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Mercury's exosphere: origin of surface sputtering and implications
Planet. Space Sci., Vol. 45, No. 1, pp. 73-79, 1997 0 1997 Elsevier Science Ltd Printed in Great Britain. All rights reserved 0032-0633/97 $17.00+0.00
Pergamon
PII: SOO32-0633(96)00097-9
Mercury’s exosphere
:
origin of surface sputtering and implications
H. Lammer’ and S. J. Bauer”* ‘Dept. of Extraterrestrial Physics, Space Research Institute, Austrian Academy Austria ‘Institute for Meteorology and Geophysics, University of Graz, Graz, Austria Received
28 September
1995; revised 18 April 1996; accepted
to: H. Lammer
1, A-8010 Graz,
13 June 1996
Abstwct. The existence of gaseous’ species II, He, Q, Na and Kl established by cuv spectroseapie observatians on Mariner 10 for the $irst three a-nd groundbased observations for the last: two species vvith column contenti of the order of I@‘2cm~2 or less, quali& the : ga+wus: envelope QfNIercury ,as an exvsphere (column content per defmition I W cm-“): Whereas II a&l He seem to have its origin in the direct supply from the solar wind, the heavier constituents can arise from particle and photon interactions with surface matetiafs. Some observations have suggested that Na extends to ,7W km above ~tie surface at &e subso&r point, imply’ i3i.g the existetice of non-thermal atomic velbcities’of at leas’i.2 kmsi’, Al*ouglX Mer@ry has a small intrinsic na_agneticfield, solar wind protons have access-to parts of its surface. Since. particle sputtering liberat& wplatiles from the surface at energies greater than t&e mean therm&l energy, the sputtered particle profile. is governed by a scale height much huger thanthe bar+ m&tic sex& $.&glit. The surface den&ies of fhese sputWed vofatiles t&en exhibit a lower expotintial. do&y’ t&m wouJid’be e-xpe@edkm a: bttrcrme& law using the mxq&x:ic ‘(sutiacej terrtpetatm@‘&&Ecuq”s emsphere is.a low attenuation medium for solar euv r&i. ation whj& therefore can penetrate to ..its surface. Under these circumstances no ionospheric “layer’” can be formed. The ionization of exosp&ric species bads to electron-ion pair format&n VuQichis balanced by chemical loss and transport processes, S&X radiative recombination is an extremel;y slow loss process for atomic ions, plasma di&&st-onunder g&vityV which has a short time constant, controls the’ &I.%& species. Thus, the densities of the ionized exospheric component do not exceed a few electronsljuns cm-3. This yields an extremely low height-integrated Pcdersen conductivity, which must be taken into account in the understanding of convection processes in the magnetosphere. 0 1997 Elsevicr Science Ltd. AB rights reserved Correspondence
of Sciences, Halbarthgasse
Introduction
The Mariner 10 spacecraft first encountered Mercury on 29 March 1974. Before the Mariner 10 fly-by, it was generally thought that the solar wind interaction with the planet was moonlike, in which case solar wind ions were expected to impinge directly on the surface (Hartle et al., 1973; Hodges, 1974; Hodges et al., 1974). In these preencounter models, the solar wind was the source for the exosphere, and its magnitude was determined by the solar wind flux intercepted by Mercury. The Mariner 10 spacecraft discovered that direct solar wind interaction was generally not possible, since the planet was found to possess a magnetosphere (Ness et al., 1974, 1976; Ogilvie et al., 1974). However, as with the Earth, a fraction of the solar wind ions penetrate into the magnetosphere and, down to the surface and thus provide an indirect source for the exosphere. The airglow experiment on board of Mariner 10 indeed identified atomic hydrogen and helium as constituents of the exosphere with a He surface number density of about 4500 cmp3 estimated by Broadfoot et al. (1976). Curtis and Hartle (1978) estimated the magnetospheric capture fraction of solar wind He*+ required for maintaining the observed He exosphere. Broadfoot et al. (1976), from Mariner 10 data, also discovered oxygen in the Mercury exosphere. Potter and Morgan (1985), from ground-based observations, discovered in the spectrum of Mercury strong emission features at the Fraunhofer sodium D lines attributed to resonant scattering of sunlight from sodium vapour in the exosphere of the planet. They estimated a Na column density of 1011-1012cmp2, which corresponds to a surface number density of 2 x 104-1.5 x 10’ Na atoms cmp3. Potter and Morgan (1986) later also discovered small concentrations of potassium in the Mercury exosphere and suggested that solar wind ion sputtering might account for their existence. Atoms or molecules can be supplied from the surface to the exosphere by thermal evaporation, or from sputtering by charged particles, or by photons. Sputtering of the surface by bombardment of solar wind protons, alpha particles and heavy ions in the aurora1 zones was first
74 discussed by McGrath et al. (1986) by Cheng et al. (1987) and by Killen (1989), but McGrath et al. (1986) suggested that photon stimulated desorption (photon sputtering) be more likely than charged particle sputtering and estimated the photon sputtering fluxes to be of the order of 2 x IO’2 x 10’ cmp2 s-l. Killen and Morgan (1993) however, argued that these fluxes are overly optimistic since they were based on data for alkali halides. The photon sputtering yield depends critically on the low energy cut-off of photons capable of ejecting sodium. Sputtering can be produced only by light with energy larger than the band gap (Townsend and Elliot, 1970) ; Killen et al. (1990) therefore assumed that photons beyond the band gap for NaCl, i.e. IOeV, can sputter Na atoms and obtained an average photon sputtering flux of about 2 x 1O’Na atoms cmp2 s-r, assuming a lunar abundance of sodium. Shemansky and Morgan (1991), on the other hand, estimated an Na photon sputtering flux of only 3.8 x 104cm-2s-1. Thus, there seems to be a discrepancy between the estimated production rates of exospheric Na by photon sputtering of nearly “four orders” of magnitude. Cheng et al. (1987), Morgan et al. (1988, 1989), Sprague (1990) and Cintala (1992) have considered Na and K supply by micrometeorite impact vaporization. Since the rate of volatile production is proportional to the elemental abundance in the regolith, Na and K should be present in the vapour in the same proportion as they exist in the target material. The production of Na by micrometeorite impacts is, however, typically only a few per cent that required for the observed atom content (Cintala, 1992; Killen and Morgan, 1993). Observations of the broad wings in the profile of the sodium D-line, however, suggest that Na extends up to about 700 km above the surface (e.g. Potter and Morgan, 1985; Smyth, 1986; Hunten et al., 1988). This requires the presence of hot Na atoms to match the broad wings of the observed line profile. Neither photon sputtering nor thermal evaporation produce a non-thermal velocity distribution enabling particles to reach such altitudes. Impact vaporization of interplanetary meteoroids or charged particle sputtering could be a source of hot Na atoms with ejection speeds > 2 km s-l (Morgan et al., 1988 ; Ip, 1990,1993). Killen and Morgan (1993) used the estimated solar wind proton fluxes given by Baker et al. (1987) and concluded that ion sputtering is five times smaller than that required to supply the exospheric content, assuming a lunar abundance of sodium in the regolith. If the Na abundance on Mercury were earthlike, which is five times the lunar value, these ion fluxes, however, would be sufficient to supply the exosphere from Mercury’s crust. The Na and 0 column densities were estimated in earlier papers (e.g. Broadfoot et al., 1976; Potter and Morgan, 1985; McGrath et al., 1986; Cheng et al., 1987; Killen et al., 1990 ; Killen and Morgan, 1993) on the basis of scale heights of about 50 and 80 km, respectively, corresponding to a subsolar surface exospheric temperature of 500-575 K.
