magnetospheric plasma bombardment, and deposition of E-ring material (Hamilton and Burns, 1994; Buratti et al., 1998). With respect to the possibility of magnetospheric plasma, Enceladus’ trailing hemisphere is preferentially bombarded due to the fact that Saturn’s magnetospheric plasma corotates faster than Enceladus at its orbital distance. In particular, Sack et al. (1992) found that bombardment of energetic ions tended to brighten the spectrum in the visible but darken it in the FUV. This process would appear to be consistent with our nondetection of a brighter trailing hemisphere in the A stripe. Another possibility is that the spectral properties of the local material deposited by the E-ring and its source, the plume, could account for the lack of hemispheric albedo difference in the FUV. Schenk et al. (2011) noted that the plume and E-ring deposition, rather than external bombardment, appeared to dominate spectral color patterns on the surface of Enceladus and block color features seen on other saturnian satellites. If plume and E-ring deposition is the dominant bombardment mechanism in determining the modification of the surface reflectance of Enceladus, the lack of a hemispheric difference in albedo in the FUV could suggest that the plume/E-ring is depositing material fast enough to counteract the effects of micrometeoroid and plasma bombardment in this spectral region.
5. Discussion and conclusions 5.1. Airglow emission The upper limits on the airglow emission of Enceladus are quite low. Even assuming a small emitting region roughly the size of Enceladus’ plume, the upper limits are in the hundreds to thousands of Rayleighs, on the same order as the faint auroral footprints reported by Pryor et al. (2011). Similarly, the column densities for C
M. Zastrow et al. / Icarus 220 (2012) 29–35
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Fig. 5. Comparisons of spectra of leading and trailing hemispheres from COS, smoothed to 25 pixels (2 Å) boxcar average. On the left, in the mid-UV, the trailing hemisphere can be seen to be brighter than the leading hemisphere. On the right, in the FUV, this effect is not seen.
are low—well below the H2O upper limit of 1016 cm2 from UVIS reported by Hansen et al. (2011). Although a plume-sized source for C I 1280 results in an upper limit for C of 3.1 1015 cm2, or roughly one-third the density of the H2 O plume, C I 1277 constrains it to 5 1012 cm2 or one part in 2000. (For comparison, Hansen et al. (2008) placed an upper limit on CO of 0.3% of the plume, or a column density of 3 1014 cm2.) While the faintness of the footprints might suggest that the bulk of the interaction could be found near Enceladus itself, we find no evidence for significant airglow/auroral emissions in the extended atmosphere of Enceladus. For future work, O 1356 Å presents an opportunity to constrain the interaction between Enceladus and Saturn’s magnetosphere. Since it is a forbidden transition, one could reasonably take its upper limit of emission to be a limit of collisional excitation. The peak UV emission cross section for electron impact excitation on O 1356 Å is 9.0 1018 cm2 at 16 eV, and for e on O2 the cross section is 8.8 1018 cm2 at 110 eV (Meier, 1991). Therefore, with knowledge of the energy distribution of the incoming electrons, a theoretical upper limit on O column density could be obtained. The electron temperature at Enceladus is quite cold—it has been measured by CAPS to be 1 eV (Sittler et al., 2008)— and the 10 eV electrons necessary to excite O 1356 would clearly be in the tail of the distribution. However, published Cassini results do not currently reveal the full energy distribution of electrons at Enceladus. Modeling would be necessary to extrapolate downward from higher energy observations, which is beyond the scope of this paper. 5.2. Reflectance spectrum Hendrix et al. (2010) noted a region at 170–190 nm that was too dark to be explained by a pure water ice spectrum. This indicated the presence of another species that was bright in the visible, but dark in the FUV. Hendrix et al. (2010) singled out an intimate mixture of water ice and ammonia ice with global coverage of Enceladus as the most promising explanation, as ammonia ice has a strong absorption edge at 190 nm. Hendrix et al. found good agreement with a model spectrum of a mixture of 1% NH3 in H2O. We find that our data is consistent with the H2O + NH3 models reported in the model spectra of Hendrix et al. (2010). Our data
