5169

DISCOVERY OF METEORITIC LAKARGIITE (CaZrO₃), A NEW ULTRAREFRCTORY MINERAL FROM THE ACFER 094 CARBONACEOUS CHONDRITE

Chi Ma. Division of Geological and Planetary Sciences, California Institute of Technology, Pasadena, CA 91125, USA. E-mail: chi@gps.caltech.edu

**Introduction:** During a nano-mineralogy investigation of the Acfer 094 carbonaceous chondrite, lakargiite (CaZrO₃) was identified as submicrometer inclusions in an isolated hibonite grain in section USNM 7233-1. The hibonite grain is described in Simon and Grossman [1]. Lakargiite is a recently found perovskite-group mineral from high-temperature skarns in ignimbrites of the Upper-Chegem volcanic structure, the North Caucasus, Russia [2]. Reported here is the first extraterrestrial occurrence of lakargiite, as a new ultra refractory mineral in a primitive meteorite, among the first solids formed in the solar system. Field-emission SEM with EDS and electron back-scatter diffraction (EBSD) was used to characterize its composition and structure and associated phases.

**Occurrence, Chemistry, and Crystallography:** Lakargiite [Ca<sub>0.95</sub>(Zr<sub>0.87</sub>Ti<sub>0.16</sub>O<sub>3</sub>)] occurs along with tazheranite [(Zr<sub>0.52</sub>Ti<sub>0.18</sub>Ca<sub>0.16</sub>Y<sub>0.06</sub>Fe<sub>0.05</sub>Sc<sub>0.03</sub>)O<sub>1.75</sub>], Zr-bearing perovskite [Ca<sub>0.94</sub>(Ti<sub>0.98</sub>Zr<sub>0.06</sub>)O<sub>3</sub>] and Os-W-alloy [Os<sub>0.81</sub>W<sub>0.13</sub>Fe<sub>0.06</sub>] as fine-inclusions scattered in the central area of the hibonite [Ca (Al<sub>11.70</sub>Ti<sub>0.14</sub>Fe<sub>0.10</sub>Mg<sub>0.06</sub>)O<sub>19</sub>] grain (Fig. 1). Lakargiite appears as irregular or subhedral lath-shaped grains, 300–900 nm in size, sometimes in contact with Os-W or perovskite. Associated tazheranite is a Ca-stabilized cubic zirconia (CSZ), likely being the first reported meteoritic CSZ. EBSD analysis revealed that the lakargiite has a perovskite Phnm structure, identical to that of synthetic CaZrO₃ [3], showing *a* = 5.591 Å, *b* = 5.762 Å, *c* = 8.017 Å, *V* = 258.3 Å³, *Z* = 4.

**Origin and Significance:** Lakargiite is a new Zr-dominant ultrarefractory mineral, joining the Zr-rich refractory minerals allendeite (Sc₄Zr₃O₁₂) [4], tazheranite (Sc- or Y- stabilized cubic zirconia) [4, 5] and panguite [6]. Texturally, lakargiite, tazheranite, Zr-bearing perovskite and Os-W formed before host hibonite. Lakargiite is likely an early condensate. Further work is under way to determine its place in the nebula.

**References:** [1] Simon S. B. and Grossman L. 2011. *Meteorities & Planetary Science*, in press. [2] Galuskin E. V. et al. 2008. *American Mineralogist* 93:1903–1910. [3] Koopmanns H. J. A. et al. 1983. *Acta Crystallographica Section C* 39:1323–1325. [4] Ma C. et al. 2009. Abstract #1402. 40th Lunar and Planetary Science Conference. [5] Ma C. and Rossman G. R. 2008. *Geochimica et Cosmochimica Acta* 72:12S, A577. [6] Ma C. et al. 2011. Abstract #1276. 42nd Lunar and Planetary Science Conference.

<figure><img src="image_20.png" /><figcaption>Fig 1. BSE image showing lakargiite in Acfer 094.</figcaption></figure>

5171

THORTVEITITE (Sc₂Si₂O₇), THE FIRST SOLAR SILICATE?

Chi Ma¹, John R. Beckett¹, Oliver Tschauner¹,² and George R. Rossman¹. ¹Division of Geological and Planetary Sciences, California Institute of Technology, Pasadena, CA 91125, USA. ²High Pressure Science and Engineering Center and Department of Geoscience, University of Nevada, Las Vegas, NV 89154, USA. E-mail: chi@gps.caltech.edu.

**Introduction:** In condensation calculations for a cooling gas of solar composition, or dust enriched variants, melilite is invariably the earliest condensing silicate except at very high dust-to-gas ratios [e.g., 1]. The possibility that silicates of the highly refractory lithophiles, such as Sc and Zr, actually form first is not tested because of a lack of thermodynamic data. Meteoritic occurrences may, therefore, provide the best clues to the earliest evolution of Si in nebular solids. During a nano-mineralogy investigation of the Murchison CM2 carbonaceous chondrite, we identified thortveitite (Sc₂Si₂O₇), along with davisite (CaScAlSiO₆), panguite [(Ti,Sc,Al,Mg)<sub>1.8</sub>O₇], spinel, and Sc-rich diopside, in the Sc-enriched ultrarefractory inclusion MUR1. This is the first meteoritic occurrence of thortveitite and the second natural occurrence of panguite, a Ti-rich oxide recently discovered in Allende [2]. Field-emission SEM with EDS, electron back-scatter diffraction, and electron microprobe were used to characterize the compositions and structures of these two minerals and associated phases.

**Occurrence, Chemistry, and Crystallography:** Thortveitite [(Sc₁.₆₀Zr₀.₁₃Ti₀.₀₈Mg₀.₀₆Y₀.₀₄Fe₀.₀₄Ca₀.₀₂)(Si₁.₉₁Ti₀.₀₈Al₀.₀₁)O₇] occurs with davisite (15 wt% Sc₂O₃), panguite [(Ti₀.₇₁Sc₀.₄₀Al₀.₂₄ Mg₀.₁₄Si₀.₀₈Fe₀.₀₆Zr₀.₀₅Cr₀.₀₅Ca₀.₀₅Y₀.₀₃) Σ₁.₈₂O₃], and MgAl-spinel. These phases form clots that are in 2-D mostly or entirely enclosed in Sc-bearing diopside (3–8 wt% Sc₂O₃). There is a step function in composition between davisites and the diopside where in contact. Panguite appears as irregular grains, 0.5–2 µm in size, along with fine-grained spinel within davisite or serpentinized regions, presumably after davisite. The thortveitite is subhedral, 1.5 × 3.5 µm to 6.5 × 9.0 µm, and is in contact with davisite/ alteration or diopside. Where in contact with davisite, thortveitite often appears embayed, suggestive of a reaction relationship. Thortveitite appears to be inclusion free.

**Origin and Significance:** With the discovery of thortveitite, the Sc-, Zr-rich menagerie in carbonaceous chondrites now includes thortveitite, panguite, davisite, allendeite (Sc₄Zr₃O₁₂) [3], tazheranite (Sc- and/or Y-stabilized cubic zirconia) [3, 4], and lakargiite (CaZrO₃) [5]. MUR1 contains three of these Sc-enriched phases, thortveitite, panguite and davisite. Texturally, thortveitite formed first in MUR1, possibly with later partial reaction with vapor (?) to form davisite + spinel + panguite. Molar Sc/Zr of MUR1 thortveitite (approximately 12) is intermediate between davisite (approximately 15) and panguite (approximately 9), consistent with conservation of Sc/Zr in a reaction between thortveitite and vapor (source of Ti, Ca, Mg, and Al) to produce davisite + panguite + spinel. Since Sc/Zr in all of these Sc-, Zr-bearing phases are much higher than solar (1.2), prior fractionation of a separate Zr-rich phase may be indicated. Further work is required to determine MUR1's place in the nebular framework but, if the thorveitite is an early condensate, it may well be the first silicate and it certainly is more refractory than melilite. Indeed, davisite, which appears to be a reaction product of thortveitite in MUR1, may also predate melilite.

**References:** [1] Ebel D. S. and Grossman L. 2000. *Geochimica et Cosmochimica Acta* 64:339–366. [2] Ma C. et al. 2011. Abstract #1276. 42nd Lunar and Planetary Science Conference. [3] Ma C. et al. 2009. Abstract #1402. 40th Lunar and Planetary Science Conference. [4] Ma C. and Rossman G. R. 2008. *Geochimica et Cosmochimica Acta* 72:12S, A577. [5] Ma C. 2011. 74th Annual Meeting of the Meteoritical Society. Abstract #5169.
## POROSITY, MAGNETIC SUSCEPTIBILITY, AND DENSITY OF LUNAR METEORITES

R. J. Macke¹, D. T. Britt¹, W. S. Kiefer², A. J. Irving³ and G. J. Consolmagno⁴. ¹Department of Physics, University of Central Florida. E-mail: macke@jesuits.net. ²Lunar and Planetary Institute. ³Department of Earth and Space Sciences, University of Washington Seattle. ⁴Specola Vaticana.

**Introduction:** The interpretation of gravimetric and topographical data from lunar missions to model the Moon's interior requires a comprehensive database of lunar rock densities and porosities. To date we have measured porosity, density and magnetic susceptibility for five Apollo lunar samples and 50 stones from 44 of the 74 known lunar meteorites, representing a wide range of lithologies and lunar regions. Our measurement methods as outlined in Consolmagno [1] are fast, non-destructive and non-contaminating. Grain density is measured by helium ideal-gas pycnometry, bulk density by the glass bead method developed by Consolmagno and Britt [2]. Porosity is calculated directly from bulk and grain densities. We also measure magnetic susceptibility, using an SM-30 magnetic susceptibility meter [3].

**Basalts:** The low-Ti (1–5 wt% TiO₂) basalts measured here group tightly in grain density (avg. 3.30 ± 0.06 g cm⁻³) and magnetic susceptibility (log χ = 2.86 ± 0.11). Porosities range from zero to 10%, averaging 5.6%. There is no discernible difference between the two Apollo basalts (12051 and 15555) and the six basaltic meteorites taken as a whole.

**Feldspathic Breccias:** Grain densities averaged 2.85 ± 0.11 g cm⁻³. Magnetic susceptibilities varied over a wide range, from log χ = 2.1 to 4.5. Two distinct populations are visible in porosity, with most below 8% porous, but a few (most notably all stones of the pairing NWA 2995/2996/4503/5151/6252) have a porosity in excess of 10%. These samples represent a mix of regolith breccias, fragmental breccias, and impact-melt breccias.