H. Lammer and S. J. Bauer: Mercury’s exosphere sity profiles of the newborne hot atoms by comparing them with the “cold” density distributions. Ionization processes are reviewed and the equilibrium concentrations of exospheric species are evaluated from considerations about production, chemical loss and plasma diffusion under gravity. Magnetospheric ions, solar wind ions, and locally produced pick up ions can impact the atmospheres or surfaces of bodies in the solar system. These ions transfer their energy by collisions with atmospheric or surface atoms or molecules, they induce an expansion of the atmospheric corona, and are responsible for the loss of a fraction of the hot particles from the atmosphere (e.g. Haff and Watson, 1979 ; Johnson, 1990 ; Luhmann et al., 1992 ; Lammer and Bauer, 1993; Johnson, 1994; Lammer, 1995). Sputtering from a solid planetary surface is calculated in exact analogy with sputtering from an atmosphere ; it produces particles which escape directly or, if their velocity is insufficient, enter ballistic trajectories. We use the energy distribution F(E,) of Sieveka and Johnson (1984) for the determination of the ejection energy E, :
fro=(E, +EeE,J3 {&
yi”“l”‘}
(1)
where E,, is the surface binding energy and Ei the energy of the incident particle. The functions represented by equation (1) have been normalized in such a way that : $ F(EJ dE, = 1.
(2)
Although the surface composition of Mercury is not well known, Na is likely to be bound to 0 with a binding energy of about 2eV, as for example in NaA1,Si30s (feldspar) (McGrath et aE., 1986 ; Cheng et al., 1987). Since 0 atoms themselves are tightly bound to the molecule, the overall binding energy Eb is about 3-4eV (Johnson, 1995). We consider that the incident solar wind protons have an energy E, of 1 keV. Figure 1 shows the shape of the energy distribution F(E,) as a function of the ejected particle energy E,. It is seen that most Na and 0 atoms are ejected with energies of about 1 and 1.7 eV, respectively. These results are in good agreement with the estimations of Ip (1986, 1993). The neutral species injected in Mercury’s exosphere move under the influence of gravity and solar radiation pressure, but collisions between exospheric particles are neglected (Smyth, 1986 ; Ip, 1990). Most atoms return to the surface after a time of flight t,, ; for a particle leaving the surface of a planetary body with a speed v, ejection energy E, and direction 0 with respect to the surface normal, we obtain (Johnson, 1990) : R
to =
max
dr (3)
is (dr/dt)
with Ejection of Na and 0 atoms by charged particle sputtering We now estimate the ejection energies of Na and 0 atoms by charged particle sputtering, as well as the number den-
$=U[l-(~‘)lsin2@-(~)(l-~)~2 where
R, is the surface radius
(4) and
u” the gravitational
15
H. Lammer and S. J. Bauer: Mercury’s exosphere
0.15.....-_...../
0.10
0.05
1 .o
0.1
10.0
E. [evl Fig. 1. Behavior of the normalized energy distribution IF(&) as a function of ejected particle energies E,.Most Na and 0 atoms are ejected into Mercury’s exosphere with about 1 or 1.7 eV. These results are in good agreement with the estimated ejection energies of Ip (1986, 1990, 1993)
(Hunten et al., 1988) and that, the critical level on Mercury is the planetary surface. At great distances above the critical level the barometric law breaks down, since a planetary exosphere populated by escaping particles is not strictly in hydrostatic equilibrium. However, Hodges (1980) has shown that on the Moon and on Mercury, elastic collisions of He atoms with a Maxwellian distribution of vibrating bound atoms produce a “nearly” Maxwellian distribution of velocities, despite the absence of He atoms impinging the surface with velocities in excess of the escape velocity. The exosphere is expanding slightly, and matter is lost, if the gravity field is small, which corresponds to evaporative loss in the kinetic theory. Therefore, the number density n(z) is the product of the barometric density and a partition function c (e.g. Chamberlain, 1963 ; Bauer, 1973; Lammer, 1989; Lammer and Bauer, 1991)
n(z) = n, epcziH)c(x)
(6)
n, is the surface density and x the escape parameter binding energy of the planet. The free flight time to is typically < 2000 s for Na atoms. A fraction of the returning particles penetrate into the surface of Mercury and are recycled into the system ; these implanted ions can cause additional neutral sputtering. Neutrals are lost from the exosphere when they are ionized and swept away under the influence of magnetospheric electromagnetic fields ; due to the eccentricity of Mercury’s orbit, the photoionization time lies in the range 103-104s. One must therefore consider an equilibrium between sources and sinks for each exospheric species. In order to maintain the exosphere in a steady state, the Na supply rate must of course equal the loss rate. If all the ions are escaping, the Na supply rate must be w lo7 atoms cm-* s- ’ (Killen and Morgan, 1993). The maintenance of the present-day exosphere over geologic time requires that a production mechanism can supply Na from great depths and that fresh Na bearing materials be brought to the surface. There is still a large uncertainty about the processes supplying Na or K to the exosphere. However, all processes which bring Na into the exosphere are dependent on the composition of the regolith, megaregolith and upper crust of the planet (Killen and Morgan, 1993).
x(4
GmM = -&-
:
(7)
G is the gravitational constant, M the planetary mass, mass, k the Boltzmann constant, r the planetocentric distance and T the exospheric temperature. The partition function can be written as
m the particle
i: =
(f-9
rbal+La,+ L,.
The first term corresponds jectories :
to particles
with ballistic
tra-
[( 1 s ( )I
i,a~w=$2 Y$X
-7
6s -3
e-Yy 3, X_-y 2
(9)
where
y=xz x+x,
(10)
x, is the escape parameter on Mercury’s y(3/2,x - y) the incomplete I-function :
surface
and
(11)
Mercury’s exosphere The detailed theory of the exosphere made by Chamberlain (1963) is valid for a Maxwellian distribution. The critical level or height, above which collisions are negligible, is called exobase and the number density at this level can be defined as (e.g. Chamberlain, 1963) :
The second orbits :
term
L(x) =
corresponds
$
to particles
in satellite
(12)
“;:“‘12ewy[y(~,x-y)]
1
n, = aH where u is the the scale height. is thus usually Spacecraft and that Mercury’s
gas-kinetic collision cross-section and H The total column content of an exosphere N, d c-’ z 2 x 10’4cm-2 (Bauer, 1973). ground-based observations determined total atmospheric content is < 10’2cm-2
is y
=XTYy3/*-l e-Y 0
dy,
(13)
76
H. Lammer and S. J. Bauer: Mercury’s exosphere
The third term corresponds to the tail of the Maxwellian distribution, i.e. particles with escaping trajectories :
- (xi~~“2e-~[I(~)
-8(:,x-y)]}.