suggests that the spectrum begins to quickly brighten longward of 1900 Å, constraining the extent of this absorption feature. (The signal-to-noise ratio from COS and STIS is too low to be reliable shortward of 1906 Å.) This matches the overall behavior of the H2O + NH3 models of Hendrix et al. (2010, Fig. 9), which exhibit a sharp edge longward of 1900 Å. The model that best fits our data in the FUV is the one using the Dressler and Schnepp (1960) ammonia data as rescaled and included in the compilation by Martonchik et al. (1984), also with a grain size of 30 lm (Hendrix et al., 2010, Fig. 9, top left). Other models by Hendrix et al. (2010) included different NH3 measurements, including optical constants by Dawes et al. (2007) and reflectance data by Hapke et al. (1981) and Pipes et al. (1974). Hendrix et al. kept the ratio of water ice to NH3 constant, and so the only differences in the models arise from the inconsistencies observed in the different lab spectra of NH3. All of them exhibit a steep absorption edge, but its location varies from 1900 to 2100 Å. In addition, a ‘‘trough’’ in reflectance (an absorption peak) is seen just shortward of the absorption edge in the optical constants measured by Dawes et al. (2007) at Enceladus-like temperatures (75 K). It also appears more weakly in the measurements by Dressler and Schnepp (1960) in the Martonchik et al. collection at 77 K (although not at 125 K) (Hendrix et al., 2010, Fig. 5). It is not seen in the reflectance measurements by Hapke et al. (1981) and Pipes et al. (1974). In the model spectra, this feature appears most strongly in the model using the Dawes et al. measurements, where the ammonia ice grain size is increased to 30 lm to darken the spectrum to match the 170–190 nm absorption band, (Hendrix et al., 2010, Fig. 9). To these models, Hendrix et al. (2010) (Fig. 9, bottom) also added 1% tholins in order to darken the NUV longward of 190 nm to match the data point at 275 nm of Verbiscer et al. (2005), varying the grain size—the larger the grain, the greater the absorption. This darkened the NUV, while making no significant difference in the wavelength range discussed in the preceding paragraph. In Fig. 6, we present a model spectrum that combines water ice with 1% NH3, and 1% tholins, but using the Martonchik et al. NH3 constants. The difference in the location and strength of the ammonia trough produces an acceptable fit to our COS and STIS data in the FUV while still darkening the NUV to match the STIS data out to 2800 Å.
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Fig. 6. Composite spectrum from Fig. 4, but with the COS data smoothed as in Fig. 3. The model using Martonchik et al. (1984) optical constants is overplotted.
Hendrix et al. noted that the dearth of published optical constants for ices and tholins at Enceladus temperatures in this spectral range hampered their analysis. Additionally, if either water or ammonia ice are actually slightly absorbing in the NUV, tholins may not be necessary to match the data of Verbiscer et al. (2005). Nor do we have any way to know which set of NH3 optical constants (Dawes et al. or Martonchik et al.) are more accurate. Nevertheless, we present Fig. 6 to demonstrate that our data is consistent with the range of H2O + NH3 + tholins model spectra of Hendrix et al. in the near- to mid-UV. We note one discrepant feature in the model spectrum (Fig. 6) around 190–195 nm, where the ammonia ice trough darkens the model relative to the UVIS, COS, and STIS spectra. One possible explanation is that, as noted by Hendrix et al. (2010), many published lab spectra of water ice also show a trough in reflectance varying in strength from measurement to measurement—similar to that discussed in the previous paragraph in ammonia ice, but at a slightly shorter wavelength: ranging from 1860 to 1930 Å as measured by Hapke et al. (1981); Pipes et al. (1974), and Gougen (obtained by Hendrix et al. via personal communication). However, this feature was not incorporated into their models as it was not present in the dataset they based them on, the optical constants of Warren and Brandt (2008). If it were, it might even out the model spectrum over the 170–200 nm range by darkening the shortwavelength side of it. Coupled with a slightly smaller ammonia grain size (to slightly brighten it to fit to the data), it could reduce the discrepant feature or eliminate it altogether. Further analysis is hampered by the difficulty in precisely measuring the spectra of ices in the lab, the resulting inconsistencies of the existing ice data, and the general lack of data on ices at Enceladus temperatures. Keeping in mind the varying characteristics of the different ice measurements, we observe that our results are in agreement with the H2O + NH3 + tholins model spectra produced by Hendrix et al. (2010), and support the possibility of the presence of small amounts of ammonia ice and tholins. Acknowledgments This work is based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the AURA Inc. for NASA. This research was supported by NASA Grant HST-GO-11645.01-A from the Space Telescope Science Institute to Boston University. Marty Snow kindly provided the solar spectrum from SORCE SOLSTICE. MZ thanks Andrew West and Supriya Chakrabarti for helpful conversations.
