Template:Confusing The origin and development of water on terrestrial planets, Venus, Earth, Mars, and the closely related Earth's Moon, varies with each planetary body, with the exact origins remaining unclear.

Water inventory of MarsEdit

A significant amount of surface hydrogen has been observed globally by the Mars Odyssey GRS.[1] Stoichiometrically estimated water mass fractions indicate that - when free of carbon dioxide - the near surface at the poles consists almost entirely of water covered by a thin veneer of fine material.[1] This is reinforced by MARSIS observations, with an estimated 1.6x106 km3 of water at the southern polar region with Water Equivalent to a Global layer (WEG) 11 meters deep.[2] Additional observations at both poles suggest the total WEG to be 30 m, while the Mars Odyssey NS observations places the lower bound at ~14 cm depth.[3] Geomorphic evidence favors significantly larger quantities of surface water over geologic history, with WEG as deep as 500 m.[3] The current atmospheric reservoir of water, though important as a conduit, is insignificant in volume with the WEG no more than 10 µm.[3] Since the typical surface pressure of the current atmosphere (~6 hPa [4]) is less than the triple point of H2O, liquid water is unstable on the surface unless present in sufficiently large volumes. Furthermore, the average global temperature is ~220 K, even below the eutectic freezing point of most brines.[4] For comparison, the highest diurnal surface temperatures at the two MER sites have been ~290 K.[5]

Water inventory of VenusEdit

The current Venusian atmosphere has only ~200 mg/kg H2O(g) in its atmosphere and the pressure and temperature regime makes water unstable on its surface. Nevertheless, assuming that early Venus's H2O had a D/H ratio similar to Earth's Vienna Standard Mean Ocean Water (VSMOW) of 1.6x10-4,[6] the current D/H isotopic ratio in the Venusian atmosphere of 1.9x10-2, at nearly x120 of Earth's, may indicate that Venus had a much larger H2O inventory.[7] While the large disparity between terrestrial and Venusian D/H ratios makes any estimation of Venus's geologically ancient water budget difficult,[8] its mass may have been at least 0.3% of Earth's hydrosphere.[7]

Water inventories of Mercury, Moon, and EarthEdit

Recent observation made by a number of spacecraft confirmed significant amounts of Lunar water. Mercury does not appear to contain observable quantities of H2O, presumably due to loss from giant impacts.[8] In contrast, Earth's hydrosphere contains ~1.46x1021 kg of H2O and sedimentary rocks contain ~0.21x1021 kg, for a total crustal inventory of ~1.67x1021 kg of H2O.[7] The mantle inventory is poorly constrained in the range of (0.5 - 4)x1021 kg.[9] Therefore, the bulk inventory of H2O on Earth can be conservatively estimated as 0.04% of Earth's mass (~6x1024 kg).

Accretion of water by Earth and MarsEdit

The D/H isotopic ratio is a primary constraint on the source of H2O of terrestrial planets. Comparison of the planetary D/H ratios with those of carbonaceous chondrites and comets enables a tentative determination of the source of H2O. The best constraints for accreted H2O are determined from non-atmospheric H2O, as the D/H ratio of the atmospheric component may be subject to rapid alteration by the preferential loss of H [4] unless it is in isotopic equilbrium with surface H2O. Earth's VSMOW D/H ratio of 1.6x10-4[6] and modeling of impacts suggest that the cometary contribution to crustal water was less than 10%. However, much of the water could be derived from Mercury-sized planetary embryos that formed in the asteroid belt beyond 2.5 AU.[10] Mars's original D/H ratio, as estimated by deconvolving the atmospheric and magmatic D/H components in Martian meteorites (e.g., QUE 94201), is x(1.9+/-0.25) the VSMOW value.[10] The higher D/H and impact modeling (significantly different from Earth due to Mars's smaller mass) favor a model where Mars accreted a total of 6% to 27% the mass of the current Earth hydrosphere, corresponding respectively to an original D/H between x1.6 and x1.2 the SMOW value.[10] The former enhancement is consistent with roughly equal asteroidal and cometary contributions, while the latter would indicate mostly asteroidal contributions.[10] The corresponding WEG would be 0.6 - 2.7 km, consistent with a 50% outgassing efficiency to yield ~500 m WEG of surface water.[10] Comparing the current atmospheric D/H ratio of x5.5 SMOW ratio with the primordial x1.6 SMOW ratio suggests that ~50 m of has been lost to space via solar wind stripping.[10]