(14)
The Mariner 10 euv airglow spectrometer measurements (Broadfoot et al., 1974) indicated very low values of surface densities y1,, of about 8 cmp3 for H atoms and about 4.5 x 103-6 x lo3 cmW3, for He atoms on the dayside of Mercury (Hunten et al., 1988). A scale height H was determined only for He ; the detection of 0 was only tentative. The He scale height corresponded to a temperature of 575K, roughly consistent with that of the subsolar region. The Na and K observations were made with ground-based optical spectrometers. Morgan et al. (1988) revised the originally reported Na column density of about 8 x 10” cm-’ (Potter and Morgan, 1985) downward to about l-2.5 x lOr* cnl-*, whereas Killen et al. (1990) estimated Na column densities of about 2.83.8 x 10” cmp2. However, these values were derived for a scale height of about 50 km, corresponding to a subsolar surface exospheric temperature of 500-575 K and under the assumption that photon sputtering would be the dominant source of Na atoms. If the population of hot Na atoms originate mostly from charged particle sputtering or meteoroid impact vaporization (Morgan et al., 1988 ; Ip, 1990, 1993) their initial speed must be considerably higher than the thermal velocities of about 0.6 km s-l. The non-thermal Na components have a broader velocity distribution with more particles exceeding the surface escape velocity of 4.3 km s-l. The motion of these hot atoms, can, however, be influenced by solar radiation pressure acceleration (Smyth, 1986; Morgan et al., 1988; Hunten et al., 1988; Ip, 1990), and reach escape trajectories if their velocities exceed 2 km s-l. Atoms with velocities less than 2 km s-l are also influenced by solar radiation and drift from the dayside to the night-side hemisphere (Ip, 1990). If one assumes an ion sputtering yield of Y = 0.1-2, a “lunar” abundance corresponding to a sodium fraction of 4 x low3 in the regolith and incident ion fluxes equal to those given by Baker et al. (1987), one derives a sputtered Na flux in the range 5 x 105-1 x 10scm~*s~’ (Killen and Morgan, 1993). The surface number density of the nonthermal Na component is estimated by dividing the mean velocity of the energy distribution function of the ejected Na atoms, about 2.9 x lO”cms-’ at 1 eV, by Na fluxes given by Killen and Morgan (1993). The surface number density of the hot component ranges between 2 and 345cmW3. If a “terrestrial” abundance of Na (i.e. five times that of the Moon) is considered, then an upper limit for the corresponding Na fluxes is about 2.5 x 1065 x 108cm-2s-’ (Killen and Morgan, 1993) and is sufficient to supply the exospheric content. The surface number densities corresponding to a “terrestrial” Na abundance in Mercury’s regolith lies between 9 and 1730 cme3. Figure 2 shows the number density profiles based on “lunar” abundance for 0, Na and hot Na atoms when
10
1000
100 n
10000
[cm-‘]
Fig. 2. Number density profiles of He, 0, Na and hot Na atoms for “lunar” surface abundance. Both charged particle sputtering and photo-sputtering are considered as possible source mechanisms for Na; particle sputtering is the source of the nonthermal component. Helium dominates over sodium above about 140 km and over oxygen (bound particles) above about 240 km. For helium a full line represents the bound particles with ballistic trajectories and satellite orbits, whereas a dashed line represents the sum of bound and escaping particles. The dot-dashed line (bound particles) shows the number density distribution of hot Na atoms estimated from a sputtering flux of about 1 x 108cm-2s-’ (Killen and Morgan, 1993). The hot Na component extending beyond 700 km could be responsible for the observations of the broad wings in the line profile of the sodium D-line
both charged particle sputtering and photo-sputtering are considered as possible source mechanisms for the Na population, including a non-thermal component. The He atoms dominate over the Na atoms above 140 km and the 0 (dotted line) atoms above about 240 km. The full line represents the density of bound particles with ballistic trajectories and satellite orbits, whereas the dashed line represents the sum of bound and escaping particles (equation (14)). The dot-dashed line gives the number density profile of hot Na atoms, as derived from an upper limit of the sputtered flux of about 1 x lo* cm-2s-i (Killen and Morgan, 1993). The hot Na component extends to altitudes beyond 700 km and could explain the observations of the broad wings in the line profile of the sodium Dline. These hot atoms, however, are influenced by solar radiation pressure acceleration. A small population of high velocity 0 atoms could also be present, but their concentration may lie below the observational limit because losses due to radiation pressure acceleration may become significant. Charged particle sputtering, however, yields a hot Na corona, similar to the N2 corona of Saturn’s large moon Titan (Lammer and Bauer, 1993) and Neptune’s moon Triton (Lammer, 1995).
Mercury’s ionized component The radio occultation experiment on Mariner 10 set an upper limit for the ionospheric density of about lo3 cmW3. As pointed out by Bauer (1975), the absence of an iono-
77
H. Lammer and S. J. Bauer: Mercury’s exosphere spheric layer is consistent with ground and spacecraft optical observations of a total column content N, of ionizable species of d 10” cm-‘. An ionospheric layer cannot form since an optical depth z = 1 is required, corresponding to a column content of absorbing, ionizable constituent NT = ((r,)-‘. Since, typical absorption crosssections B, are of the order z lo-” cm*, z z 10m6and the ionizing radiation penetrates nearly unattenuated down to the planetary surface. The ionization of a species X by the action of a photon with energy hv>IP (ionization potential) is written X = hv(TIP)+X++e.
(15)
Since, the optical depth of Mercury’s exosphere is very small, the ion production function takes the simple form : q = Jn(z)
(16)
where the photoionization rate coefficient J, is about 1.5 x 10-4s-’ for Na and 4 x 10-7s-’ for He (Cheng et al., 1987). The number density distribution of electrons and ions in a planetary “atmosphere” is governed by the equation of continuity (e.g. Bauer, 1973) : (17) where q is the ion pair production, ziDthe plasma diffusion velocity and L(Ni) the chemical loss. If we assume a steady state, two limiting cases can be considered for the continuity equation : (1) “chemical equilibrium”, when chemical processes are predominant and (2) “diffusive equilibrium” when chemical processes can be neglected compared to plasma diffusion. The applicability of either of these to Mercury can be determined by estimating the appropriate time constants. The plasma diffusion time zn is : H2 TD E D,
v-9
where D, = D,/n is the ambipolar diffusion coefficient with Do z 1 x 10’gcm-l s-’ (e.g. Bauer, 1973). Using values of 12and H, representative of Mercury’s exosphere, it is found that the diffusion time for Na ions zn < 1 s. The chemical loss of atomic ions, in the absence of neutral molecular species, is determined by radiative recombination with a coefficient a of the order of lo-‘*cm3 s-’ (Banks and Kockarts, 1973) and the chemical lifetime is rc=p.