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Ralchenko, Y., Kramida, A.E., Reader, J., Team, N.A., 2011. NIST Atomic Spectra Database, Version 4.1. National Institute of Standards and Technology, Gaithersburg, MD. Available from:. Roesler, F.L. et al., 1999. Far-ultraviolet imaging spectroscopy of Io’s atmosphere with HST/STIS. Science 283, 353–357. Sack, N.J., Johnson, R.E., Boring, J.W., Baragiola, R.A., 1992. The effect of magnetospheric ion bombardment on the reflectance of Europa’s surface. Icarus 100, 534–540. Schenk, P., Hamilton, D.P., Johnson, R.E., McKinnon, W.B., Paranicas, C., Schmidt, J., Showalter, M.R., 2011. Plasma, plumes and rings: Saturn system dynamics as recorded in global color patterns on its midsize icy satellites. Icarus 211, 740– 757. Shaw, B., Massa, D., Kaiser, M.E., 2009. COS Data Handbook, Version 1.0. STScI, Baltimore. Sittler, E.C. et al., 2008. Ion and neutral sources and sinks within Saturn’s inner magnetosphere: Cassini results. Planet. Space Sci. 56, 3–18. Teolis, B.D. et al., 2010. Cassini finds an oxygen–carbon dioxide atmosphere at Saturn’s icy moon Rhea. Science 330, 1813–1815. Tokar, R.L. et al., 2006. The Interaction of the atmosphere of Enceladus with Saturn’s plasma. Science 311, 1409–1412. Verbiscer, A.J., French, R.G., McGhee, C.A., 2005. The opposition surge of Enceladus: HST observations 338–1022 nm. Icarus 173, 66–83. Verbiscer, A.J., Peterson, D.E., Skrutskie, M.F., Cushing, M., Helfenstein, P., Nelson, M.J., Smith, J.D., Wilson, J.C., 2006. Near-infrared spectra of the leading and trailing hemispheres of Enceladus. Icarus 182, 211–223. Verbiscer, A.J., Skrutskie, M., Wilson, J., Nelson, M., Helfenstein, P., 2008. Ammonia hydrate in the saturnian system. Bull. Am. Astron. Soc. 40, 510. Waite, J.H. et al., 2006. Cassini ion and neutral mass spectrometer: Enceladus plume composition and structure. Science 311, 1419–1422. Waite, J.H. et al., 2009. Liquid water on Enceladus from observations of ammonia and 40Ar in the plume. Nature 460, 487–490. Wannawichian, S., Clarke, J.T., Pontius, D.H., 2008. Interaction evidence between Enceladus’ atmosphere and Saturn’s magnetosphere. J. Geophys. Res. 113, 07217. Warren, S.G., Brandt, R.E., 2008. Optical constants of ice from the ultraviolet to the microwave: A revised compilation. J. Geophys. Res. 113, 14220.
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UV spectrum of Enceladus Mark Zastrow a,⇑, John T. Clarke a, Amanda R. Hendrix b, Keith S. Noll c a
Center for Space Physics, Boston University, 725 Commonwealth Ave., Boston, MA 02215, United States Jet Propulsion Laboratory, California Institute of Technology, Mail Stop 230-250, Pasadena, CA 91109, United States c Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, United States b
a r t i c l e
i n f o
Article history: Received 1 December 2011 Revised 1 April 2012 Accepted 3 April 2012 Available online 21 April 2012 Keywords: Saturn, Magnetosphere Saturn, Satellites Enceladus Ices, UV spectroscopy
a b s t r a c t We present a far ultraviolet (FUV) spectrum of Saturn’s moon Enceladus from the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST). We have put upper limits on emission from C, N, and O lines in Enceladus’ atmosphere and column densities for the C lines assuming solar resonance scattering. We find these upper limits to be relatively low—on the order of tens to thousands of Rayleighs and with C column densities on the order of 108–1015 cm2, depending on the assumed source size. We also present a segment of a reflectance spectrum in the FUV from 1900–2130 Å. This region was sensitive to the different ice mixtures in the model spectra reported by Hendrix et al. (Hendrix, A.R., Hansen, C.J., Holsclaw, G.M. [2010]. Icarus, 206, 608). We find the spectrum brightens quickly longward of 1900 Å, constraining the absorption band observed by Hendrix et al. from 170 to 190 nm. We find our data is consistent with the suggestion of Hendrix et al. of the presence of ammonia ice (or ammonia hydrate) to darken that region, and also possibly tholins to darken the mid-UV, as reported by Verbiscer et al. (Verbiscer, A.J., French, R.G., McGhee, C.A. [2005]. Icarus, 173, 66). Ó 2012 Elsevier Inc. All rights reserved.