**Other Breccias:** Diverse lithologies (including KREEP-rich breccias) occupy a different range of density and magnetic susceptibility than the other two groups. The meteorites in this group have low porosities (like the feldspathic breccias and basalts), but in contrast Apollo breccias have porosities above ~15%.

**Discussion:** The data presented here are early results of an ongoing study. Much work still remains to be done, including further measurements already in progress of both meteorites and Apollo samples. Expanding the database of Apollo samples will help address issues about the porosity of in-situ breccias as opposed to meteorite breccias. It will also fill in missing data with the inclusion of thus-far unrepresented high-Ti (9–14 wt% TiO₂) basalts as well as expanding the representation of low-Ti basalts. In addition, the Apollo samples have known geologic contexts and include lithologies that have not been recognized in the lunar meteorite collection.

**Acknowledgments:** CAPTEM, Caroline Smith (NHM), Carl Agee (UNM), Kevin Righter (NASA-JSC), Ted Bunch (NAU), Randy Korotev (WUSTL), Ian Nicklin (ROM), Philip Mani, Gregory Hupé, Don Edwards, and Andrew Abraham. NASA Planetary Geology and Geophysics Program grant NNX09AD91G.

**References:** [1] Consolmagno G. J. et al. 2008. *Chemie der Erde*—*Geochemistry* 68:1–29. [2] Consolmagno G. J. and Britt D. T. 1998. *Meteoritics & Planetary Science* 33:1231–1241. [3] Gattacceca J. et al. 2004. *Geophysical Journal International* 158:42–49.

## SCHREIBERSITES IN MINCY AND VACA MUERTA MESOSIDERITES

D. E. Magnelli¹, M. E. Varela¹ and F. Brandstätter². ¹ICATE-CONICET, San Juan, Argentina. ²Naturhistorisches Museum, Vienna, Austria.

**Introduction:** The iron-nickel phosphide schreibersite is a common accessory phase in many iron and iron-rich meteorites [1]. In mesosiderites, it is found as disseminated crystals characterized by a high Ni content (approximately 50 wt%) [2]. This contribution is a report on petrographic features and chemical compositions of schreibersites from Mincy (nine objects) and Vaca Muerta (four objects) mesosiderites.

**Samples and Methods:** The studied PTS are Vaca Muerta K and Mincy L from NHMV. Electron microprobe measurements were performed with an ARL-SEMQ at NHMV.

**Results:** Phosphide (Ph) crystals are totally or partially (e.g., metal + silicate) enveloped by metal. The schreibersite/silicate interfaces are irregular with schreibersite copying silicates edges. Conversely, schreibersite shows regular – occasionally euhedrally outlined-faces when it is enveloped by metal (kamacite [αFe] and/or taenite [γFe]). The grain size is quite variable and ranges from 20 µm up to 300 µm. Nickel profiles performed across Ph + αFe on Mincy (Fig.) are in agreement with those previously measured on Emery [3]. We also show the distribution of Ni when schreibersite is in contact with taenite (γFe).

**Acknowledgment:** The present research was supported by Agencia PICT 212.

**References:** [1] Skála R. and Cisarová I. 2005. *Physics and Chemistry of minerals* 31:721–732. [2] Powell B. N. 1971. *Geochimica et Cosmochimica Acta* 35:5–34. [3] Kulpecz A. A. and Hewins R. H. 1978. *Geochimica et Cosmochimica Acta* 42:1495–1500.

<figure><img src="image_21.png" /></figure>
5157

EQUILIBRIUM TEMPERATURE ESTIMATIONS ON COEXISTING
PHASES IN MESOSIDERITES

D. E. Magnelli¹, M. E. Varela¹, F. Brandstätter² and R. B. Scorzelli³. ¹ICATE-CONICET, San Juan, Argentina. E-mail: dmagnelli@icate-conicet.gob.ar. ²Naturhistorisches Museum, Vienna, Austria. ³Centro Brasileiro de Pesquisas Físicas, RJ, Brazil.

**Introduction:** The thermal history of mesosiderites is not totally understood yet. Materials with different physical and chemical properties (e.g., silicates, oxides, sulfides, and metal) make up the whole rock, showing a heterogeneous mineral distribution with variable degrees of recrystallization. Here we report on new equilibrium temperature calculations made on the silicate pairs high-Ca/low-Ca pyroxene and low-Ca pyroxene/olivine. We took into account only clinopyroxene and orthopyroxene exsolutions present as lamellae and/or blebs.

**Methods:** To assess equilibria among pyroxene–pyroxene and low-Ca pyroxene–olivine pairs, we have used the oxythermobarometry program QUILF [1]. The program was run using 1 bar as pressure condition since mesosiderites silicates have been considered to be part of shallow crustal levels [e.g., 2, 3].

**Samples:** The studied PTS are Vaca Muerta K, Vaca Muerta O, Crab Orchard L3771, Mincy L, Estherville M2072 (from NHMV) and ALHA77219,74-10, Asuka-87106.51.3 and Asuka-882023.51-3 (from NIPR).

**Results and Discussion:** For thermometric purposes, the molar compositions of all involved phases were recalculated into the scheme Mg-Fe-Ca according to QUILF projection. Almost all compositions corresponding to both, high- and low-Ca pyroxenes nearly fit the equilibrium model in Fe-Mg exchange. The equilibrium temperatures fall in the range from 872 °C ± 49 °C to 1057 °C ± 28 °C. This temperature range is in agreement with that previously estimated for Budulan and Mincy [4]. For Crab Orchard, QUILF gives temperatures with relatively large uncertainty values (1018 °C ± 75 °C to 1025 °C ± 121 °C) suggesting that equilibrium conditions were not established.

If olivines and orthopyroxene rims are taking into account the estimated temperatures give similar results. Concerning Fe-Mg exchange, the highest temperature corresponds to Vaca Muerta K 1038 °C ± 0 °C and the lowest to Mincy L (848 °C ± 0 °C) with a mean value placed at 950.5 °C. In all olivine-orthopyroxene pairs an exact solution (±0 °C uncertainty) was verified indicating that Fe-Mg partitioning has achieved a near equilibrium state. However, if Fe-Mg-Ca is taken into account, the measured amount of calcio-olivine component is too low to reach equilibrium conditions. Did olivine lose Ca during cooling?

**Acknowledgment:** The authors are greatly indebted to Dr. Donald Lindsley for the helpful discussions. This work was supported by Agencia PICT 212.

**References:** [1] Andersen D. J. et al. 1993. *Computer & Geosciences* 19:1333–1390. [2] Powell B. N. 1971. *Geochimica et Cosmochimica Acta* 35:5–34. [3] Ganguly J. et al. 1994. *Geochimica et Cosmochimica Acta* 58:2711–2723. [4] Hewins R. H. and Harriot T. A. 1986. *Proceedings, 16th Lunar and Planetary Science Conference*. pp. D365–D372.

5303

Li Be B IN REFRACTORY HIBONITE OBJECT

K. K. Marhas and Y. Kadlag. Physical Research Laboratory, Navrangpura, Ahmedabad, Gujarat, India 380009. E-mail: kkmarhas@prl.res.in.

Presence of ¹⁰Be (t₁/₂~1.5 Ma) in the early solar system has been confirmed with initial ¹⁰Be/⁹Be ranging from 3 × 10⁻⁴–18 × 10⁻⁴ [1–5] in various early solar system objects. Potential source to the origin of this short-lived nuclide in the early solar system has been argued between irradiation of gas and dust in the solar nebula by solar energetic particles [6, 7] and ¹⁰Be-GCR trapped within the proto-solar molecular cloud [8]. Possible presence of ⁷Be (t₁/₂~53 days) in the Allende CAI [9] favors the energetic particle interaction as the source of these (⁷Be, ¹⁰Be) short-lived nuclides whereas the absence of ⁷Be in platelet hibonite gives an additional outlook on mixing of ⁷Li produced from spallation and the chondritic material [5]. Platelet hibonites with the absence of ²⁶Al and ⁴¹Ca are considered as the first forming solids in the early solar system [10] and hence should have seen the early irradiation in turn producing ⁷Be along with ¹⁰Be within the grain. We have initiated Li Be B study of platelet type hibonites with absence of ²⁶Al and ⁴¹Ca in order to check on the possible presence of ⁷Be and to understand the irradiation scenario in the early solar system.

SH-7 is a 2 mm-sized hibonite rich refractory object found on a broken surface of the Murchison meteorite [11]. It is devoid of short-lived nuclide ²⁶Al [(²⁶Al /²⁷Al)ᵢ < 1.5 × 10⁻⁶] and ⁴¹Ca [(⁴¹Ca /⁴⁰Ca)ᵢ < 3.3 × 10⁻⁹] and indicates presence of ¹⁰Be with (¹⁰Be/⁹Be)ᵢ ~4 × 10⁻⁴ [10, 3]. ¹⁶O and group III REE pattern [1] imply object to be refractory with negligible effect of secondary processes after its formation.

Utilizing the high spatial resolution capability of NanoSIMS, PRL nanoSIMS was used in multicollection mode (⁶Li⁺, ⁷Li⁺, ⁹Be⁺, ¹⁰B⁺, ¹¹B⁻) with an O⁻ primary beam of approximately 2 nA rastering the sample for more than 1.5 h in order to achieve better counting statistics. Eight different areas of the refractory object (SH7) were analyzed with varying ⁹Be/⁶Li values (ranging from 90 to 280). Concentration of Be for these analyzed areas varied from 140 to 500 ppb. (¹⁰Be/⁹Be)ᵢ values obtained by PRL nanoSIMS are reproducible (within errors) compared to the earlier SIMS measurements on the same sample [3].

No definite excess in ⁷Li has been observed for SH7 hibonite. Lower value of the ⁷Li/⁶Li (minimum approximately -100‰) as compared to chondritic value (⁷Li/⁶Li = 12.02) in most of the analyses probably indicates the mixing of the spallation product (⁷Li/⁶Li = 0.8) with the chondritic material (⁷Li/⁶Li = 12.02). Further understanding definitely requires a detailed study of diffusion processes of Li in hibonite grains with different conditions.