1 aNi
(19)
If photochemical equilibrium prevails, Ni can be as high as 4 x 104cm-3. Since chemical lifetime the zc x 2.5 x lo7 s, one can see that recombination is extremely slow compared to plasma diffusion even for high values of Ni. Thus, neglecting the chemical loss term, L(Ni), the continuity equation reduces to : q= Jn=$(Nio)
where J is the photoionization
rate coefficient and n the
number density of the neutral gas. With vD z Do/(d) ion density Ni becomes : Jn2H2 Nip:--.
the
(21)
The resulting ion surface densities for Na, 0 and He ions are < 1 cmw3. Thus, the ionized component of Mercury’s exosphere will not exceed densities of a few electrons/ions cmm3. These results are in a good agreement with Mariner 10 observations of electron densities of about 0.1 cmM3 in the polar cap region and 1 crn3 in Mercury’s magnetotail, although Ogilvie et al. (1977) suggest that the observed electrons within the polar flux tubes at low altitudes are suprathermal electrons and not of atmospheric origin. For understanding convection processes in Mercury’s magnetosphere, one must first know the height-integrated Pedersen conductivity, since it controls the rate of field line merging and hence the ionospheric electric field ; the Pedersen conductivity is written as
cp=e2[& (v;,i$J]
(22)
with ion mass mi, gyro frequency ooY and neutral-ion collision frequency (Banks and Kockarts, 1973), i.e. : v,,~= 2.6 x lo-‘n
0
I/*
2
S-l
(23)
PA
where a, is the polarizability in units of 10-24cm3 of the neutral gas with density n and :
pA =
(
m,mi
(m,+m,)
)
(24)
is the reduced mass in amu. The electron/ion densities encountered on Mercury, lead to an extremely low heightintegrated Pedersen conductivity (up to 1000 km) of about 5 x 10-6mho. Reflectance spectra from Mercury’s surface display the same slope as those from the Moon, a feature which is attributed to the presence of Fe- and Ti-bearing agglutinates in the lunar regolith (Vilas, 1988). The conductivity of the lunar crust is about 10-7mhom-1 (Schwerer et al., 1972 ; Goldstein, 1974) and if one assumes that this figure is also typical of Mercury, then the height integrated Pedersen conductivity is shown to be equivalent to the conductivity of a crustal layer 50 m in thickness. A conducting ionosphere hinders magnetospheric convection because of the drag force exerted by ion-neutral collisions. The ionosphere shorts out the convection electric field induced by solar wind, this effect is known as “ionospheric line tying” and has been discussed in connection with the Earth’s magnetosphere by Cole (1963) and Atkinson (1967). Rassbach et al. (1974) suggested that the maximum possible convection potential is one that drives ionospheric currents sufficiently large to perturb significantly the magnetic field near the dayside magnetopause. However, ionospheric line tying plays an important role only when the height-integrated conductivity is comparable to or greater than a critical conductivity g’cof about 20mho (Hill et al., 1976). The lack
78
H. Lammer and S. J. Bauer: Mercury’s
of ionospheric line tying might mean a very short magnetospheric convection timescale in response to the solar wind condition and hence the particle acceleration process. Since both the derived height-integrated Pedersen conductivity and the estimated surface conductivity are negligibly small compared to the critical conductivity o, of about 20 mho, ionospheric line tying must be virtually absent in Mercury’s magnetosphere. Magnetospheric processes should therefore not be strongly affected by ionospheric conductivity at Mercury as they are at Earth and possibly even at Mars (Rassbach et al., 1974; Hill et al., 1976).
Conclusion Our results are given for an exosphere composed of He, Na and 0 atoms with surface number densities of about 6 x 103, 3.8 x 104, and 4.4 x 104cme3, respectively, which is consistent with observed column contents and assuming lunar abundance for Mercury’s surface constituents. We show that most Na and 0 atoms originating via charged particle sputtering are ejected into Mercury’s exosphere with energies of about 1 and 1.7 eV, respectively, in good agreement with the estimated ejection energies of Ip (1986, 1990, 1993). A hot Na component extending to altitudes beyond 700 km could be responsible for the observations of broad wings in the line profile of the sodium D-line. The density distribution of the hot Na atoms depends on the incident solar wind flux, regolith composition (lunar or terrestrial) and estimated sputtering rates. A small population of high velocity 0
a sub-
exosphere
Mercury, Venus, and Mars, ed. N. F. Ness, NASA SP-397, pp. 47-62. Broadfoot, A. L., Kumar, S., Belton, M. J. S. and McElroy, M. B. (1974) Mercury’s atmosphere from Mariner 10 : preliminary results. Science 185, 1666169. Broadfoot, A. L., Shemansky, D. E. and Kumar, S. (1976) Mariner 10: Mercury atmosphere. Geophys. Res. Left. 3, 577-580. Chamberlain, J. W. (1963) Planetary coronae and atmospheric evaporation. Planet. Space Sci. 11,901-960. Cheng, A. F., Johnson, R. E., Krimigis, S. M. and Lanzerotti, L. J. (1987) Magnetosphere, exosphere and surface of Mercury. Icarus 7L430-140. Cintala, M. J. (1992) Impact induced thermal effects in the lunar and Mercurian regoliths. Geophys. Res. Lett. 97,207-208. Cole, K. D. (1963) Damping of magnetospheric motions by the ionosphere. J. Geophys. Res. 68, 3231-3235. Curtis, S. A. and Hartle, R. E. (1978) Mercury’s helium exosphere after Mariner 10’s third encounter. J. Geophys. Res. 83,1551-1557. GoIdstein, B. E. (1974) Observations of electrons at the Lunar surface. J. Geophys. Res. 79,23-35. Half, P. K. and Watson, C. C. (1979) The erosion of planetary and satellite atmospheres by energetic atomic particles. J. Geophys. Res. 84, 84368442. Hartle, R. E., Ogilvie, K. W. and Wu, C. S. (1973) Neutral and ion exospheres in the solar wind with applications to Mercury. Planet. Space Sci. 21,2181-2191.
Hill, T. W., Dessler, A. J. and Wolf, R. A. (1976) Mercury and Mars : the role of ionospheric conductivity in the acceleration of magnetospheric particles. Geophys. Res. Lett. 3,429432. Hodges Jr, R. R. (1974) Model atmospheres for Mercury based on a lunar analogy. J. Geophys. Res. 79,2881-2885. Hodges Jr, R. R. (1980) Methods for Monte Carlo simulation of the exosphere of the Moon. J. Geophys. Res. 85,164170.
Hodges Jr, R. R., Hoffman, J. H. and Johnson, F. S. (1974) The lunar atmosphere. fcarus 21,415426. Hunten, D. M., Morgan, T. H. and Shemansky, D. E. (1988) The Mercury atmosphere. In Mercury, pp. 562-613. University of Arizona Press, Tucson. Ip, W.-H. (1986) The sodium exosphere and magnetosphere Mercury. Geophys. Res. Lett. 13,423426.
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Ip, W.-H. (1990) On solar radiation-driven surface transport of sodium atoms at Mercury. Astrophys. J. 356,675-68 1.
A diffusion a negligibly
The authors wish to thank R. E. Johnson (Department of Nuclear Engineering and Engineering Physics), University of Virginia, U.S.A., for discussions relating to this work and two referees for their helpful comments.
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Ip, W.-H. (1993) On the surface sputtering effects of magnetospheric charged particles at Mercury. Astrophys. J. 418, 451456. Johnson, R. E. (1990) Energetic Charged-particle Interactions with Atmospheres and Surfaces. Springer, Berlin.
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Johnson, R. E. (1995) Private communication. Killen, R. M. (1989) Crustal diffusion of gases out of Mercury and the Moon. Geophys. Res. Lett. 16, 171-174. Killen, R. M. and Morgan, T. H. (1993) Maintaining the Na atmosphere of Mercury. Icarus 101,294312. Killen, R. M., Potter, A. E. and Morgan, T. H. (1990) Spatial distribution of sodium vapor in the atmosphere of Mercury. Icarus 85, 145-167. Lammer, H. (1989) Planetary coronae and nonthermal escape from magnetic field free planets. M.Sc. thesis, University of Graz, Graz, Austria. Lammer, H. (1995) Mass loss of Nz molecules from Triton by magnetospheric plasma interaction. Planet. Space Sci. 43, 845-850. Lammer, H. and Bauer, S. J. (1991) Nonthermal atmospheric escape from Mars and Titan. J. Geophys. Res. 96, 18191825. Lammer, H. and Bauer, S. J. (1993) Atmospheric mass loss from Titan by sputtering. Planet. Space Sci. 41, 657-663.