1. Introduction Seven months after the Cassini spacecraft arrived in the saturnian system in July 2004, its first flyby of Enceladus encountered striking magnetic disturbances. The field lines were draped around the moon as if it were obstructing the plasma flow, suggesting a strong electromagnetic interaction (Dougherty et al., 2006). Closer encounters on later passes revealed spectacular geysers of fine icy particles feeding a large plume spewing from the ‘‘tiger stripes’’— surface fractures in geologically youthful and rugged terrain in the southern polar region of Enceladus (Porco et al., 2006). This finding raised the tantalizing possibility of a subsurface ocean of liquid water. It also shed light on Enceladus’ role within Saturn’s magnetosphere, suggesting it is one of its main sources of plasma. Stellar occultation observations by Cassini’s Ultraviolet Imaging Spectrograph (UVIS) and direct measurements from the Ion Neutral Mass Spectrometer (INMS) confirmed the southerly existence and water ice composition of the plume and suggested a source of 200 kg s1 of H2O (Hansen et al., 2006, 2011; Waite et al., 2006). Burger et al. (2007) used a 3-D Monte Carlo model to simultaneously model the UVIS column density and INMS water gas density results. They found the data were consistent with a twocomponent atmosphere: a thin global H2O atmosphere with a slant column 1013 cm2 neutral density and a source rate of ⇑ Corresponding author. E-mail address: [email protected] (M. Zastrow). 0019-1035/$ - see front matter Ó 2012 Elsevier Inc. All rights reserved. http://dx.doi.org/10.1016/j.icarus.2012.04.002
1026 H2O s1, and a much larger localized plume with a source of 1028 H2O s1 (300 kg s1). From this neutral water source, they estimated a magnetospheric plasma mass loading rate (via charge exchange) of 2–3 kg s1. Khurana et al. (2007), working from magnetic field flyby data, estimated a similar plasma mass loading rate of 0.6–2.8 kg s1, based upon the inference of a pickup source size of 4 REn (where 1 REn = 250 km). However, the plasma deflections at Enceladus were recorded by the Cassini Plasma Spectrometer (CAPS) at distances as large as 30 REn from Enceladus (Dougherty et al., 2006; Tokar et al., 2006), implying an electrodynamic interaction between Saturn’s magnetosphere and Enceladus that requires a much larger mass loading rate. As detailed by Pontius and Hill (2006), depending on the coupling model used, the source rate can be taken to be proportional to either the Pedersen conductivity of Saturn’s ionosphere or the Alfvén conductivity in the Alfvén wings near Enceladus. The plausible range of those parameters implies, at minimum, a plasma mass loading rate of order 100 kg s1. According to this evidence, Enceladus’ role at Saturn appears to be analogous in nature to that of Io’s at Jupiter—a mass loading region comparable in extent, and interacting strongly with Saturn’s corotating magnetic field and plasma. Such magnetospheric planet–satellite interactions have been the subject of intensive research for decades, with particular insight gleaned from the Galilean moons of Jupiter (see Kivelson et al., 2004). The electric potential induced by a satellite’s movement relative to the magnetic field drives ionospheric currents,
30
M. Zastrow et al. / Icarus 220 (2012) 29–35
as well as a pickup current as newly ionized particles are incorporated into the plasma flow (Goertz, 1980; Hill et al., 1983). These currents propagate as shear Alfvén waves along magnetic flux tubes to the planet’s ionosphere (Neubauer, 1980). At Jupiter, these produce bright auroral emissions (10–100 kR, where 1 kR = 109 photons cm2 s1 into 4p steradians) located equatorward of the main auroral oval at the magnetic footprint of the field line threading each satellite (Clarke et al., 2004). The currents can also produce auroral emission at the satellites, which has been detected at Io (Roesler et al., 1999) and Ganymede (Hall et al., 1998; Feldman et al., 2000). Auroral evidence for the magnetospheric interaction has also been sought in the case of Enceladus. A study by Wannawichian et al. (2008) put an upper limit on any auroral footprint on Saturn to be a few kR, and the auroral footprint was definitively detected by Pryor et al. (2011) using Cassini UVIS. Three detections of the footprint in the extreme ultraviolet (EUV) and the far ultraviolet (FUV) were reported, with brightnesses varying by a factor of 3, from 450 Rayleighs to 1550 R. They suggest that this range in brightness is driven by variability in the cryovolcanism at the tiger stripes. However, the relative faintness of the footprints suggests that the bulk of the interaction could be localized near Enceladus itself, rather than driving strong currents through the ionosphere of Saturn. With regards to our study, this implies significant airglow/auroral emissions where the currents close near Enceladus. The surface composition of Enceladus has also been long studied. Photometric and spectroscopic observations in the IR have established that Enceladus’ surface is nearly pure water ice (Cruikshank, 1980; Clark et al., 1984; Cruikshank et al., 1998, 2005), but the presence of other compounds has not been ruled out. In particular, ammonia has been suggested as another species and has been definitively detected in the plume (Waite et al., 2009). Emery et al. (2005) reported a possible detection of surface NH3 ice based on spectral features near 2.0 and 2.25 lm. Verbiscer et al. (2006) presented a near-IR spectrum consistent with a mixture of as much as 20% ammonia hydrate (1% NH3H2O). Tentative detections of ammonia in the near-IR have also been made throughout the saturnian system—on Tethys in ground-based observations by Verbiscer et al. (2008) and in dark material seen in Cassini VIMS (Visual and Infrared Mapping Spectrometer) observations of Dione, Phoebe, and Iapetus by Clark et al. (2008). In the UV, Cassini UVIS observations from Hendrix et al. (2010) found that Enceladus’ spectrum in the region 170–190 nm was too dark be made of water ice alone. Hendrix et al. (2010), through modeling of the spectrum, suggested that this characteristic, too, could be explained by the presence of a small amount of ammonia ice (1% in H2O). Hendrix et al. also introduced a small amount of tholins into the mixture in order to darken the NUV portion of the spectrum from 200 nm to 300 nm to be consistent with photometric observations by Verbiscer et al. (2005). We present a segment of a UV spectrum of Enceladus taken with the Cosmic Origins Spectrograph (COS) (Green et al., 2003) on the Hubble Space Telescope (HST). We have measured possible
atomic lines of airglow species and placed upper limits on their column densities. (We consider no molecular features due to their relative weakness.) We also present a segment of a reflectance spectrum from 1900 to 2130 Å, a key region that displays many spectral features of important molecules, including the absorption feature that Hendrix et al. (2010) suggests could be due to the presence of ammonia and/or tholins.