**References:** [1] McKeegan K. D. et al. 2000. *Science* 289: 1334–1337. [2] Sugiura, N. et al. 2001. *Meteoritics & Planetary Science* 36: 1397–1408. [3] Marhas K. K. et al. 2002. *Science* 298:2182–2185. [4] MacPherson, G. J. et al. 2003. *Geochimica et Cosmochimica Acta* 67: 3165–3179. [5] Liu M. -C. et al. 2010. *The Astrophysical Journal Letters* 719: L99–L103. [6] Goswami J. N. et al. 2001. *The Astrophysical Journal* 549:1151–1159. [7] Gounelle M. et al. 2001. *The Astrophysical Journal* 548:1051–1070. [8] Desch S. J. et al. 2003. *The Astrophysical Journal* 602: 528–542. [9] Chaussidon M. et al. 2006. *Geochimica et Cosmochimica Acta* 70:224–245. [10] Sahijpal S. et al. 1998. *Nature* 391:559–561. [11] Hashimoto A. et al. 1986. 17th Lunar and Planetary Science Conference. pp. 317–318.
5183

Si ISOTOPIC COMPOSITION OF EFREMOVKA E101.1 CAI

J. Marin-Carbonne¹, J. Aléon², K. D. McKeegan¹ and A. El Goresy³.
¹Department of Earth and Space Sciences, University of California, Los Angeles, CA 90095-1567, USA. ²CSNSM, CNRS/IN2P3/ Univ. Paris Sud, Bat 104, 91405 Orsay Campus, France. ³Bayerisches Geoinstitut, Bayreuth, Germany. E-mail: jmarin@ess.ucla.edu, mkeegan@ess.ucla.edu

Inclusion E101.1 from the CV3 carbonaceous chondrite Efremovka is a compact type A CAI, highly enriched in ultrarefractory oxides [1]. This inclusion consists of several different lithological subunits with a complex petrography. The core contains ultrarefractory fassaitic clinopyroxene, enriched in Sc and Zr, enclosed in large melilite crystals [1]. This core also contains some impact glass [1, 2]. Surrounding the inclusion is a well-defined Wark-Lowering rim (WLR). This inclusion has been extensively studied for O [2] and Mg isotopes [3].

Here, we investigate the Si isotopic composition of the different sub-units of this inclusion. Our goals are to better understand the conditions of formation of the CAI and the relationship between the different lithological units. We performed Si isotope analyses using the Cameca ims 1270 at UCLA in multicollection mode. San Carlos olivine, pyroxene and melilite reference materials were used to correct for instrumental mass fractionation. The reproducibility obtained on these reference materials is better than 0.4‰ at one sigma.

One profile of 200 µm from the rim to the core of the inclusion was performed. Melilite close to the WLR is fairly homogeneous, around—4.8 ± 1.6‰, in contrast to the melilite from the center of the core. Only two analyses were performed on the center of the core and we will perform more analyses to investigate this discrepancy. The impact glass is isotopically lighter than the melilite with a mean value of -7.7‰. The WR is more fractionated than the core of the inclusion with values between +7.3‰ and +14.7‰.

No correlation is apparent between Mg and Si isotopic compositions in the few data that we have. As shown by the O [2] and the Mg [3] studies, this inclusion could have experienced some evaporation event. Further Si isotope analyses are currently ongoing to investigate in detail the relationship between the sub-units and the nature of the evaporation event(s) forming the inclusion rim.

**References:** [1] El Goresy et al. 2002. *Geochimica et Cosmochimica Acta* 66:1459-1491. [2] Aléon et al. 2010. *Meteoritics & Planetary Science* 45:A7. [3] Aléon et al. 2011.

5383

Mg AND Si ISOTOPIC COMPOSITIONS OF VIGARANO FUN INCLUSION 1623-5

J. Marin-Carbonne¹, K. D. McKeegan¹, A. M Davis², G. J. MacPherson³, R. A. Mendybaev², and F. M. Richter².
¹Department of Earth and Space Sciences, University of California, Los Angeles, CA 90095-1567, USA. (jmarin@ess.ucla.edu).
²Department of the Geophysical Sciences and Enrico Fermi Institute, University of Chicago, IL 60637-1433, USA.
³Department of Mineral Sciences, National Museum of Natural History, Smithsonian Institution, Washington DC 20560-0119, USA.

Among the CAIs, some inclusions have extreme isotopic properties, with large mass-dependant isotopic fractionation of major elements (Mg, Si, O) accompanied by large nonradiogenic nuclear anomalies in various minor and trace elements [1]. The origin of these fractionations and the composition of the precursors of these so-called FUN inclusions are still unknown. Vigarano 1623-5 has strong isotopic affinities to the Allende FUN inclusion C1 and is the only “classical” FUN inclusion that still exists in thin section and so can be studied by in situ methods [2]. This object, which shows the first unequivocal correlated petrologic and isotopic evidence for volatilization [2, 3], is composed of two chemically distinct parts, separated by a melilite-rich mantle and a hibonite-bearing vesicle. One part is pyroxene-rich and the other is a melilite-rich. A discontinuous 200 µm thick melilite mantle surrounds these cores [2]. Here, we investigate in situ isotopic compositions of Mg and Si of the different parts of this unique inclusion. Our goal is to better understand the relationship between the mantle and the core of this inclusion to constrain its origin and environment of formation.

The analyses were done on a Cameca ims-1270 at UCLA at high mass resolution in monocollection and multicollection modes. Olivine, pyroxene, and melilite standards were used to correct for instrumental mass fractionation. The matrix effect due to variation of Åk content [4] was also investigated by analyzing synthetic zoned melilite.

Individual minerals show a large range of mass fractionation in δ³⁰Si and δ²⁵Mg, forsterite from +6.9 to +23.2‰ and +23.1 to +29.0‰, respectively, and mantle melilite from +6.8 to +26.2‰ and from +30.6 to +32.9‰, respectively. Olivines show less mass-dependent fractionation for Mg than do pyroxene and melilite, but the same range of fractionation for Si is present in all of these minerals. There is no apparent relationship between the magnitude of fractionation and spatial distribution in either core of the inclusion. Olivines show a correlation between Mg and Si isotopes, which is slightly different from the relationship defined by the free evaporation of Mg₂SiO₄ residues [5]. The relationship between the degree of evaporation and the isotopic mass fractionation can be modeled using the Rayleigh equation. Our data imply an evaporation fraction (taking α = 0.99009, [6]) of approximately 72% for the pyroxene-rich core and approximately 84% for the mantle. The precursor of the core of 1623-5 was at least partially evaporated as it was crystallizing.

**References:** [1] Loss, R. D. et al. 1994. *The Astrophysical Journal*. 436:193-196. [2] Davis A. M. et al. 1991. *Geochimica et Cosmochimica Acta*, **55**, 621-637. [3] McKeegan K. D. et al. 2005. Abstract #2077. 36th Lunar and Planetary Science Conference. [4] Knight K. B. et al. 2009. *Geochimica et Cosmochimica Acta*, **73**:6390-6401. [5] Davis A. M. et al. 1990. *Nature*, **347**:655-657. [6] Mendybaev et al. 2009. Abstract #2461. 40th Lunar and Planetary Science Conference.
5148

ORIGIN OF XENON ISOTOPIC COMPOSITION OF PHASE Q

Y. Marrocchi¹, M. Kuga¹, B. Marty¹ and P. Reinhardt². ¹CRPG-CNRS, BP20, 54501 Vandoeuvre les Nancy, France. ²Laboratoire de Chimie Théorique, UPMC-Paris VI, 75252 Paris, E-mail: yvesm@crpg.cnrs-nancy.fr.

**Introduction:** Noble gases trapped in primitive meteorites are isotopically mass-dependant fractionated with respect to solar noble gases. They are quantitatively hosted by a poorly defined, organic phase [1], labeled phase Q. Xenon, which is the least prone to isotopic fractionation due to its lowest relative mass difference between isotopes, is enriched in heavy isotopes by +1.26 ± 0.05% per atomic mass unit in the carbonaceous chondrite Orgueil relative to Solar. In order to understand the origin of phase Q and its xenon isotopic composition, we have performed an experimental study coupled with theoretical calculation.

**Experimental Method and Theoretical Calculation:** We performed adsorption experiments (< 1 eV) using Xe atoms and ions, ionized in a radiofrequency plasma. Within the reaction vessel, anthracite was heated in a furnace and the resulting organic particles were deposited onto the walls of the vessel, resulting in carbon-rich films. Xe was trapped either as ions in the ionization zone of the vessel, or as neutral atoms outside this zone.

Using the finite nucleus model [2], we determined the first ionization potential (FIP) by solving the relativistic Schrödinger–Dirac equation for each Xe isotopes for neutral atoms and for ions.

**Results:** Xe trapped from ionized Xe is tightly bound and is enriched by +1.36 ± 0.06% amu⁻¹ (2σ) in heavy isotopes, reproducing the isotopic fractionation of Xe trapped in chondrites relative to Solar. Neutral Xe is more loosely trapped, is much lower in abundance, and is not isotopically fractionated. Ionized conditions allow the monotonic Xe isotopic composition observed in meteorites during stepwise heating release to be reproduced.

**Discussion:** Xe was not implanted into the film, because the energies of the incident Xe atoms and ions are far too low (< 1 eV). From the difference of behavior between ionic and neutral forms, we propose that Xe ions were trapped *via* chemical bonding at the surface of the newly created C-rich film. The observed mass-dependant fractionation of Xe is unlikely to have occurred in the gas phase as the difference in FIPs correspond to ≈ 1 × 10⁻⁴ eV between ¹²⁴Xe and ¹³⁶Xe. As the mean experimental first ionization energy of xenon is 12.13 eV, the isotopic effect makes only ≈ 0.00008% difference in ionization energy between two Xe isotopes. Xe isotopic fractionation is more probably related to variations in chemical bonding strengths of Xe isotopes as chemical bonds involving heavy Xe isotopes are more stable than those involving light ones. The extent of fractionation could have been kinetically enhanced as a result of competition between film growth and Xe dissociation rate. The Q-noble gas fractionation and trapping characteristics are likely related to UV irradiation (from nearby young stars) of the surface of growing organic grains in the outer part of the solar system. For young stars, photons emitted in the far UV energy range (< 100 nm) that are capable of ionizing noble gases were orders of magnitude more abundant than for the Present-day Sun, allowing efficient ionization of gaseous species

**References:** [1] Marrocchi et al. 2005. *Earth and Planetary Science Letters* 236:569–578. [2] Dyall K.G. et al. 1989 *Comp. Phys. Comm.* 55:425–446.

5335

A Fe-RICH PHASE ENRICHED IN REE AND HFSE: NEW CLUES ON CHONDRITE HISTORY?

Céline Martin¹, Vinciane Debaille², Philippe Claeys¹, ¹Earth System Sciences, Vrije Universiteit Brussel, Belgium. ²Département des Sciences de la Terre et de l'Environnement, Université Libre de Bruxelles, Belgium.