H. Lammer
and S. J. Bauer: Mercury’s
exosphere
Luhmann, J. G., Johnson, R. E. and Zhang, M. H. G. (1992) Evolutionary impact of sputtering of the Martian atmosphere by O+ pickup ions. Geophys. Res. Lett. 19,2151-2154. McGrath, M. A., Johnson, R. E. and Lanzerotti, L. J. (1986) Sputtering of sodium on the planet Mercury. Nature 323, 694-696. Morgan, T. H., Hunten, D. M. and Shemansky, D. E. (1986) Atmosphere of Mercury. In Mercury. University of Arizona Press, Tucson. Morgan, T. H., Zook, H. A. and Potter, A. E. (1988) Impactdriven supply of sodium to the atmosphere of Mercury. Icarus 75, 156170. Morgan, T. H., Zook, H. A. and Potter, A. E. (1989) Production of sodium vapor from exposed regolith in the inner solar system. Proc. Lunar Planet. Sci. Con. 19,297-304. Ness, N. F., Behannon, K. W., Lepping, R. P., Whang, Y. C. and Schatten, K. H. (1974) Magnetic field observations near Mercury : preliminary results from Mariner 10. Science 185, 151-160. Ness, N. F., Behannon, K. W., Lepping, R. P. and Whang, Y. C. (1976) Observations of Mercury’s magnetic field. Icarus 28,479%488. Ogilvie, K. W., Scudder, J. D., Hartle, R. E., Siscoe, G. L., Bridge, H. S., Lazarus, A. J., Asbridge, J. R., Bame, S. J. and Yeates, C. M. (1974) Observations at Mercury encounter by the plasma science experiment on Mariner 10. Science 185, 146-1.52. Ogilvie, K. W., Scudder, J. D., Vasyliunas, V. M., Hartle, R. E. and Siscoe, G. L. (1977) Observations at the planet Mercury
79 by plasma electron experiment : Mariner 10. J. Geophys. Res. 82, 1807-1824. Potter, A. and Morgan, T. H. (1985) Discovery of sodium in the atmosphere of Mercury. Science 229, 651-653. Potter, A. and Morgan, T. H. (1986) Potassium in the atmosphere of Mercury. Icarus 67, 336-340. Rassbach, M. E., Wolf, R. A. and Daniel1 Jr, R. E. (1974) Convection in a Martian magnetosphere. J. Geophys. Res. 79, 1125-l 127. Schwerer, F. C., Huffman, G. P., Fisher, R. M. and Nagata, T. (1972) DC electrical conductivity of lunar surface rocks. The Moon 4,187. Shemansky, D. E. and Morgan, T. H. (1991) Source processes for the alkali metals in the atmosphere of Mercury. Geophys. Res. Lett. 18, 1659-1662. Sieveka, E. M. and Johnson, R. E. (1984) Ejection of atoms and molecules from IO by plasma-ion impact. Astrophys. J. 287, 418426. Smyth, W. H. (1986) Nature and variability of Mercury’s sodium atmosphere. Nature 323, 696-699. Sprague, A. L. (1990) A diffusion source for sodium and potassium in the atmospheres of Mercury and the Moon. Icarus 84,93-105. Townsend, P. D. and Elliot, D. J. (1970) Low energy electron and UV sputtering of KI. In Atomic Collision Phenomena in Solids, pp. 328-338. American Elsevier, New York. Vilas, F. (1988) Surface composition of Mercury from reflectance spectrophotometry. In Mercury, pp. 562613. University of Arizona Press, Tucson.
Pergamon
PII: SOO32-0633(96)00097-9
Mercury’s exosphere
:
origin of surface sputtering and implications
H. Lammer’ and S. J. Bauer”* ‘Dept. of Extraterrestrial Physics, Space Research Institute, Austrian Academy Austria ‘Institute for Meteorology and Geophysics, University of Graz, Graz, Austria Received
28 September
1995; revised 18 April 1996; accepted
to: H. Lammer
1, A-8010 Graz,
13 June 1996
Abstwct. The existence of gaseous’ species II, He, Q, Na and Kl established by cuv spectroseapie observatians on Mariner 10 for the $irst three a-nd groundbased observations for the last: two species vvith column contenti of the order of I@‘2cm~2 or less, quali& the : ga+wus: envelope QfNIercury ,as an exvsphere (column content per defmition I W cm-“): Whereas II a&l He seem to have its origin in the direct supply from the solar wind, the heavier constituents can arise from particle and photon interactions with surface matetiafs. Some observations have suggested that Na extends to ,7W km above ~tie surface at &e subso&r point, imply’ i3i.g the existetice of non-thermal atomic velbcities’of at leas’i.2 kmsi’, Al*ouglX Mer@ry has a small intrinsic na_agneticfield, solar wind protons have access-to parts of its surface. Since. particle sputtering liberat& wplatiles from the surface at energies greater than t&e mean therm&l energy, the sputtered particle profile. is governed by a scale height much huger thanthe bar+ m&tic sex& $.&glit. The surface den&ies of fhese sputWed vofatiles t&en exhibit a lower expotintial. do&y’ t&m wouJid’be e-xpe@edkm a: bttrcrme& law using the mxq&x:ic ‘(sutiacej terrtpetatm@‘&&Ecuq”s emsphere is.a low attenuation medium for solar euv r&i. ation whj& therefore can penetrate to ..its surface. Under these circumstances no ionospheric “layer’” can be formed. The ionization of exosp&ric species bads to electron-ion pair format&n VuQichis balanced by chemical loss and transport processes, S&X radiative recombination is an extremel;y slow loss process for atomic ions, plasma di&&st-onunder g&vityV which has a short time constant, controls the’ &I.%& species. Thus, the densities of the ionized exospheric component do not exceed a few electronsljuns cm-3. This yields an extremely low height-integrated Pcdersen conductivity, which must be taken into account in the understanding of convection processes in the magnetosphere. 0 1997 Elsevicr Science Ltd. AB rights reserved Correspondence
of Sciences, Halbarthgasse
Introduction
The Mariner 10 spacecraft first encountered Mercury on 29 March 1974. Before the Mariner 10 fly-by, it was generally thought that the solar wind interaction with the planet was moonlike, in which case solar wind ions were expected to impinge directly on the surface (Hartle et al., 1973; Hodges, 1974; Hodges et al., 1974). In these preencounter models, the solar wind was the source for the exosphere, and its magnitude was determined by the solar wind flux intercepted by Mercury. The Mariner 10 spacecraft discovered that direct solar wind interaction was generally not possible, since the planet was found to possess a magnetosphere (Ness et al., 1974, 1976; Ogilvie et al., 1974). However, as with the Earth, a fraction of the solar wind ions penetrate into the magnetosphere and, down to the surface and thus provide an indirect source for the exosphere. The airglow experiment on board of Mariner 10 indeed identified atomic hydrogen and helium as constituents of the exosphere with a He surface number density of about 4500 cmp3 estimated by Broadfoot et al. (1976). Curtis and Hartle (1978) estimated the magnetospheric capture fraction of solar wind He*+ required for maintaining the observed He exosphere. Broadfoot et al. (1976), from Mariner 10 data, also discovered oxygen in the Mercury exosphere. Potter and Morgan (1985), from ground-based observations, discovered in the spectrum of Mercury strong emission features at the Fraunhofer sodium D lines attributed to resonant scattering of sunlight from sodium vapour in the exosphere of the planet. They estimated a Na column density of 1011-1012cmp2, which corresponds to a surface number density of 2 x 104-1.5 x 10’ Na atoms cmp3. Potter and Morgan (1986) later also discovered small concentrations of potassium in the Mercury exosphere and suggested that solar wind ion sputtering might account for their existence. Atoms or molecules can be supplied from the surface to the exosphere by thermal evaporation, or from sputtering by charged particles, or by photons. Sputtering of the surface by bombardment of solar wind protons, alpha particles and heavy ions in the aurora1 zones was first