2. Observations The observations were carried out over four HST orbits on May 5–6, 2010 (Table 1). Enceladus’s leading and trailing hemispheres were observed half an Enceladus orbit (16 h) apart, near its western and eastern elongations. The solar phase angle was 4.3°. COS is a highly sensitive UV spectrograph that features a unique slitless design. In lieu of a slit, it has a small aperture, a small field of view, and derives its resolving power from Hubble’s stability and pointing capabilities (Osterman et al., 2011). All exposures were made through COS’ Primary Science Aperture (PSA), a circular aperture 700 lm (2.5 arcsec) in diameter (Shaw et al., 2009). COS has both far ultraviolet (FUV) and near ultraviolet (NUV) channels. Our observations include spectra from each channel, using the G140L and G230L gratings in the FUV and NUV respectively. The FUV channel’s detector is a large format cross delay line (XDL) micro-channel plate consisting of two 16,384 1024 pixel segments (designated A and B). The NUV detector is a single 1024 1024 multi-anode micro-channel array (MAMA) that images three separate, noncontiguous stripes of a spectrum (designated A, B, and C). In the following results, the airglow emission analysis was conducted on segment A of the FUV channel, over a wavelength range of 1267–1800 Å, shown in Fig. 1 smoothed to a 25 pixel (2 Å) boxcar average. Wavelengths longward of 1800 Å and all of segment B (which covers the wavelength range 300–1095 Å) were ignored as the signal-to-noise ratio was too low to be reliable. In Fig. 1, we restrict the wavelengths shortward of 1320 Å to the data from exposures taken while HST was on the nightside of the Earth to avoid contamination of the 1304 Å feature from the geocorona. We believe, however, that much of what is detected is still nightglow from Earth and possibly emission from the greater saturnian magnetosphere; we have no reason to believe that any of the observed emission is from Enceladus or its plume, as detailed below. For the surface composition analysis, a reflectance spectrum was obtained from stripe A of the NUV channel, restricted to the wavelengths longward of 1900 Å up to its limit at 2130 Å (Fig. 2). Shortward of 1900 Å, again, the instrument sensitivity was too low to be reliable. In considering possible hemispheric differences in albedo, we also used the spectrum from stripe B of the NUV channel (2800–3200 Å). Throughout this paper, we will refer to far- and near-UV wavelength ranges according to the COS nomenclature for the NUV and FUV detectors.
Table 1 Log of HST observations of Enceladus. For each grating, the exposures were split between eastern and western elongations. The wavelength ranges listed are the ranges presented here; unreliable data was truncated from the COS spectra and the low wavelength end of the STIS spectrum. The range of resolving power for COS data was interpolated from the published range, where the power increases roughly linearly with wavelength (Dixon and Niemi, 2010). Epoch
Instrument
No. of exposures
Exposure time (s)
Grating
Wavelength range (Å)
Resolving power
Designation
Aperture
Aperture size (arcsec)
May 5–6, 2010
COS
4
1330
G140L
1267–1800
2300–3500
Segment A
PSA
2.5
4
840
G230L
1906–2130 2800–3200
2400–2650 3400–3900
Stripe A Stripe B
2 2
2400 600
G230L
1906–2700
500
52 0.5
0.5
1999–2000
STIS
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M. Zastrow et al. / Icarus 220 (2012) 29–35
Fig. 1. FUV spectrum, used for airglow analysis; disk-averaged, smoothed to 25 pixel (2 Å) boxcar average.
Fig. 2. Segment of NUV spectrum used for surface analysis (disk-averaged, unsmoothed).
In our analysis, we used the solar spectrum for the date of the observation from SOLSTICE on the SORCE satellite at 1 Å resolution (McClintock et al., 2000), convolved with the line spread function (LSF) of COS.