Chondrites are among the most pristine objects of the solar system and their average chemical composition estimates the initial composition of the Earth. Particularly, the Sm-Nd and Lu-Hf isotopic systems are used together as tracers of early planetary processes because they are not fractionated either by condensation/volatilization processes in the early solar system or by metal/silicate separation during core formation in terrestrial planets. The CHUR (CHondritic Uniform Reservoir) serves as the established reference value for these two systems. It is clearly demonstrated that Sm-Nd values are well constrained and homogeneous in chondrites. However, recent data shows that this is not the case for Lu-Hf system that displays large Lu/Hf ratio heterogeneities (up to 28%), hence restraining its use in cosmo/geochemistry. A better knowledge of these meteorites and the geological processes they underwent will ensure a better understanding of the first million years of the solar system, including the early evolution of terrestrial planets and more importantly the Earth.

The study of accessory phases provides significant information on the PT conditions of their formation as well as on their alteration state. Moreover, such minerals are assumed to contain a large part of rare earth elements (REE) and high field strength elements (HFSE). Among those accessory phases, an undescribed Fe(±Ni)-rich phase, founded as single grain as well as rim of sulfide or as veins has been found almost ubiquitously in chondrites (H3, H4, L6). The interest of this kind of interstitial phase will be discussed in order to better constrain Sm-Nd and Lu-Hf systematics, as it contains significant amount of LREE (up to 75 ppm of Nd) and Hf (up to 40 ppm).
THREE-DIMENSIONAL OBSERVATION AND IMAGE ANALYSIS
OF ORGANIC NANOGLOBULES IN CARBONACEOUS
CHONDRITES USING X-RAY MICROTOMOGRAPHY

T. Matsumoto¹, A. Tsuchiyama¹, K. Nakamura-Messenger², M. E. Zolensky², T. Nakano³, and K. Uesugi⁴. ¹Department of Earth and Space Science, Osaka University, Osaka, Japan. E-mail: matsumoto@as-troboy.ess.sci.osaka-u.ac.jp. ²NASA JSC, Houston, TX 77058, USA. ³GSJ/AIST. ⁴JASRI/SSpring-8.

**Introduction:** In carbonaceous chondrites, IDPs, and dust from comet 81P/Wild 2, spherical organic objects of a few hundred nanometers, called organic nanoglobules were found [1–3]. They have usually hollow structures, and it has been proposed that they were formed from composite particles of organic materials and ice in molecular clouds or aqueous alteration of organic materials in asteroids [4]. In the both origins, hollow regions of the globules might be filled with H₂O-rich fluids. However, any fluids have not been detected due to destructive analysis, such as by TEM, in the previous studies. In order to determine the existence of such fluids, we tried to observe organic nanoglobules nondestructively using synchrotron radiation-based X-ray microtomography.

**Experiments:** The imaging experiments were done at BL47XU of SSpring-8, JAPAN. The samples were pieces of Tagish Lake meteorite of 60 µm in size. Their 3-D structures were obtained by the tomography system with the spatial resolution of about 250 nm. Then, the samples were microtomed and observed under a transmission electron microscope (TEM) to compare with the CT images.

**Results and Discussions:** Comparison between the CT and TEM images showed that nanoglobules were identified by the CT images. However, CT images of nanoglobules are not exactly the same as TEM images probably due to X-ray refraction in the sample. In order to evaluate CT images of nanoglobules and to determine the presence of fluids in the hollows, simulation of CT images was made by considering X-ray absorption and refraction for nanoglobules with or without water in the hollows. The simulation suggests that at least the observed nanoglobule do not seem to contain any water in the hollow. The simulation also revealed that the external shapes of nanoglobules can be identified although internal structures are hardly recognized by CT images. Based on the simulation, three-dimensional shapes of nanoglobules in the Tagish Lake samples were examined. Three-axial ellipsoids approximation of the nanoglobule shapes revealed that they are almost spherical (the ratio of the longest and shortest axis is no less than approximately 0.7). If nanoglobules have molecular clouds origin based on N and H isotopic features [4], icy dust in the molecular clouds should be spherical.

**References:** [1] Nakamura K. et al. 2002. *International Journal of Astrobiology* 1:179. [2] Messenger S. et al. 2008. Abstract #2391. 34th Lunar and Planetary Science Conference. [3] De Gregorio B. T. 2009. Abstract #1130. 40th Lunar and Planetary Science Conference. [4] Nakamura-Messenger K. et al. 2006. *Science* 314:1439–1442.

SILICATE-BEARING SULFIDE-METAL ASSEMBLAGES IN
PRIMITIVE ORDINARY CHONDRITES

M. Matthes¹ and A. Bischoff¹. ¹Institut für Planetologie, Wilhelm-Klemm-Str. 10, 48149 Münster, Germany. E-mail: max.matthes@uni-muenster.de.

**Introduction:** The ordinary chondrites Adrar 003 (L/LL3.10), Krymka (LL3.2), and Semarkona (LL3.00) belong to the most primitive chondrites known [1]. These rocks were found to contain different types of silicate-bearing sulfide-metal assemblages in terms of texture, size, and relative abundance of silicates and opaques (sulfides, metal). We have studied 53 aggregates within the three chondrites and found that troilite is the most abundant phase within the assemblages. As major silicates olivine and pyroxene occur.

**Mineralogy:** The assemblages were identified and documented in thin section by optical microscopy in reflected and transmitted light. Most of the assemblages occur as individual objects between chondrules. In several cases these aggregates form broad chondrule rims presumably indicating a close genetic relationship. The mineral sizes within the aggregates vary extremely as well as the relative abundances of silicates and opaques. Besides troilite, pentlandite is of minor abundance and the metal grains analyzed so far are primarily kamacite; taenite only occurs sporadically.

341 olivines and 157 low-Ca pyroxenes were analyzed within these sulfide-metal assemblages by SEM-EDX. Although the range and average composition of pyroxene within the aggregates are similar to those of the bulk chondrites, significant differences were obtained for olivine. Considering all three samples, olivine provides a much broader range in Fa-content in the assemblages than in the bulk rocks. In addition, the average olivine composition of the grains within the aggregates of all three samples (~Fa₃₀) is significantly different from the average composition of olivine in the three bulk rocks (~Fa₁₂; Table 1).

**Discussion:** The different textural and mineralogical characteristics of the silicate-bearing sulfide-metal aggregates suggest different formation conditions. The differences in olivine compositions (bulk chondrites vs. aggregates) will be considered in detail in ongoing studies.

**Reference:** [1] Grossman J. N. and Brearley A. J. 2005. *Meteoritics & Planetary Science* 40:87–122.

<table><caption>Table 1. Compositions of olivine (Ol) and low-Ca pyroxene (Px) in the studied bulk chondrites and within sulfide-metal aggregates (agg.); n = number of analyzed grains; data in mol%.</caption><thead><tr><th></th><th>Ol (Fa)</th><th>n</th><th>Px (Fs)</th><th>n</th></tr></thead><tbody><tr><td><strong>Adrar 003</strong></td><td></td><td></td><td></td><td></td></tr><tr><td>Bulk chondrite (range)</td><td>0.4–35.8</td><td>98</td><td>0.6–40.3</td><td>74</td></tr><tr><td>Average</td><td>13.2 ± 8.1</td><td></td><td>10.4 ± 8.0</td><td></td></tr><tr><td>Sulfide-metal agg. (range)</td><td>0.3–60.0</td><td>155</td><td>1.4–27.4</td><td>69</td></tr><tr><td>Average</td><td>32.4 ± 15.1</td><td></td><td>10.1 ± 7.0</td><td></td></tr><tr><td><strong>Krymka</strong></td><td></td><td></td><td></td><td></td></tr><tr><td>Bulk chondrite (range)</td><td>0.4–47.1</td><td>92</td><td>1.0–32.5</td><td>54</td></tr><tr><td>Average</td><td>12.0 ± 8.5</td><td></td><td>8.4 ± 7.2</td><td></td></tr><tr><td>Sulfide-metal agg. (range)</td><td>0.6–60.5</td><td>128</td><td>0.6–63.2</td><td>65</td></tr><tr><td>Average</td><td>29.8 ± 16.0</td><td></td><td>10.0 ± 10.9</td><td></td></tr><tr><td><strong>Semarkona</strong></td><td></td><td></td><td></td><td></td></tr><tr><td>Bulk chondrite (range)</td><td>0.3–26.3</td><td>65</td><td>0.6–32.3</td><td>46</td></tr><tr><td>Average</td><td>10.8 ± 6.9</td><td></td><td>8.9 ± 7.4</td><td></td></tr><tr><td>Sulfide-metal agg. (range)</td><td>1.1–37.3</td><td>58</td><td>2.6–23.4</td><td>23</td></tr><tr><td>Average</td><td>26.3 ± 8.1</td><td></td><td>9.9 ± 7.2</td><td></td></tr></tbody></table>
Fe²⁺-Mg PARTITIONING BETWEEN OLIVINE AND MARTIAN BASALTIC MELTS

A. K. Matzen, J. R. Beckett, M. B. Baker and E. M. Stolper. California Institute of Technology, MC 170-25, Pasadena, CA 91125, USA. E-mail:amatzen@caltech.edu.

**Introduction:** The partitioning of elements between olivine and liquid is key to understanding igneous processes in mafic and ultramafic systems [1]. The landmark study of Roeder and Emslie [2] concluded that the olivine (ol)-liquid (liq) exchange coefficient, $K_{D,Fe_{2+}-Mg} = (FeO/MgO)^{ol} / (FeO/MgO)^{liq}$ (by weight), is 0.30 ± 0.03, independent of temperature (T) and liquid composition. Subsequent work [3] has shown a dependence on liquid composition, but the canonical value of 0.30 is still widely used in both terrestrial and Martian applications [e.g., 4, 5]. Here, we use previously published experiments on bulk compositions relevant to Martian magmas to determine the appropriate $K_{D,Fe_{2+}-Mg}$ (s) for Martian systems; we then use this value to test whether any of the ol-phyric shergottites represent liquids.