74 discussed by McGrath et al. (1986) by Cheng et al. (1987) and by Killen (1989), but McGrath et al. (1986) suggested that photon stimulated desorption (photon sputtering) be more likely than charged particle sputtering and estimated the photon sputtering fluxes to be of the order of 2 x IO’2 x 10’ cmp2 s-l. Killen and Morgan (1993) however, argued that these fluxes are overly optimistic since they were based on data for alkali halides. The photon sputtering yield depends critically on the low energy cut-off of photons capable of ejecting sodium. Sputtering can be produced only by light with energy larger than the band gap (Townsend and Elliot, 1970) ; Killen et al. (1990) therefore assumed that photons beyond the band gap for NaCl, i.e. IOeV, can sputter Na atoms and obtained an average photon sputtering flux of about 2 x 1O’Na atoms cmp2 s-r, assuming a lunar abundance of sodium. Shemansky and Morgan (1991), on the other hand, estimated an Na photon sputtering flux of only 3.8 x 104cm-2s-1. Thus, there seems to be a discrepancy between the estimated production rates of exospheric Na by photon sputtering of nearly “four orders” of magnitude. Cheng et al. (1987), Morgan et al. (1988, 1989), Sprague (1990) and Cintala (1992) have considered Na and K supply by micrometeorite impact vaporization. Since the rate of volatile production is proportional to the elemental abundance in the regolith, Na and K should be present in the vapour in the same proportion as they exist in the target material. The production of Na by micrometeorite impacts is, however, typically only a few per cent that required for the observed atom content (Cintala, 1992; Killen and Morgan, 1993). Observations of the broad wings in the profile of the sodium D-line, however, suggest that Na extends up to about 700 km above the surface (e.g. Potter and Morgan, 1985; Smyth, 1986; Hunten et al., 1988). This requires the presence of hot Na atoms to match the broad wings of the observed line profile. Neither photon sputtering nor thermal evaporation produce a non-thermal velocity distribution enabling particles to reach such altitudes. Impact vaporization of interplanetary meteoroids or charged particle sputtering could be a source of hot Na atoms with ejection speeds > 2 km s-l (Morgan et al., 1988 ; Ip, 1990,1993). Killen and Morgan (1993) used the estimated solar wind proton fluxes given by Baker et al. (1987) and concluded that ion sputtering is five times smaller than that required to supply the exospheric content, assuming a lunar abundance of sodium in the regolith. If the Na abundance on Mercury were earthlike, which is five times the lunar value, these ion fluxes, however, would be sufficient to supply the exosphere from Mercury’s crust. The Na and 0 column densities were estimated in earlier papers (e.g. Broadfoot et al., 1976; Potter and Morgan, 1985; McGrath et al., 1986; Cheng et al., 1987; Killen et al., 1990 ; Killen and Morgan, 1993) on the basis of scale heights of about 50 and 80 km, respectively, corresponding to a subsolar surface exospheric temperature of 500-575 K.
H. Lammer and S. J. Bauer: Mercury’s exosphere sity profiles of the newborne hot atoms by comparing them with the “cold” density distributions. Ionization processes are reviewed and the equilibrium concentrations of exospheric species are evaluated from considerations about production, chemical loss and plasma diffusion under gravity. Magnetospheric ions, solar wind ions, and locally produced pick up ions can impact the atmospheres or surfaces of bodies in the solar system. These ions transfer their energy by collisions with atmospheric or surface atoms or molecules, they induce an expansion of the atmospheric corona, and are responsible for the loss of a fraction of the hot particles from the atmosphere (e.g. Haff and Watson, 1979 ; Johnson, 1990 ; Luhmann et al., 1992 ; Lammer and Bauer, 1993; Johnson, 1994; Lammer, 1995). Sputtering from a solid planetary surface is calculated in exact analogy with sputtering from an atmosphere ; it produces particles which escape directly or, if their velocity is insufficient, enter ballistic trajectories. We use the energy distribution F(E,) of Sieveka and Johnson (1984) for the determination of the ejection energy E, :
fro=(E, +EeE,J3 {&
yi”“l”‘}
(1)
where E,, is the surface binding energy and Ei the energy of the incident particle. The functions represented by equation (1) have been normalized in such a way that : $ F(EJ dE, = 1.
(2)
Although the surface composition of Mercury is not well known, Na is likely to be bound to 0 with a binding energy of about 2eV, as for example in NaA1,Si30s (feldspar) (McGrath et aE., 1986 ; Cheng et al., 1987). Since 0 atoms themselves are tightly bound to the molecule, the overall binding energy Eb is about 3-4eV (Johnson, 1995). We consider that the incident solar wind protons have an energy E, of 1 keV. Figure 1 shows the shape of the energy distribution F(E,) as a function of the ejected particle energy E,. It is seen that most Na and 0 atoms are ejected with energies of about 1 and 1.7 eV, respectively. These results are in good agreement with the estimations of Ip (1986, 1993). The neutral species injected in Mercury’s exosphere move under the influence of gravity and solar radiation pressure, but collisions between exospheric particles are neglected (Smyth, 1986 ; Ip, 1990). Most atoms return to the surface after a time of flight t,, ; for a particle leaving the surface of a planetary body with a speed v, ejection energy E, and direction 0 with respect to the surface normal, we obtain (Johnson, 1990) : R
to =
max
dr (3)
is (dr/dt)
with Ejection of Na and 0 atoms by charged particle sputtering We now estimate the ejection energies of Na and 0 atoms by charged particle sputtering, as well as the number den-
$=U[l-(~‘)lsin2@-(~)(l-~)~2 where
R, is the surface radius
(4) and
u” the gravitational
15
H. Lammer and S. J. Bauer: Mercury’s exosphere
0.15.....-_...../
0.10
0.05
1 .o
0.1
10.0
E. [evl Fig. 1. Behavior of the normalized energy distribution IF(&) as a function of ejected particle energies E,.Most Na and 0 atoms are ejected into Mercury’s exosphere with about 1 or 1.7 eV. These results are in good agreement with the estimated ejection energies of Ip (1986, 1990, 1993)
(Hunten et al., 1988) and that, the critical level on Mercury is the planetary surface. At great distances above the critical level the barometric law breaks down, since a planetary exosphere populated by escaping particles is not strictly in hydrostatic equilibrium. However, Hodges (1980) has shown that on the Moon and on Mercury, elastic collisions of He atoms with a Maxwellian distribution of vibrating bound atoms produce a “nearly” Maxwellian distribution of velocities, despite the absence of He atoms impinging the surface with velocities in excess of the escape velocity. The exosphere is expanding slightly, and matter is lost, if the gravity field is small, which corresponds to evaporative loss in the kinetic theory. Therefore, the number density n(z) is the product of the barometric density and a partition function c (e.g. Chamberlain, 1963 ; Bauer, 1973; Lammer, 1989; Lammer and Bauer, 1991)
n(z) = n, epcziH)c(x)
(6)
n, is the surface density and x the escape parameter binding energy of the planet. The free flight time to is typically < 2000 s for Na atoms. A fraction of the returning particles penetrate into the surface of Mercury and are recycled into the system ; these implanted ions can cause additional neutral sputtering. Neutrals are lost from the exosphere when they are ionized and swept away under the influence of magnetospheric electromagnetic fields ; due to the eccentricity of Mercury’s orbit, the photoionization time lies in the range 103-104s. One must therefore consider an equilibrium between sources and sinks for each exospheric species. In order to maintain the exosphere in a steady state, the Na supply rate must of course equal the loss rate. If all the ions are escaping, the Na supply rate must be w lo7 atoms cm-* s- ’ (Killen and Morgan, 1993). The maintenance of the present-day exosphere over geologic time requires that a production mechanism can supply Na from great depths and that fresh Na bearing materials be brought to the surface. There is still a large uncertainty about the processes supplying Na or K to the exosphere. However, all processes which bring Na into the exosphere are dependent on the composition of the regolith, megaregolith and upper crust of the planet (Killen and Morgan, 1993).