3. Upper limits on airglow emission We present upper limits on airglow emissions of O, C, and N at various lines in Table 2. Enceladus appears essentially as a point
source to COS; with a radius of 250 km it appears with a 0.1 arcsec diameter, close to the instrument’s full width half maximum. The calculated FWHM of the G140L grating is 1 Å across its entire wavelength range. The upper limits on airglow were determined by smoothing over a 2 Å (25-pixel) boxcar average, adding 2r above the noise level from the surrounding 10 Å, and integrating over a 2 Å wide segment for each potential feature. (Empirical LSFs have been released for the G130M and G160M gratings in the FUV mode, but are not yet available for the G140L grating.) Since COS cannot spatially resolve any of the structure of the plume or Enceladus’ atmosphere, and the actual strength of the emission depends on the size of the emitting region, we attempt to give a sense of the plausible range of values for the emission by presenting three columns in Table 2. These correspond to three cases: a global atmosphere with a simple circular cross section of 1.5 REn (375 km); a localized plume source with a circular cross section of diameter 150 km (the width of the plume derived by Hansen et al. (2011) from a Cassini UVIS observation of a solar occultation on May 18, 2010, twelve days after our observations); and the E-ring torus, which would fill the field of view of COS. Although the PSA has a physical radius of 1.25 arcsec, aberrated light up to 2 arcsec from its center can enter the system. However, to derive an upper limit, we assume a field of view simply corresponding to the size of the PSA. Any local emission that is observed could be due to either: atomic or molecular emission, solar resonance fluorescence scattering, or reflected sunlight from the surface. By assuming solar resonance scattering, we can calculate upper limits on the column densities from resonance lines of carbon, also given in Table 2. We note that O I 1356 Å is a forbidden transition and should rather be taken as a measure of electron collisions (discussed in S5). The solar flux was determined for each emission line from the spectrum for the date of the observations from the SOLSTICE instrument. A prime candidate species is atomic oxygen, which has been detected in the saturnian magnetosphere from dissociation of water molecules by Hansen et al. (2006) and Melin et al. (2009). Melin et al. found it to be distributed in a torus centered at Enceladus’ orbit (a distance of 4 Saturn radii.) O emission has been previously detected in the atmospheres of Europa and Ganymede at 1304 and 1356 Å by Hall et al. (1998); in the saturnian system, O2 produced by radiolysis has been detected in the atmosphere of Rhea by Teolis et al. (2010). We note that although O I emission at 1304 Å can be seen in the disk-averaged spectrum (Fig. 1), even though we have restricted that portion of the spectrum to nightside exposures, it is easily explained by nightglow from Earth’s atmosphere and/or background emission from the saturnian magnetosphere. Both sources would fill the field of view of COS; in this case, we derive an O I 1304 Å emission of 1 Rayleigh. This is significantly less than the nightglow of 7 Rayleighs observed by Chakrabarti et al. (1984) at this line in uplooking spectra from a satellite in a polar orbit
Table 2 Upper limits on airglow emission lines obtained from COS, with column densities for C assuming solar resonance fluorescence scattering. 1 Rayleigh = 106 photons cm2 s1 into 4p steradians. Wavelength data from Ralchenko et al. (2011). Species
OI CI CI CI C II NI NI NI
k (Å)
1356 1277 1280 1561 1335 1419 1493 1743
Atmosphere (r = 1.5 RE)
Plume (r = 75 km)
Emission (Rayleighs)
Emission (Rayleighs)
88 17 13 45 33 42 34 360
Column density (atoms cm2) 5.2 1012 1.2 1014 1.2 1012 4.0 1011
2200 420 320 1200 830 1100 870 9200
E-ring torus (aperture radius = 1.25 arcsec) Column density (atoms cm2) 1.3 1014 3.1 1015 3.0 1013 1.0 1013
Emission (Rayleighs) 0.20 0.04 0.03 0.10 0.07 0.10 0.08 0.82
Column density (atoms cm2) 1.2 1010 2.8 1011 2.6 109 8.9 108
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M. Zastrow et al. / Icarus 220 (2012) 29–35
Fig. 4. Composite spectrum of COS and UVIS data with STIS spectrum overplotted (K. Noll).
Fig. 3. Top: Composite disk-averaged spectrum, of Cassini UVIS (Hendrix et al., 2010) and unsmoothed COS data >1900 Å. Bottom: Same, but COS data smoothed over 25 pixel (2 Å) boxcar average.