**Experiments on Martian Compositions:** One-atm experiments on model Martian compositions provide the tightest constraint on $K_{D,Fe_{2+}-Mg}$ because T and oxygen fugacity ($fO_2$) are well known, allowing us to use [6] to predict the amount of ferrous iron present in each liquid. The median of 17 published 1-atm experiments yields $K_{D,Fe_{2+}-Mg} = 0.354 \pm 0.008$ (error is one mean absolute deviation, MAD). There are higher pressure (P) experiments on model Martian compositions, but their $fO_2$s are less well constrained. High-P experiments in graphite capsules are reducing [7], but $Fe^{3+}$ still materially affects estimates of $K_{D,Fe_{2+}-Mg}$; correcting for the $Fe^{3+}$ using [6, 7] increases $K_{D,Fe_{2+}-Mg}$, on average, by 0.016. We see no strong correlations of $K_{D,Fe_{2+}-Mg}$ with liquid or ol composition for the high or low-P experiments and a weak increase in $K_{D,Fe_{2+}-Mg}$ with increasing P. A marked increase in $K_{D,Fe_{2+}-Mg}$ for high-P experiments at $T \sim \sim 1150$ °C is likely due to lack of equilibrium. Accepting high-P experiments with $T > 1150$ °C leads to a median $K_{D,Fe_{2+}-Mg}$ of 0.369 ± 0.024 (MAD), higher than the 1-atm experiments but notably higher than the canonical value of 0.30 [2]. Combining 1-atm. and high-P, high-T experiments yields a median $K_{D,Fe_{2+}-Mg}$ of 0.360 ± 0.013 (MAD), close to the recently reported value of 0.35 ± 0.01 [8], which was obtained without correcting for $Fe^{3+}$ in the liquid.

**Olivine-Phyric Shergottites:** Using $K_{D,Fe_{2+}-Mg} = 0.30$ leads to the result that none of the ol cores in ol-phyric shergottites is in equilibrium with liquids equivalent to their bulk. Applying our $K_{D,Fe_{2+}-Mg}$ to these shergottites leads to the possibility that Y980459, NWA 5789 and 2990 are liquid compositions (others are not), identical to the results of Filiberto and Dasgupta [8]. Accounting for $Fe^{3+}$ in the liquid also allows us to constrain plausible $fO_2$s during cooling when independent measures are unavailable or not yet determined; e.g., if the olivines in NWA 2990 are in equilibrium with a liquid whose composition is that of the bulk meteorite, crystallization must have occurred under reducing conditions where the $Fe^{3+}/\sum Fe$ is small, $\sim IW + 0.5$ at 1 atm or $IW + 1$ at 1 GPa.

**References:** [1] Basaltic Volcanism Study Project. 1981. [2] Roeder P. L. and Emslie R. F. 1970. *Contributions to Mineralogy and Petrology* 29:275-289. [3] Sack R. O. et al. 1987. *Contributions to Mineralogy and Petrology* 96:1-23. [4] Falloon T. J. et al. 2007. *Chemical Geology* 241:207-233. [5] Peslier A. H. et al. 2010. *Geochemistry Geophysics Geosystems* 11:404-412. [6] O'Neill et al. 2006. *American Mineralogist* 93:1838-1844. [8] Filiberto J. and Dasgupta R. 2011. *Earth and Planetary Science Letters* 304:527-537.

EXPLORING XRF AS A NEW TECHNIQUE FOR BASIC METEORITE CLASSIFICATIONS

R. G. Mayne¹, A. J. Ehlmann¹ and K. C. Daviau². ¹Monnig Meteorite Collection, School of Geology, Energy, and the Environment, Texas Christian University, TCU Box 298830, Fort Worth, TX 76129, USA. E-mail: r.g.mayne@tcu.edu. ²Physics Department, Bard College, PO Box 5000, Annandale-on-Hudson, NY 12504, USA.

**Current Classification Methods:** In the 2010/2011 season, the Antarctic Search for Meteorites (ANSMET) collected around 1200 meteorites. These will all be classified over the next 2 years, according to the U.S. Federal Regulation on Antarctic Meteorites [1], by the curatorial staff at the Smithsonian Institution's National Museum of Natural History. The vast majority of the meteorites collected in Antarctica are equilibrated ordinary chondrites and the current method employed for their identification is that of visual examination combined with oil immersion of olivine grains. This technique requires that a small amount of each meteorite be powdered, sieved, and then examined in refractive index oils under a petrographic microscope. This is a timely process and the Smithsonian's main goal is to be able to "separate the approximately 10% of meteorites that deserve further scientific study" [1].

**XRF:** The Monnig Meteorite Collection deals with many requests from the general public for meteorite identification. The vast majority of these samples are industrial slag. The newly purchased Bruker handheld XRF will be used to analyze these samples to confirm that their composition is not meteoritic.

The goal of this study is to explore XRF as a new technique for meteorite identification. In this preliminary study we will present data on the first two parts of this study (1) estimating the Ni-content of iron meteorites (2) examining previously classified ordinary chondrites to establish a calibration technique for unknowns. The eventual aim is to see if the XRF technique outlined here can be used to distinguish LL, L, and H chondrites from one another, to provide an alternative streamlined process to oil immersion.

**Methodology:** All meteorites will be analyzed on the Bruker Tracer III-SD handheld XRF system at Texas Christian University. For part (1) only meteorites from the Monnig Meteorite Collection with a published Ni content were selected for analysis. All Ni contents were taken from [2]. A minimum of three analyses were taken for each sample. The number of counts for Ni $K\alpha$ were tabulated each time and a calibration curve was constructed. Previously classified ordinary chondrites were selected from the Monnig Collection for part (2) of the study. Twenty meteorites of each ordinary chondrite group (H, L, LL) will be measured.

**Preliminary Results:** Initial results relating Ni $K\alpha$ counts to Ni content of iron meteorites show that these two components do appear to be directly related and no further reduction of the data (into percentage Ni) is required. The calibration curve constructed from the preliminary data has an $R^2$ value of 0.94.

At the time of writing, only five H, L, and LL meteorites have been analyzed on the XRF. While there appear to be some differences between the ordinary chondrite groups, there is currently not enough data for statistical analysis to see if these differences are significant enough to allow their identification.

**References:** [1] Corrigan C. M. et al. 2008. *Meteoritics & Planetary Science* 43:A180. [2] Wasson J. T. 2010. *Geochimica et Cosmochimica Acta* 75:1757-1772.
5179

THE 100th ANNIVERSARY OF THE FALL OF NAKHLA: THE SUBDIVISION OF BM1913,25

K. M. McBride¹ and K. Righter². ¹ESCG Jacobs Technology, NASA Johnson Space Center, 2101 NASA Pkwy, Houston, TX 77058, USA. E-mail: kathleen.mcbride-1@nasa.gov. ²NASA-JSC, 2101 NASA Pkwy, Houston, TX 77058, USA.

This year marks the 100th anniversary of the fall of Nakhla, a cumulate clinopyroxenite of Martian origin that fell near Alexandria, Egypt, in 1911. Multiple fragments of the meteorite were seen to fall over an area of 4.5 km in diameter. Approximately 40 stones were recovered with a combined weight of about 10 kg [1]. Most of the larger specimens found their way to museums and meteorite collections in Cairo, Paris, Berlin, and the Smithsonian, to name a few.

In 1998, the British Museum sent a 641 g (BM1913,25), fully fusion crusted stone of Nakhla to the Johnson Space Center (JSC) for processing in the Antarctic Meteorite Lab in order to allocate samples to the scientific community. The stone was split in half in a dry nitrogen glove box. One half of the stone was sent back to the museum and the other half (346 g) was used for sample allocations. From 1998–2001, 37 scientists requested 65 separate samples of Nakhla, including 2 thin sections.

This set of allocations was especially important in that all of the sample splits are from the same piece of Nakhla and it had a known history since it was acquired by the museum in 1913. With the multiple fragments of Nakhla, it is not known from what pieces the main bulk of research has been done, what variation may exist between all the pieces and to what contaminants the fragments may have been exposed, (i.e., water, solvents or cutting fluids, etc.). All of the allocations prepared at JSC were processed in a nitrogen cabinet using only stainless steel, aluminum, and Teflon tools and containers to reduce the chance of introducing any new contaminants.

The focused effort to subdivide and distribute samples of Nakhla to the meteorite community resulted in enhanced understanding of Nakhla and nakhlites in general: organic geochemistry [2, 3], weathering [4], sulfur isotopes [5], radiometric age [6], and magmatic history [7, 8].

There are 13 Nakhlites that have been recovered to date: Nakhla, Lafayette, Governador Valdares, three from NW Africa, three from Yamato and four from Miller Range regions in Antarctica. The Yamato Nakhlites are paired [9] as are the Miller Range samples [10].

**References:** [1] Meyer C. 2004. *Mars Meteorite Compendium* II1–II14. [2] Jull A. J. T. et al. 2000. *Geochimica et Cosmochimica Acta* 64:3763–3772. [3] Glavin D. et al. 1999. *Proceedings of the National Academy of Sciences* 96:8835–8838. [4] Bridges J. and Grady M. M. 2000. *Earth and Planetary Science Letters* 176:267–279. [5] Farquhar J. 2000. *Nature* 404:50–54. [6] Shih C.-Y. et al. 2010. Abstract #1367. 41st Lunar and Planetary Science Conference. [7] Stockstill K. et al. 2005. *Meteoritics & Planetary Science* 40:377–396. [8] Neal C. R. et al. 2001. Abstract #1682. 32nd Lunar and Planetary Science Conference. [9] Imae N. et al. 2005. *Meteoritics & Planetary Science* 40:1581–1598. [10] *Antarctic Meteorite Newsletter* 33, no. 2, 2010.

5040

METEORITE MOVEMENT IN ANTARCTICA AND CLIMATE CHANGE

G. J. H. McCall. Honorary Associate Western Australian Museum (retired); 44 Robert Franklin Way, South Cerney, Glos., UK, GL7 5UD E-mail: joemccall@tiscali.co.uk.

**Introduction:** Large ice sheets lose ice in two ways; slow movement of ice by processes of crystal deformation [1]; and summersurface melting and base flow [2, 3]. The Japanese made the first meteorite finds in the Yamato Mountains, Antarctica, in 1969. Antarctic meteorites can be shown by cosmic ray studies to have fallen thousands, even millions of years ago and are flowing at a very slow rate. Haack et al. [4], who failed to find meteorites on blue (de-oxygenated) ice surfaces in Greenland, suggested that the meteorites were possibly washed down by the melting. Greenland has a higher average temperature than Antarctica. The meteorites may be removed from the surface by the summer melt before they can enter the deep ice flow system; or they may have entered the deep ice as the snow built up and been carried in the blue ice, as in Antarctica, but are not found on the surface. Katabatic winds that are believed to bring meteorites to the surface of blue ice in Antarctica [5], after the very slow flow is impeded by nunataks etc, may not operate in Greenland's climate. The presence or absence of arresting barriers such as nunataks and the amount of snowfall forming ice inland are other possible. Where the blue ice is other possible exposed, impeded and slowed up by nunataks, it is isoclinally folded, and this may also bring meteorites to the surface: also the Greenland ice sheet may have been almost entirely lost in the Eemian interglacial [100,000–70,000 years ago] - the East Antarctic Ice sheet is millions of years old. West Antarctica is losing ice, in East Antarctica there is equilibrium between loss and increment from snowfall inland. The slow movement of meteorites is not affected by the fact that the ice sheet overlies kilometer deep basins, because of the positive topography of the ice surface, which builds up faster than the weight of the ice pushes the land surface down, and in both Greenland and Antarctica the land surface originally (7 and 34 million years ago respectively) had positive topography before the ice formed. Slow movement of the meteorites has no bearing on the arguments about global warming, but surface melt and basal flow does. Because of deep depression of the ice base, rapid surface melt and basal flow can only operate at the coastal margins of the sheet (as is the case in Greenland). The difference in climate between coastal regions and centre of the ice sheet is an important factor also in the distribution of ice loss from surface melting and base flow. Summer melt and basal flow may be insignificant except at the coastal margins in E. Antarctica. The absence in West Antarctica is likely due to climatic or topographic causes.