x(4
GmM = -&-
:
(7)
G is the gravitational constant, M the planetary mass, mass, k the Boltzmann constant, r the planetocentric distance and T the exospheric temperature. The partition function can be written as
m the particle
i: =
(f-9
rbal+La,+ L,.
The first term corresponds jectories :
to particles
with ballistic
tra-
[( 1 s ( )I
i,a~w=$2 Y$X
-7
6s -3
e-Yy 3, X_-y 2
(9)
where
y=xz x+x,
(10)
x, is the escape parameter on Mercury’s y(3/2,x - y) the incomplete I-function :
surface
and
(11)
Mercury’s exosphere The detailed theory of the exosphere made by Chamberlain (1963) is valid for a Maxwellian distribution. The critical level or height, above which collisions are negligible, is called exobase and the number density at this level can be defined as (e.g. Chamberlain, 1963) :
The second orbits :
term
L(x) =
corresponds
$
to particles
in satellite
(12)
“;:“‘12ewy[y(~,x-y)]
1
n, = aH where u is the the scale height. is thus usually Spacecraft and that Mercury’s
gas-kinetic collision cross-section and H The total column content of an exosphere N, d c-’ z 2 x 10’4cm-2 (Bauer, 1973). ground-based observations determined total atmospheric content is < 10’2cm-2
is y
=XTYy3/*-l e-Y 0
dy,
(13)
76
H. Lammer and S. J. Bauer: Mercury’s exosphere
The third term corresponds to the tail of the Maxwellian distribution, i.e. particles with escaping trajectories :
- (xi~~“2e-~[I(~)
-8(:,x-y)]}.
(14)
The Mariner 10 euv airglow spectrometer measurements (Broadfoot et al., 1974) indicated very low values of surface densities y1,, of about 8 cmp3 for H atoms and about 4.5 x 103-6 x lo3 cmW3, for He atoms on the dayside of Mercury (Hunten et al., 1988). A scale height H was determined only for He ; the detection of 0 was only tentative. The He scale height corresponded to a temperature of 575K, roughly consistent with that of the subsolar region. The Na and K observations were made with ground-based optical spectrometers. Morgan et al. (1988) revised the originally reported Na column density of about 8 x 10” cm-’ (Potter and Morgan, 1985) downward to about l-2.5 x lOr* cnl-*, whereas Killen et al. (1990) estimated Na column densities of about 2.83.8 x 10” cmp2. However, these values were derived for a scale height of about 50 km, corresponding to a subsolar surface exospheric temperature of 500-575 K and under the assumption that photon sputtering would be the dominant source of Na atoms. If the population of hot Na atoms originate mostly from charged particle sputtering or meteoroid impact vaporization (Morgan et al., 1988 ; Ip, 1990, 1993) their initial speed must be considerably higher than the thermal velocities of about 0.6 km s-l. The non-thermal Na components have a broader velocity distribution with more particles exceeding the surface escape velocity of 4.3 km s-l. The motion of these hot atoms, can, however, be influenced by solar radiation pressure acceleration (Smyth, 1986; Morgan et al., 1988; Hunten et al., 1988; Ip, 1990), and reach escape trajectories if their velocities exceed 2 km s-l. Atoms with velocities less than 2 km s-l are also influenced by solar radiation and drift from the dayside to the night-side hemisphere (Ip, 1990). If one assumes an ion sputtering yield of Y = 0.1-2, a “lunar” abundance corresponding to a sodium fraction of 4 x low3 in the regolith and incident ion fluxes equal to those given by Baker et al. (1987), one derives a sputtered Na flux in the range 5 x 105-1 x 10scm~*s~’ (Killen and Morgan, 1993). The surface number density of the nonthermal Na component is estimated by dividing the mean velocity of the energy distribution function of the ejected Na atoms, about 2.9 x lO”cms-’ at 1 eV, by Na fluxes given by Killen and Morgan (1993). The surface number density of the hot component ranges between 2 and 345cmW3. If a “terrestrial” abundance of Na (i.e. five times that of the Moon) is considered, then an upper limit for the corresponding Na fluxes is about 2.5 x 1065 x 108cm-2s-’ (Killen and Morgan, 1993) and is sufficient to supply the exospheric content. The surface number densities corresponding to a “terrestrial” Na abundance in Mercury’s regolith lies between 9 and 1730 cme3. Figure 2 shows the number density profiles based on “lunar” abundance for 0, Na and hot Na atoms when
10
1000
100 n
10000
[cm-‘]
Fig. 2. Number density profiles of He, 0, Na and hot Na atoms for “lunar” surface abundance. Both charged particle sputtering and photo-sputtering are considered as possible source mechanisms for Na; particle sputtering is the source of the nonthermal component. Helium dominates over sodium above about 140 km and over oxygen (bound particles) above about 240 km. For helium a full line represents the bound particles with ballistic trajectories and satellite orbits, whereas a dashed line represents the sum of bound and escaping particles. The dot-dashed line (bound particles) shows the number density distribution of hot Na atoms estimated from a sputtering flux of about 1 x 108cm-2s-’ (Killen and Morgan, 1993). The hot Na component extending beyond 700 km could be responsible for the observations of the broad wings in the line profile of the sodium D-line
both charged particle sputtering and photo-sputtering are considered as possible source mechanisms for the Na population, including a non-thermal component. The He atoms dominate over the Na atoms above 140 km and the 0 (dotted line) atoms above about 240 km. The full line represents the density of bound particles with ballistic trajectories and satellite orbits, whereas the dashed line represents the sum of bound and escaping particles (equation (14)). The dot-dashed line gives the number density profile of hot Na atoms, as derived from an upper limit of the sputtered flux of about 1 x lo* cm-2s-i (Killen and Morgan, 1993). The hot Na component extends to altitudes beyond 700 km and could explain the observations of the broad wings in the line profile of the sodium Dline. These hot atoms, however, are influenced by solar radiation pressure acceleration. A small population of high velocity 0 atoms could also be present, but their concentration may lie below the observational limit because losses due to radiation pressure acceleration may become significant. Charged particle sputtering, however, yields a hot Na corona, similar to the N2 corona of Saturn’s large moon Titan (Lammer and Bauer, 1993) and Neptune’s moon Triton (Lammer, 1995).