at local midnight from an altitude of 650 km (higher than HST’s 559 km); it is also less than the saturnian magnetospheric emission at this line detected by Melin et al. (2009) in the vicinity of Enceladus’ orbit at an intensity of 3 Rayleighs. Furthermore, the FWHM of the feature is 7–8 Å, which is consistent with the FWHM of the same feature as measured in previous observations of Earth nightglow contamination in HST (observing programs 11860 and 12414). Based on this evidence, we have no reason to believe that any of our observed 1304 Å emission is from Enceladus. 4. Reflectance spectrum The disk-averaged reflectance spectrum in the wavelength range 1900–2130 Å is obtained by dividing the radiance measured at HST (units of photons cm2 s1) by the radiance expected from a Lambertian surface normally illuminated by the Sun:
I=F ¼
F En =XEn F =p
ð1Þ
where FEn is the measured irradiance of Enceladus, and XEn is the solid angle of Enceladus. F/p is the radiance expected from a Lambertian disk, where F is the solar irradiance (units photons cm2 s1) scaled to 9.503 AU (the heliocentric distance of Saturn on the date of the observation). In Fig. 3, at wavelengths shortward of 1906 Å we include the Enceladus reflectance spectrum from Cassini UVIS taken by Hendrix et al. (2010). We note that these reflectance spectra are given in terms of the canonical I/F with no correction for phase angle; they are equivalent to the product of the geometric albedo times the phase function. In Fig. 4, we overplot a portion of a spectrum from observations from HST Space Telescope Imaging Spectrograph (STIS) carried out over 1999–2000 by Keith Noll. Our phase
angle was 4.3 degrees, while those of Hendrix et al. (2010) and Noll were both 2 degrees. However, we do not expect this difference to factor into our discussion of the broad spectral features. The trailing hemisphere of Enceladus is known to be slightly brighter than the leading hemisphere by 0.06–0.12 mag (approximately 6–12%), both in the visible (Verbiscer et al., 2005) and the near-IR (Buratti et al., 1998; Momary et al., 2000; Verbiscer et al., 2005, 2006). We see this effect in the B stripe of the NUV channel (2800–3200 Å); however, in the A stripe of the NUV channel (1906–2130 Å), we find no significant difference in the continuum reflectivity of the two hemispheres shortward of 2105 Å (Fig. 5). Possible mechanisms for any difference in albedo between hemispheres include micrometeoroid impacts, magnetospheric plasma bombardment, and deposition of E-ring material (Hamilton and Burns, 1994; Buratti et al., 1998). With respect to the possibility of magnetospheric plasma, Enceladus’ trailing hemisphere is preferentially bombarded due to the fact that Saturn’s magnetospheric plasma corotates faster than Enceladus at its orbital distance. In particular, Sack et al. (1992) found that bombardment of energetic ions tended to brighten the spectrum in the visible but darken it in the FUV. This process would appear to be consistent with our nondetection of a brighter trailing hemisphere in the A stripe. Another possibility is that the spectral properties of the local material deposited by the E-ring and its source, the plume, could account for the lack of hemispheric albedo difference in the FUV. Schenk et al. (2011) noted that the plume and E-ring deposition, rather than external bombardment, appeared to dominate spectral color patterns on the surface of Enceladus and block color features seen on other saturnian satellites. If plume and E-ring deposition is the dominant bombardment mechanism in determining the modification of the surface reflectance of Enceladus, the lack of a hemispheric difference in albedo in the FUV could suggest that the plume/E-ring is depositing material fast enough to counteract the effects of micrometeoroid and plasma bombardment in this spectral region.
5. Discussion and conclusions 5.1. Airglow emission The upper limits on the airglow emission of Enceladus are quite low. Even assuming a small emitting region roughly the size of Enceladus’ plume, the upper limits are in the hundreds to thousands of Rayleighs, on the same order as the faint auroral footprints reported by Pryor et al. (2011). Similarly, the column densities for C
M. Zastrow et al. / Icarus 220 (2012) 29–35
33
Fig. 5. Comparisons of spectra of leading and trailing hemispheres from COS, smoothed to 25 pixels (2 Å) boxcar average. On the left, in the mid-UV, the trailing hemisphere can be seen to be brighter than the leading hemisphere. On the right, in the FUV, this effect is not seen.
are low—well below the H2O upper limit of 1016 cm2 from UVIS reported by Hansen et al. (2011). Although a plume-sized source for C I 1280 results in an upper limit for C of 3.1 1015 cm2, or roughly one-third the density of the H2 O plume, C I 1277 constrains it to 5 1012 cm2 or one part in 2000. (For comparison, Hansen et al. (2008) placed an upper limit on CO of 0.3% of the plume, or a column density of 3 1014 cm2.) While the faintness of the footprints might suggest that the bulk of the interaction could be found near Enceladus itself, we find no evidence for significant airglow/auroral emissions in the extended atmosphere of Enceladus. For future work, O 1356 Å presents an opportunity to constrain the interaction between Enceladus and Saturn’s magnetosphere. Since it is a forbidden transition, one could reasonably take its upper limit of emission to be a limit of collisional excitation. The peak UV emission cross section for electron impact excitation on O 1356 Å is 9.0 1018 cm2 at 16 eV, and for e on O2 the cross section is 8.8 1018 cm2 at 110 eV (Meier, 1991). Therefore, with knowledge of the energy distribution of the incoming electrons, a theoretical upper limit on O column density could be obtained. The electron temperature at Enceladus is quite cold—it has been measured by CAPS to be 1 eV (Sittler et al., 2008)— and the 10 eV electrons necessary to excite O 1356 would clearly be in the tail of the distribution. However, published Cassini results do not currently reveal the full energy distribution of electrons at Enceladus. Modeling would be necessary to extrapolate downward from higher energy observations, which is beyond the scope of this paper. 5.2. Reflectance spectrum Hendrix et al. (2010) noted a region at 170–190 nm that was too dark to be explained by a pure water ice spectrum. This indicated the presence of another species that was bright in the visible, but dark in the FUV. Hendrix et al. (2010) singled out an intimate mixture of water ice and ammonia ice with global coverage of Enceladus as the most promising explanation, as ammonia ice has a strong absorption edge at 190 nm. Hendrix et al. found good agreement with a model spectrum of a mixture of 1% NH3 in H2O. We find that our data is consistent with the H2O + NH3 models reported in the model spectra of Hendrix et al. (2010). Our data