**References:** [1] Ollier C. 2010. *Geoscientist* 20:16–21. [2] Hambrey M., Bamber J., Christofferson P., Glasser N., Hubbard A., and Hubbard B. 2010. *Geoscientist* 20:18–23. [3] Jenkins M. 2010. *National Geographic Magazine* 217:34–47. [4] Haack H., Schutt J., Meibom A., and Harvey R. 2007. *Meteoritics & Planetary Science* 42:1727–1733. [5] Zolensky M. 1998. In: *Meteorites: Flux with time and impact effects*. Edited by Grady M. M., Hutchison R., McCall G. J. H., and Rothery D. A. Geological Society, London, Special Publication 140:93–104. [6] Naish T. et al. 2009. *Nature* 458:322–328.
ARE WE ALONE? THE EVOLUTION FACTOR IN THE CONTEXT
OF THE EXOPLANET-PROLIFERATED UNIVERSE

G. J. H. McCall. Honorary Associate Western Australian, Museum
(Retired) 44 Robert Franklin Way, South Cerney, GL7 5UD, UK.
E-mail: joemccall@ticali.co.uk

**Introduction:** Recent discussions of the question "Are we alone in the universe" heavily emphasized the astronomical and physical arguments, but seemed largely to ignore the arguments related to the long history of evolution on Earth, and the likelihood or otherwise on exoplanets of any repetition or "mirror image"[1] of the path of evolution, as we know it on Earth, in the habitable zone in the Milky Way galaxy or galaxies more distant. The discovery of about ~500 exoplanets in our galaxy is a scientific revolution of great importance, and these are only "the tip of the iceberg." Myriads more must exist. Yet such is the nature of the variables relating to evolution, the possibility of it taking numerous pathways—even the possibility of different patterns of evolution on other planets in the habitable zone—that it is concluded that there is no chance at all of a repetition of our own planet's evolutionary pathway being repeated on any exoplanet in the habitable zone [2]. Evolution is subject to innumerable constraints, internally within the development of life through the ages, externally, through different geological histories of the host planet and interaction with different coexistent life populations at any point in geological time; and on exoplanets through differences in the parent star, arrangement of the planetary system, and the physical conditions that these impose on the planet: in addition the Earth has the peculiarity of shields such as the magnetosphere and the ozone layer, and the peculiarity of surface liquids in the form of seas (within which much of evolutionary history took place) and a peculiar secondary oxygenic atmosphere, a product of early photosynthetic unicellular life itself. Plate tectonics allows the escape of the Earth's internal heat, and absorbs excess carbon dioxide [3], preventing a Venus-like runaway greenhouse effect, in which advanced life is impossible: Plate tectonics is unique in the solar system, and it may well have appeared mid-way through Earth's life history [4], and allowed eukaryote development? [5-8] It may be unique in the galaxy? The conclusion is reached that, without the likelihood of repetition of the Earth's evolutionary pathways (which may be unique, despite the myriads of exoplanets out there), there will surely be no sufficiently similar intelligent beings out there, who can communicate back to us? Probably the evolutionary pathway on any exoplanet possessed of life is unique. Evolution is essentially a random process. As R. A. Fortey has concluded "There is nothing inexorable about the course of Evolution: rather it is a pas de deux between the changing environment and the capacity of organisms to respond to these changes" [9].

**References:** [1] Gramling C. 2009. *Earth* 54:36–43. [2] Marquardt M. 2010. *Earth* 55:50. [3] Baker D, Ratcliff T. 2010. The 50 most extreme places in the Solar System. Belknap Press of Harvard University Press, London England and Cambridge, Massachusetts 151–155. [4] McCall G. J. H. 2010. Australian Journal of Earth Sciences 57:349–355. [5] Katz L., Bhattacharya, D. 2006. Genomics and the evolution of the microbial eukaryotes. Oxford University Press, 243. [6] Margolis L. 1970. Origin of Eukaryotic Cells. Yale University Press, Newhaven. [7] Keeling P. J. et al. 2006. *TRENDS in ecology and evolution*, 20:670–676. [8] Albani A. et al. 2010. *Nature* 466:100–104. [9] Fortey R. A. (*Geoscientist*, in press).

THE RELATIONSHIP BETWEEN IIE IRONS AND H
CHONDRITES: PETROLOGIC AND OXYGEN ISOTOPE
CONSTRAINTS

K. H. McDermott¹, R. C. Greenwood¹, I. A. Franchi¹, M. Anand¹ and E. R. D. Scott². ¹CEPSAR, The Open University, Milton Keynes MK7 6AA, UK. E-mail: k.h.mcdermott@open.ac.uk. ²HIGP, University of Hawai'i, Honolulu, HI 96822, USA.

**Introduction:** Silicate inclusions have been observed in about half of all currently identified IIE irons (21 specimens) [1]. Based mainly on their similar oxygen isotope compositions, a genetic relationship between the IIE irons and H chondrites had been proposed [2, 3]. However, this relationship was subsequently questioned due to the limited overlap between laser fluorination data for equilibrated H chondrites and the earlier IIE silicate data [4, 5]. We are undertaking a detailed oxygen isotope and geochemical study of both the IIE silicate inclusions and H chondrites, with the aim of investigating further the relationship between the two groups.

**Analytical Techniques:** Oxygen isotope analysis of both mineral separates and bulk silicate inclusions from 9 IIE irons, as well as a suite of 12 H chondrite samples, was undertaken by infrared laser-assisted fluorination [6]. Textural and quantitative analysis of IIE silicate inclusions was carried out using a Cameca SX-100 Electron Microprobe and a FEI Quanta 200 FIB-ASEM.

**Petrographical Results:** Detailed petrographic study of the chondritic IIE-an Netschaevo, (Fa<sub>14</sub>, excluding Fa<sub>25</sub> secondary olivines) reveals that it not only contains well-developed chondrule relicts, but also extensive evidence of partial melting. The more primitive IIE silicates have olivine compositions that fall within the H chondrite range, i.e., Techado (Fa<sub>16</sub>) and Watson 001 (Fa<sub>20</sub>), both of which have also experienced melting.

**Oxygen Isotopic Compositions:** Oxygen isotopes in the IIEs show a range of Δ¹⁷O values from 0.58‰ to 0.90‰, with a mean of 0.72 ± 0.11‰ (2σ). δ¹⁸O for the 9 samples studied also show a considerable range from 3.06‰ to 6.25‰, the mean δ¹⁸O value is 4.69 ± 0.59‰ (2σ). The wide range in δ¹⁸O values reflects the fact that some of the IIE silicate inclusions are highly differentiated (Kodaikanal and Colomera), as discussed by [5]. For comparison, the oxygen isotope analyses of the 12 equilibrated H chondrites examined gave a mean δ¹⁸O value of 4.11 ± 0.56‰ (2σ) and Δ¹⁷O value of 0.71 ± 0.12‰ (2σ).

**Discussion:** The presence of relict chondrules in Netschaevo suggests that the group as a whole may be derived from a chondritic parent body. The nearly identical mean Δ¹⁷O values for the IIE irons and H chondrites support the possibility of a genetic link between these two groups. In addition, the substantial variation in Δ¹⁷O seen in both groups suggests that they were derived from parent bodies with similar levels of primary isotopic heterogeneity. The wide variation in δ¹⁸O values seen in some IIE silicate indicates that they underwent significant degrees of differentiation. In conclusion, our new laser fluorination data show that the H chondrites and IIEs could both have originated on the same unequilibrated H chondrite parent body.

**References:** [1] Mittlefehldt D. W. et al. 1998. *Reviews in Mineralogy* 36. [2] Goldstein J. I. et al. 2009. *Chemie der Erde* 69:293–325. [3] Clayton R. N. and Mayeda T. K. 1996. *Geochimica et Cosmochimica Acta* 60:1999–2018. [4] Folco L. et al. 2004. *Geochimica et Cosmochimica Acta* 68:2379–2397. [5] Franchi I. A. 2008. *Reviews in Mineralogy and Geochemistry*, vol. 68. pp. 345–397. [6] Miller M. F. et al. 1999. *Rapid Communications in Mass Spectrometry* 13:1211–1217.
LASER ABLATION MAPPING OF SIDEROPHILE ELEMENTS IN
THE LATE ARCHAAN REIVOLO IMPACT SPHERULE LAYER,
SOUTH AFRICA

McDonald¹, B. Simonson², C. Koeberl³ and W. U. Reimold⁴. ¹School of Earth & Ocean Sciences, Cardiff University, Cardiff CF10 3YE, UK, E-mail: mcdonald1@cf.ac.uk. ²Oberlin College, Oberlin, OH 44074-1052, USA. ³University of Vienna, Althanstrasse 14, A-1090 Vienna, Austria. ⁴Museum für Naturkunde-Leibniz Institute at Humboldt University Berlin, Invalidenstrasse 43, 10115 Berlin, Germany.

**Introduction:** Layers of spherules formed during late Archaean to early Proterozoic impacts have been recognized in the coeval Hamersley Basin of western Australia and the Griqualand West basin of South Africa [1]. Three spherule layers in the Hamersley Basin (the Jeerinah, Paraburdoo and Dales Gorge layers) can be correlated with equivalents (Monteville, Reivilo and Kuruman) in the Griqualand West Basin [2]. The Paraburdoo-Reivilo layer is remarkable in that it contains an extremely high (up to 40%) extraterrestrial component (ETC), represented by platinum-group element (PGE) contents, similar to some early Archaean spherule layers [3, 4] from Barberton. The high ETC present in the layer allows in-situ trace element mapping of the distribution of the extraterrestrial components within and around impact spherules in order to constrain the processes by which they formed and the manner by which projectile material became incorporated.