Mercury’s ionized component The radio occultation experiment on Mariner 10 set an upper limit for the ionospheric density of about lo3 cmW3. As pointed out by Bauer (1975), the absence of an iono-
77
H. Lammer and S. J. Bauer: Mercury’s exosphere spheric layer is consistent with ground and spacecraft optical observations of a total column content N, of ionizable species of d 10” cm-‘. An ionospheric layer cannot form since an optical depth z = 1 is required, corresponding to a column content of absorbing, ionizable constituent NT = ((r,)-‘. Since, typical absorption crosssections B, are of the order z lo-” cm*, z z 10m6and the ionizing radiation penetrates nearly unattenuated down to the planetary surface. The ionization of a species X by the action of a photon with energy hv>IP (ionization potential) is written X = hv(TIP)+X++e.
(15)
Since, the optical depth of Mercury’s exosphere is very small, the ion production function takes the simple form : q = Jn(z)
(16)
where the photoionization rate coefficient J, is about 1.5 x 10-4s-’ for Na and 4 x 10-7s-’ for He (Cheng et al., 1987). The number density distribution of electrons and ions in a planetary “atmosphere” is governed by the equation of continuity (e.g. Bauer, 1973) : (17) where q is the ion pair production, ziDthe plasma diffusion velocity and L(Ni) the chemical loss. If we assume a steady state, two limiting cases can be considered for the continuity equation : (1) “chemical equilibrium”, when chemical processes are predominant and (2) “diffusive equilibrium” when chemical processes can be neglected compared to plasma diffusion. The applicability of either of these to Mercury can be determined by estimating the appropriate time constants. The plasma diffusion time zn is : H2 TD E D,
v-9
where D, = D,/n is the ambipolar diffusion coefficient with Do z 1 x 10’gcm-l s-’ (e.g. Bauer, 1973). Using values of 12and H, representative of Mercury’s exosphere, it is found that the diffusion time for Na ions zn < 1 s. The chemical loss of atomic ions, in the absence of neutral molecular species, is determined by radiative recombination with a coefficient a of the order of lo-‘*cm3 s-’ (Banks and Kockarts, 1973) and the chemical lifetime is rc=p.
1 aNi
(19)
If photochemical equilibrium prevails, Ni can be as high as 4 x 104cm-3. Since chemical lifetime the zc x 2.5 x lo7 s, one can see that recombination is extremely slow compared to plasma diffusion even for high values of Ni. Thus, neglecting the chemical loss term, L(Ni), the continuity equation reduces to : q= Jn=$(Nio)
where J is the photoionization
rate coefficient and n the
number density of the neutral gas. With vD z Do/(d) ion density Ni becomes : Jn2H2 Nip:--.
the
(21)
The resulting ion surface densities for Na, 0 and He ions are < 1 cmw3. Thus, the ionized component of Mercury’s exosphere will not exceed densities of a few electrons/ions cmm3. These results are in a good agreement with Mariner 10 observations of electron densities of about 0.1 cmM3 in the polar cap region and 1 crn3 in Mercury’s magnetotail, although Ogilvie et al. (1977) suggest that the observed electrons within the polar flux tubes at low altitudes are suprathermal electrons and not of atmospheric origin. For understanding convection processes in Mercury’s magnetosphere, one must first know the height-integrated Pedersen conductivity, since it controls the rate of field line merging and hence the ionospheric electric field ; the Pedersen conductivity is written as
cp=e2[& (v;,i$J]
(22)
with ion mass mi, gyro frequency ooY and neutral-ion collision frequency (Banks and Kockarts, 1973), i.e. : v,,~= 2.6 x lo-‘n
0
I/*
2
S-l
(23)
PA
where a, is the polarizability in units of 10-24cm3 of the neutral gas with density n and :
pA =
(
m,mi
(m,+m,)
)
(24)
is the reduced mass in amu. The electron/ion densities encountered on Mercury, lead to an extremely low heightintegrated Pedersen conductivity (up to 1000 km) of about 5 x 10-6mho. Reflectance spectra from Mercury’s surface display the same slope as those from the Moon, a feature which is attributed to the presence of Fe- and Ti-bearing agglutinates in the lunar regolith (Vilas, 1988). The conductivity of the lunar crust is about 10-7mhom-1 (Schwerer et al., 1972 ; Goldstein, 1974) and if one assumes that this figure is also typical of Mercury, then the height integrated Pedersen conductivity is shown to be equivalent to the conductivity of a crustal layer 50 m in thickness. A conducting ionosphere hinders magnetospheric convection because of the drag force exerted by ion-neutral collisions. The ionosphere shorts out the convection electric field induced by solar wind, this effect is known as “ionospheric line tying” and has been discussed in connection with the Earth’s magnetosphere by Cole (1963) and Atkinson (1967). Rassbach et al. (1974) suggested that the maximum possible convection potential is one that drives ionospheric currents sufficiently large to perturb significantly the magnetic field near the dayside magnetopause. However, ionospheric line tying plays an important role only when the height-integrated conductivity is comparable to or greater than a critical conductivity g’cof about 20mho (Hill et al., 1976). The lack
78
H. Lammer and S. J. Bauer: Mercury’s
of ionospheric line tying might mean a very short magnetospheric convection timescale in response to the solar wind condition and hence the particle acceleration process. Since both the derived height-integrated Pedersen conductivity and the estimated surface conductivity are negligibly small compared to the critical conductivity o, of about 20 mho, ionospheric line tying must be virtually absent in Mercury’s magnetosphere. Magnetospheric processes should therefore not be strongly affected by ionospheric conductivity at Mercury as they are at Earth and possibly even at Mars (Rassbach et al., 1974; Hill et al., 1976).
Conclusion Our results are given for an exosphere composed of He, Na and 0 atoms with surface number densities of about 6 x 103, 3.8 x 104, and 4.4 x 104cme3, respectively, which is consistent with observed column contents and assuming lunar abundance for Mercury’s surface constituents. We show that most Na and 0 atoms originating via charged particle sputtering are ejected into Mercury’s exosphere with energies of about 1 and 1.7 eV, respectively, in good agreement with the estimated ejection energies of Ip (1986, 1990, 1993). A hot Na component extending to altitudes beyond 700 km could be responsible for the observations of broad wings in the line profile of the sodium D-line. The density distribution of the hot Na atoms depends on the incident solar wind flux, regolith composition (lunar or terrestrial) and estimated sputtering rates. A small population of high velocity 0
a sub-
exosphere
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The authors wish to thank R. E. Johnson (Department of Nuclear Engineering and Engineering Physics), University of Virginia, U.S.A., for discussions relating to this work and two referees for their helpful comments.
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