suggests that the spectrum begins to quickly brighten longward of 1900 Å, constraining the extent of this absorption feature. (The signal-to-noise ratio from COS and STIS is too low to be reliable shortward of 1906 Å.) This matches the overall behavior of the H2O + NH3 models of Hendrix et al. (2010, Fig. 9), which exhibit a sharp edge longward of 1900 Å. The model that best fits our data in the FUV is the one using the Dressler and Schnepp (1960) ammonia data as rescaled and included in the compilation by Martonchik et al. (1984), also with a grain size of 30 lm (Hendrix et al., 2010, Fig. 9, top left). Other models by Hendrix et al. (2010) included different NH3 measurements, including optical constants by Dawes et al. (2007) and reflectance data by Hapke et al. (1981) and Pipes et al. (1974). Hendrix et al. kept the ratio of water ice to NH3 constant, and so the only differences in the models arise from the inconsistencies observed in the different lab spectra of NH3. All of them exhibit a steep absorption edge, but its location varies from 1900 to 2100 Å. In addition, a ‘‘trough’’ in reflectance (an absorption peak) is seen just shortward of the absorption edge in the optical constants measured by Dawes et al. (2007) at Enceladus-like temperatures (75 K). It also appears more weakly in the measurements by Dressler and Schnepp (1960) in the Martonchik et al. collection at 77 K (although not at 125 K) (Hendrix et al., 2010, Fig. 5). It is not seen in the reflectance measurements by Hapke et al. (1981) and Pipes et al. (1974). In the model spectra, this feature appears most strongly in the model using the Dawes et al. measurements, where the ammonia ice grain size is increased to 30 lm to darken the spectrum to match the 170–190 nm absorption band, (Hendrix et al., 2010, Fig. 9). To these models, Hendrix et al. (2010) (Fig. 9, bottom) also added 1% tholins in order to darken the NUV longward of 190 nm to match the data point at 275 nm of Verbiscer et al. (2005), varying the grain size—the larger the grain, the greater the absorption. This darkened the NUV, while making no significant difference in the wavelength range discussed in the preceding paragraph. In Fig. 6, we present a model spectrum that combines water ice with 1% NH3, and 1% tholins, but using the Martonchik et al. NH3 constants. The difference in the location and strength of the ammonia trough produces an acceptable fit to our COS and STIS data in the FUV while still darkening the NUV to match the STIS data out to 2800 Å.
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M. Zastrow et al. / Icarus 220 (2012) 29–35
Fig. 6. Composite spectrum from Fig. 4, but with the COS data smoothed as in Fig. 3. The model using Martonchik et al. (1984) optical constants is overplotted.
Hendrix et al. noted that the dearth of published optical constants for ices and tholins at Enceladus temperatures in this spectral range hampered their analysis. Additionally, if either water or ammonia ice are actually slightly absorbing in the NUV, tholins may not be necessary to match the data of Verbiscer et al. (2005). Nor do we have any way to know which set of NH3 optical constants (Dawes et al. or Martonchik et al.) are more accurate. Nevertheless, we present Fig. 6 to demonstrate that our data is consistent with the range of H2O + NH3 + tholins model spectra of Hendrix et al. in the near- to mid-UV. We note one discrepant feature in the model spectrum (Fig. 6) around 190–195 nm, where the ammonia ice trough darkens the model relative to the UVIS, COS, and STIS spectra. One possible explanation is that, as noted by Hendrix et al. (2010), many published lab spectra of water ice also show a trough in reflectance varying in strength from measurement to measurement—similar to that discussed in the previous paragraph in ammonia ice, but at a slightly shorter wavelength: ranging from 1860 to 1930 Å as measured by Hapke et al. (1981); Pipes et al. (1974), and Gougen (obtained by Hendrix et al. via personal communication). However, this feature was not incorporated into their models as it was not present in the dataset they based them on, the optical constants of Warren and Brandt (2008). If it were, it might even out the model spectrum over the 170–200 nm range by darkening the shortwavelength side of it. Coupled with a slightly smaller ammonia grain size (to slightly brighten it to fit to the data), it could reduce the discrepant feature or eliminate it altogether. Further analysis is hampered by the difficulty in precisely measuring the spectra of ices in the lab, the resulting inconsistencies of the existing ice data, and the general lack of data on ices at Enceladus temperatures. Keeping in mind the varying characteristics of the different ice measurements, we observe that our results are in agreement with the H2O + NH3 + tholins model spectra produced by Hendrix et al. (2010), and support the possibility of the presence of small amounts of ammonia ice and tholins. Acknowledgments This work is based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the AURA Inc. for NASA. This research was supported by NASA Grant HST-GO-11645.01-A from the Space Telescope Science Institute to Boston University. Marty Snow kindly provided the solar spectrum from SORCE SOLSTICE. MZ thanks Andrew West and Supriya Chakrabarti for helpful conversations.
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