**Methods:** Individual spherules (or clusters of spherules) and surrounding matrix from the Reivilo layer were mapped by time resolved laser ablation ICP-MS using a 12 micron laser beam and similar procedures to those described by [5]. The isotopes used for analysis were ²⁴Mg, ²⁷Al, ²⁹Si, ³⁴S, ³⁹K, ⁴⁴Ca, ⁴⁷Ti, ⁵¹V, ⁵²Cr, ⁵⁵Mn, ⁵⁷Fe, ⁵⁹Co, ⁶¹Ni, ⁶⁵Cu, ⁶⁶Zn, ⁷⁵As, ⁹⁵Mo, ⁹⁹Ru, ¹⁰³Rh, ¹⁰⁶Pd, ¹⁰⁷Ag, ¹²¹Sb, ¹²⁵Te, ¹⁴⁰Ce, ¹⁸⁵Re, ¹⁸⁹Os, ¹⁹³Ir, ¹⁹⁵Pt, ¹⁹⁷Au and ²⁰⁹Bi. NIST610 was used as a reference standard and for repeat analyses to monitor drift over the analytical run.

**Results:** Element maps produced to date suggest that Os, Ir, Ru, Rh, and Pt, Ni and Co are confined to the interiors of spherules but also that some spherules are free of siderophiles. The PGE occur most notably as microparticles (nuggets?) that occur close to spherule rims and not evenly through the spherule. Palladium and Au are present at low concentrations and concentrated in the matrix between spherules. Some spherules have been infilled or partially replaced by sulfides (millerite, pyrite, chalcopyrite). These have elevated Ni, Cu, and Co contents, but the sulfides are free of PGE and PGE peaks elsewhere never coincide with elevated S. Other elements, such as Mg, Mn, and Fe, are depleted in the spherules and concentrated in the matrix.

**Discussion:** The presence of spherules with high temperature PGE nuggets and also siderophile-free spherules is significant. If these findings are confirmed by further work, it suggests the spherules formed from a mix of projectile-rich and projectile-free impact melts, not a homogenized cloud of rock vapor [1].

**References:** [1] Simonson B. et al. 2009. *Precambrian Research* 169:100-111. [2] Hassler S. W. et al. 2011. *Geology*, 39:307-310. [3] Simonson B. et al. 2009. *Precambrian Research* 175:51-76. [4] Kyte F. et al. 1992. *Geochimica et Cosmochimica Acta* 56:1365-1372. [5] Ulrich T. et al. 2009. *Canadian Mineralogist* 47:1001-1012.

GEOCHEMISTRY OF POST-KINEMATIC MAFIC INTRUSIONS AT
MANIITSOQ, WEST GREENLAND: CONTAMINATED MANTLE
MELTS OR IMPACT MELTS?

McDonald¹ and A. A. Garde². ¹School of Earth & Ocean Sciences, Cardiff University, Cardiff CF10 3YE, U.K., E-mail: mcdonald1@cf.ac.uk. ²Geological Survey of Denmark and Greenland, Øster Voldgade 10, 1350 Copenhagen K, Denmark.

**Introduction:** The ~2.975 Ga Maniitsoq structure in West Greenland may be a Mesoarchaean impact structure [1]. It is centred on a ~200 km wide aeromagnetic anomaly and contains a near-circular, > 40 km wide homogenised granitoid core with a diffused boundary, surrounded by a still wider belt of shattered and brecciated country rock with variably deformed and downgraded quartz PDFs, dykes of fluidised microbreccia, abundant evidence of partial melting, and massive hydrothermal alteration. A 70 km long contemporaneous curvilinear belt of unshocked norite-diorite sills and diorite plugs with proto-orbicular textures, indicative of rapid cooling of superheated melt, occurs east of the core. The norites and diorites are geochemically diverse [2], with zones of sulfide mineralization and the aim of this study has been to test whether any could be downward injections of impact melt below the crater in the manner of offset dykes at Sudbury [3].

**Results:** Analyses for major elements, trace elements and platinum-group elements (PGE) were performed by ICP and ICP-MS. The norites-diorites display an enormous range of MgO (5.9-23.7wt.%) and Cr (120-3600 ppm) over a limited range of SiO₂ (47-59wt.%). On Harker plots the intrusions fall on a general array, between primitive mantle melts and local gneisses/amphibolites, that is more extensive, dispersed and shows a stronger mantle component than typically found with impact melts [3]. Unmineralized high-Mg norites contain low concentrations of PGE (Ir = 0.03-0.2 ppb). Mineralized norites may carry up to 10 ppb Ir but sulfide metal tenors are low and high Cu/Pd ratios (10⁴-10⁶) indicate removal of PGE by immiscible sulphides before final emplacement. Norites and diorites typically display strongly suprachondritic Ru/Ir, Rh/Ir, Ru/Rh and Pt/Ir typical of mantle melts. However, one low-Mg diorite (283761) south of the structure core plots off the main compositional array, and contains the highest concentration of Ir (0.59 ppb), along with chondritic Ru/Ir, Rh/Ir and Ru/Rh ratios.

**Discussion:** The evidence suggests that most of the 2.975 Ga norite-diorite intrusions are not impact melts but represent contaminated mafic magma, intruding upwards from the mantle below the structure, possibly in response to post-impact decompression. However one intrusion sampled to date may represent impact melt injected downwards from a former overlying melt sheet.

**References:** [1] Garde A. A. 2010. Abstracts and Proceedings of the Geological Society of Norway 1:57-58. [2] Garde A. A. 1997. Geological Survey of Greenland Bulletin 177:115. [3] Wood C. R. and Spray J. G. 1998. *Meteoritics & Planetary Science* 33:337-347.
5362

DOES VESTA HAVE MOONS?

L. A. McFadden¹, M. Sykes², S. Joy³, P. Tricarico², D. O'Brien², J. Y. Li⁴, M. Mutchler⁵, Nargess Memarsadeghi¹, H. Safavi¹, P. Gutierrez-Marques⁶, A. Nathues⁶, S. Mottola⁷, H. Sierks⁶, S. Schroder⁶, C. Polansky⁸, R. Jacobson⁸, C. T. Russell³, C. A. Raymond⁸, M. Rayman⁸, S. Weinstein-Weiss⁸, E. Palmer². ¹NASA, Goddard Space Flight Center. E-mail: lucy.mcfadden@nasa.gov. ²Planetary Science Institute, Tucson, USA. ³IGPP, UCLA. ⁴U. Maryland. ⁵Space Telescope Science Institute. ⁶Max-Planck-Institut für Sonnensystemforschung, Lindau, Germany. ⁸California Inst. Technology, Jet Propulsion Laboratory.

Previous searches for moons around Vesta have found nothing to an upper limit of 22.5 magnitude, that corresponds to $44 \pm 4$ m diameter assuming the same albedo as Vesta. The Dawn mission's approach phase has dedicated satellite search observations consisting of two mosaic sequences bracketing the first observations of a complete rotation of Vesta scheduled for early July, 2011. In addition, we use the approach optical navigation image sequences for initial satellite searches. We will report any findings from these observations, and upper limits of magnitude and size.

5051

NAKHLA: A MARTIAN METEORITE WITH INDIGENOUS ORGANIC CARBONACEOUS FEATURES

D. S. McKay¹, E. K. Gibson¹, K. L. Thomas-Keprta², S. J. Clemett², L. Le², Z. Rahman² and S. J. Wentworth². ¹ARES, NASA/JSC, Mail Code KA, Houston, TX 77058, USA. ²ESCG at NASA/JSC, Mail Code JE23, Houston, TX 77058, USA. E-mail: david.s.mckay@nasa.gov.

The Nakhla meteorite possesses discrete, well defined, structurally coherent morphologies of carbonaceous phases present within iddingsite alteration zones. Based upon both isotopic measurements and analysis of organic phases the presence of pre-terrestrial organics is now recognized. Within the microcrystalline layers of Nakhla's iddingsite, discrete clusters of salt crystals are present. These salts are predominantly halite (NaCl) with minor $MgCl_2$ crystals. Some $CaSO_4$, likely gypsum, appears to be partially intergrown with some of the halite. EDX mapping shows discrete C-rich features are interspersed among these crystals.

A hollow semi-spherical 'bowl' structure (~3 µm) has been identified and analyzed after using a focused ion beam (FIB) to cut a transverse TEM thin section of the feature and the underlying iddingsite. TEM/EDX analysis reveals that the feature is primarily carbonaceous containing C with lesser amounts of Si, S, Ca, Cl, F, Na, and minor Mnand Fe; additionally a small peak consistent with N, which has been previously seen in Nakhla carbonaceous matter, is also present. Selected area electron diffraction (SAED) shows that this C-rich material is amorphous (lacking any long-range crystallographic order) and is not graphite or carbonate. Micro-Raman spectra acquired from the same surface from which the FIB section was extracted demonstrate a typical kerogen-like D and G band structure with a weak absorption peak at 1350 and a stronger peak at 1600 cm⁻¹. The C-rich feature is intimately associated with both the surrounding halite and underlying iddingsite matrix.

Both iddingsite and salts are interpreted as having formed as evaporate assemblages from progressive evaporation of water bodies on Mars. This assemblage, sans the carbonaceous moieties, closely resembles iddingsite alteration features previously described which were interpreted as indigenous Martian assemblages. These distinctive macromolecular carbonaceous structures in Nakhla may represent one of the sources of the high molecular weight organic material previously identified in Nakhla.

While we do not speculate on the origin of these unique carbonaceous structures, we note that the significance of such observations is that it may allow us to construct a C-cycle for Mars based on the chemistry of the Martian meteorites with obvious implications for astrobiology and the prebiotic evolution of Mars. In any case, our observations strongly suggest that organic C exists as micrometer-size, discrete structures on Mars.

**Acknowledgment:** Receipt of the Nakhla samples from Caroline Smith and the British Museum Natural History is acknowledged.
5113

Mn-Cr SYSTEMATICS OF PALLASITE OLIVINE BY SIMS:
CONTRASTING Cr IN BRAHIN AND BRENHAM

S. J. McKibbin, T. R. Ireland, Y. Amelin, P. Holden and P.
Lanc. Research School of Earth Sciences, Australian National
University, Australia. E-mail: seann.mckibbin@anu.edu.au.