The cometary and asteroidal delivery of water to accreting Earth and Mars has significant caveats, even though it is favored by D/H isotopic ratios.[8] Key issues include:[8]

  1. The higher D/H ratios in Martian meteorites could be a consequence of biased sampling since Mars may have never had an effective crustal recycling process
  2. Earth's Primitive Upper Mantle estimate of the 187Os/188Os isotopic ratio exceeds 0.129, significantly greater than that of carbonaceous chondrites, but similar to anhydrous ordinary chondrites. This makes it unlikely that planetary embryos compositionally similar to carbonaceous chondrites supplied water to Earth
  3. Earth's atmospheric content of Ne is significantly higher than would be expected had all the rare gases and H2O been accreted from planetary embryos with carbonaceous chondritic compositions.[9]

An alternative to the cometary and asteroidal delivery of H2O would be the accretion via physisorption during the formation of the terrestrial planets in the solar nebula. This would be consistent with the thermodynamic estimate of ~2 earth masses of water vapor within 3AU of the solar accretionary disk, which would exceed by a factor of 40 the mass of water needed to accrete the equivalent of 50 Earth hydrospheres (the most extreme estimate of Earth's bulk H2O content) per terrestrial planet.[8] Even though much of the nebular H2O(g) may be lost due to the high temperature environment of the accretionary disk, it is possible for physisorption of H2O on accreting grains to retain nearly 3 Earth hydrospheres of H2O at 500 K temperatures.[8] This adsorption model would effectively avoid the 187Os/188Os isotopic ratio disparity issue of distally-sourced H2O. However, the current best estimate of the nebular D/H ratio spectroscopically estimated with Jovian and Saturnian atmospheric CH4 is only 2.1x10-5, a factor of 8 lower than Earth's VSMOW ratio.[8] It is unclear how such a difference could exist if physisorption were indeed the dominant form of H2O accretion for Earth in particular and the terrestrial planets in general.

Development of Mars' water inventoryEdit

The variation in Mars's surface water content is strongly coupled to the evolution of its atmosphere and may have been marked by several key stages.

Early Noachian (4.6 to 4.1 Ga)[4] "phyllosian"[11] era
Atmospheric loss to space from heavy meteoritic bombardment and hydrodynamic escape.[4] Ejection by meteorites may have removed ~60% of the early atmosphere.[11][4] Significant quantities of phyllosilicates may have formed during this period requiring a sufficiently dense to sustain surface water, as the spectrally dominant phyllosilicate group, smectite, suggests moderate water: rock ratios.[12] However, the pH-pCO2 equilibria between smectite and carbonate show that the precipitation of smectite would constrain pCO2 to a value not more than 10-2 atm.[12] As a result, the dominant component of a dense atmosphere on early Mars becomes uncertain if the clays formed in contact with the Martian atmosphere,[13] particularly given the lack of evidence for carbonate deposits. An additional complication is that the ~25% lower brightness of the young Sun would have required an ancient atmosphere with a significant greenhouse effect to raise surface temperatures to sustain liquid water.[13] Higher CO2 content alone would have been insufficient, as CO2 precipitates at partial pressures exceeding 1.5 atm, reducing its effectiveness as a greenhouse gas.[13]

Middle to late Noachian (4.1 to 3.8 Ga)[4]
Potential formation of a secondary atmosphere by outgassing dominated by the Tharsis volcanoes, including significant quantities of H2O, CO2, and SO2.[11][4] Martian valley networks date to this period, indicating globally widespread and temporally sustained surface water as opposed to catastrophic floods.[4] The end of this period coincides with the termination of the internal magnetic field and a spike in meteoritic bombardment.[11][4] The cessation of the internal magnetic field and subsequent weakening of any local magnetic fields allowed unimpeded atmospheric stripping by the solar wind. For example, when compared with their terrestrial counterparts, 38Ar/36Ar, 15N/14N, and 13C/12C ratios of the Martian atmosphere are consistent with ~60% loss of Ar, N2, and CO2 by solar wind stripping of an upper atmosphere enriched in the lighter isotopes via Rayleigh fractionation.[4] Supplementing the solar wind activity, impacts would have ejected atmospheric components in bulk without isotopic fractionation. Nevertheless, cometary impacts in particular may have contributed volatiles to the planet.[4]

Hesperian to the present (the "theiikian" era from ~3.8 Ga[4] to ~3.5 Ga and the "siderikian" era postdating ~3.5Ga[14] )
Atmospheric enhancement by sporadic outgassing events were countered by solar wind stripping of the atmosphere, albeit less intensely than by the young Sun.[11] Catastrophic floods date to this period, favoring sudden subterranean release of volatiles, as opposed to sustained surface flows.[4] While the earlier portion of this era may have been marked by aqueous acidic environments and Tharsis-centric groundwater discharge[15] dating to the late Noachian, much of the surface alteration processes during the latter portion is marked by oxidative processes including the formation of Fe3+ oxides that impart a reddish hue to the Martian surface.[11] Such oxidation of primary mineral phases can be achieved by low-pH (and possibly high temperature) processes related to the formation of palagonitic tephra,[16] by the action of H2O2 that forms photochemically in the Martian atmosphere,[17] and by the action of water,[12] none of which require free O2. The action of H2O2 may have dominated temporally given the drastic reduction in aqueous and igneous activity in this recent era, making the observed Fe3+ oxides volumetrically small, though pervasive and spectrally dominant.[14] Nevertheless, aquifers may have driven sustained but highly localized surface water in recent geologic history, as evident in the geomorphology of craters such as Mojave.[18] Furthermore, the Lafayette Martian meteorite shows evidence of aqueous alteration as recently as 650 Ma.[4]

See alsoEdit


  1. 1.0 1.1 Boynton, W. V. et al. (2007), Concentration of H, Si, Cl, K, Fe, and Th in the low and mid latitude regions of Mars, Journal of Geophysical Research Planets, in press doi 10.1029/2007JE002887
  2. Plaut, J. J. et al. (2007), doi 10.1126/science.1139672
  3. 3.0 3.1 3.2 Feldman, W. C. et al. (2004), doi 10.1029/2003JE002160
  4. 4.00 4.01 4.02 4.03 4.04 4.05 4.06 4.07 4.08 4.09 4.10 4.11 4.12 4.13 4.14 Jakosky, B. M. and Phillips, R. J. (2001), doi 10.1038/35084184
  5. Spanovich, N. et al. (2006), doi 10.1016/j.icarus.2005.09.014
  6. 6.0 6.1 National Institute of Standards and Technology (2005), Report of Investigation
  7. 7.0 7.1 7.2 Kulikov, Yu. N. et al. (2006), doi 10.1016/j.pss.2006.04.021
  8. 8.0 8.1 8.2 8.3 8.4 8.5 8.6 Drake, M. J. (2005) Origin of water in the terrestrial planets, Meteoritics and Planetary Science 40 (4), 515-656
  9. 9.0 9.1 Morbidelli, A. et al. (2000), Source regions and timescales for the delivery of water to the Earth, Meteoritics and Planetary Science, 35, 1309-1320
  10. 10.0 10.1 10.2 10.3 10.4 10.5 Lunine, J. I. et al. (2003), doi 10.1016/S0019-1035(03)00172-6
  11. 11.0 11.1 11.2 11.3 11.4 11.5 Chaufray, J. Y. et al. (2007), doi 10.1029/2007JE002915
  12. 12.0 12.1 12.2 Chevrier, V. et al. (2007), doi 10.1038/nature05961
  13. 13.0 13.1 13.2 Catling, D. C. (2007), doi 10.1038/448031a
  14. 14.0 14.1 Bibring, J-P. et al. (2006), doi 10.1126/science.1122659
  15. Andrews-Hanna, J. C. et al. (2007), doi 10.1038/nature05594
  16. Morris, R. V. et al. (2001), doi 10.1029/2000JE001328
  17. Chevrier, V. et al. (2006), doi 10.1016/j.gca.2006.06.1368
  18. McEwen, A. S. et al. (2007), doi 10.1126/science.1143987

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