**Introduction:** Pallasite meteorites are mixtures of Fe-Ni metal and olivine and would appear to be good candidates for $^{53}$Mn-$^{53}$Cr dating (t$_{1/2}$ 3.7 Ma). However, conflicting ion-probe results have been reported for these rocks, with some workers finding elevated $^{53}$Cr/$^{52}$Cr correlated with $^{55}$Mn/$^{52}$Cr [1, 2] while others do not [3]. Recently, it has been suggested that a statistical effect at low count rates arising from the reduction of data as a series of ratios could introduce positive bias into isotope measurements for some applications [4], leading [5] to retract earlier findings of elevated $^{53}$Cr/$^{52}$Cr in pallasite olivine [1, 2]. To investigate whether in-situ accumulated radiogenic $^{53}$Cr exists in these meteorites, we have re-analyzed olivine from two petrographically contrasting pallasites: Brahin (which has a fragmental texture), and Brenham (with rounded olivine).

**Methods:** In situ isotopic analysis was performed using sensitive high-mass resolution ion microprobe reverse geometry (SHRIMP-RG) on olivine and chromite from Brahin and olivine from Brenham. The primary sputtering beam was run at high current (17-66 nA as a mixture of O⁻ and O₂⁻) and focused to a ~30 micron diameter spot, giving strong Cr secondary ion beams and hence minimizing the effect described by [4]. $^{52}$Cr⁺, $^{53}$Cr⁻ and $^{55}$Mn⁺ were measured with an electron multiplier, and San Carlos olivine was used as an external standard to correct for instrumental fractionation on $^{53}$Cr⁺/$^{52}$Cr⁺ and $^{55}$Mn⁺/$^{52}$Cr⁺.

**Results:** We find that olivine and chromite in Brahin has normal Cr isotopic composition within analytical error (~1% 1σ). In contrast, Brenham olivine shows Cr-isotope excesses (up to 6 ± 1% 1σ) within a few 10's of microns of olivine rims. These are well correlated with $^{55}$Mn/$^{52}$Cr and imply an initial $^{53}$Mn/$^{55}$Mn of 2.55 (±0.62) × 10⁻⁵ at the time of isotopic closure of Cr in Brenham olivine rims.

**Discussion:** Our finding of isotopically normal Cr in Brahin is in agreement with [3], while our initial $^{53}$Mn/$^{55}$Mn for Brenham is even higher than originally reported by [1, 2]. Previous ion-probe studies had 30–250 counts per cycle for $^{52}$Cr⁺ over 150 cycles [5], but with higher beam currents and improved ion transmission we have achieved 1–4 × 10⁴ counts per cycle, with 80–200 cycles for each analysis. We estimate bias arising from the effect highlighted by [4] to be ~0.1% for our measurements.

The contrasting Mn-Cr systematics found here is good evidence for localised Cr-isotope variation in some, but not all, pallasites. The initial $^{53}$Mn/$^{55}$Mn found for Brenham is not compatible with current understanding of extinct nuclide chronology (with solar system initial $^{53}$Mn/$^{55}$Mn of 9.1 (±1.7) × 10⁻⁶; [6]) and probably represents concentration of radiogenic $^{53}$Cr at olivine rims. This appears to be related to the arrangement of olivine, with the interconnected networks in Brenham [7] exhibiting anomalous Cr while isolated olivine fragments in Brahin do not.

**References:** [1] Hsu et al. 1997. Abstract #1783. 28th Lunar and Planetary Science Conference. [2] Hsu W. 2005. *Geochemical Journal* 39:311. [3] Tomiyama T. et al. 2007. Abstract #2007. 38th Lunar and Planetary Science Conference. [4] Ogliore et al. 2011. Abstract #1592. 42nd Lunar and Planetary Science Conference. [5] Huss et al. 2011. Abstract #2608. 42nd Lunar and Planetary Science Conference. [6] Nyquist et al. 2009. *Geochimica et Cosmochimica Acta* 73:5115–5136. [7] Spinsby et al. 2008. *Computers & Geosciences* 34:1–7.

5114

SIMS Mn–Cr SYSTEMATICS OF OLIVINE FROM THE ANGRITE METEORITE D'ORBIGNY: AN INTER-LABORATORY COMPARISON

S. J. McKibbin¹, T. R. Ireland¹, Y. Amelin¹, N. Sugiura², P. Holden¹ and P. Lanc¹. ¹Research School of Earth Sciences, Australian National University, Australia. E-mail: seann.mckibbin@gmail.com. ²Department of Earth and Planetary Science, University of Tokyo, Tokyo, Japan.

**Introduction:** In recent years, the angrite meteorites have proven suitable for dating by the $^{53}$Mn-$^{53}$Cr extinct nuclide system using *in situ* techniques [1–3]. However, analyses of the D'Orbigny angrite using two secondary ion mass spectrometers (SIMS) [1,2] revealed a significant inter-laboratory bias. The initial $^{53}$Mn/$^{55}$Mn for this meteorite is reported as 2.84(±0.24) × 10⁻⁶ by [1] and ~3.5 × 10⁻⁶ [2]. To test whether heterogeneous Mn–Cr systematics or differences in methodology caused this bias, we have re-analyzed the piece of D'Orbigny from [1] using the method of [2].

**Methods:** The D'Orbigny sample was mounted inside a circular standard holder containing reference materials to monitor instrumental mass and inter-element fractionation. We have previously developed a method for in situ measurement of Mn–Cr systematics in olivine using Sensitive High-mass Resolution Ion Micro Probe (SHRIMP; models RG and II) which we apply here. Mn–Cr relative sensitivity was obtained from synthetic Fe- and Ca-rich olivines [2]. Our correction for relative sensitivity is in the range of 0.54–0.80 depending on the matrix.

**Results:** As previously reported, D'Orbigny olivine is strongly zoned in major and trace elements, and we find a good correlation between $^{55}$Mn⁺/$^{52}$Cr⁺ and $^{53}$Cr⁺/$^{52}$Cr⁺. After correction for inter-element sensitivity (using a three-component mixing model to account for the continuously varying major-element matrix) and mass fractionation of Cr-isotopes, the gradient of our regression gives an initial $^{53}$Mn/$^{55}$Mn of 3.61 (±0.39) × 10⁻⁶.

**Discussion:** Our new result agrees with our previous finding for another piece of D'Orbigny [2], but disagrees with previous work on this chip at the University of Tokyo [1]. This indicates that the discrepancy is not caused by partial resetting of the Mn–Cr system or presence of cosmogenic isotopes. The Mn–Cr relative sensitivity factors of 0.54–0.80, determined in this study, do not seem to be responsible either, since they are similar to that of 0.617 ± 0.046 reported in [1]. Potential biases arising during data reduction of low count rate measurements [4,5] could influence Cr-isotope ratios of low-Cr olivine in angrites, but rejection of analyses with the lowest Cr concentrations (and hence highest Mn/Cr) does not change the gradient of our regression to systematically lower values (e.g., for analyses with $^{55}$Mn/$^{52}$Cr < 10⁵, the same range as in [1], we find initial $^{53}$Mn/$^{55}$Mn of 3.63(±0.47) × 10⁻⁶). We are currently investigating the effects of different data reduction methods. The discrepancy might also be related to differing styles of instrumental mass fractionation in CAMECA and SHRIMP ion probes, in which case the two-isotope external correction method may not accurately account for Cr-isotope fractionation in one or both of these instruments.

**References:** [1] Sugiura N. et al. 2005 *Earth Planets and Space* 57: e13–e16. [2] McKibbin S. J. et al. 2008. Abstract #5132. 71st MetSoc. [3] McKibbin S. J. et al. 2010. Abstract #5172. 73rd MetSoc. [4] Ogliore R. C. et al. 2011. Abstract #1592. 42nd Lunar and Planetary Science Conference. [5] Huss G. R. et al. 2011. Abstract #2608. 42nd Lunar and Planetary Science Conference.
## DAWN ARRIVES AT VESTA

H. Y. McSween¹, C. T. Russell² and C. A. Raymond³. ¹Department of Earth and Planetary Sciences, University of Tennessee, Knoxville, TN 37996-1410, USA. ²Institute of Geophysics and Planetary Physics, University of California, Los Angeles, CA 90095-1567, USA. ³Jet Propulsion Laboratory, Pasadena, CA 91109, USA.

The Dawn spacecraft, launched in 2007 and fueled by ion engines, will perform a rendezvous with asteroid 4 Vesta, a differentiated body thought to be the source of the HED meteorites. Dawn will spend roughly a year in Vesta orbit studying its surface and gravity field to determine its geologic history and to establish its relationship with the HEDs. The Framing Camera (FC) will obtain image mosaics in multiple color filters at < 100 m px⁻¹ spatial resolution, and partial coverage at 25 m px⁻¹ resolution, to map the geology, geomorphology, and tectonic fabric, and to characterize space weathering. Multi-angle clear filter images will be used to construct a topographic model to 100 m spatial resolution and 10 m height accuracy. The Visible and InfraRed Spectrometer (VIR) maps spectral variations between 0.25 and 5 µm to resolve mineralogy, geologic units, weathering effects and thermal brightness variations. The elemental composition of the surface is resolved to 10 spatial pixels by the Gamma Ray and Neutron Detector (GRAND). Together with the spherical harmonic gravity field derived from navigation, these observations will provide a comprehensive data set to investigate Vesta's differentiation, magmatic and tectonic history, effects of impacts, and weathering of the surface. During approach, rotational characterizations will be performed to update Vesta's spin pole, finalize instrument calibrations, and tune exposure times for imaging instruments.

At the time of the 2011 METSOC meeting, Dawn will have just entered its Survey orbit, observing Vesta for seven orbits over 21 days at a mean radius of 3000 km, while the Sun is at its highest southern latitudes. After achieving global color and multispectral maps with FC and VIR, the spacecraft transfers to the High-Altitude Mapping Orbit (HAMO, mean radius 950 km). In HAMO FC performs detailed global mapping in multiple color filters, while VIR collects high-resolution multispectral data in the southern hemisphere. After 30 days in HAMO, Dawn transfers to the Low-Altitude Mapping Orbit (LAMO, mean radius 460 km), where GRAND maps Vesta's surface for nominally 70 days, while daily tracking passes map the gravity field. Following LAMO, Dawn spirals out and stops at HAMO altitude to map the newly illuminated terrain in the northern polar region. Departure for Ceres is scheduled for July 2012 to meet an arrival in February 2015.

## ANALYSIS OF ALMAHATA SITTA METEORITE

H. T. Meadni¹ and G. M. Taha² ¹Student Department of Physics.
E-mail: haboyaasl@yahoo.com. ²Student/department of Chemistry.
E-mail: ghadamhmd11@yahoo.com.

Fragments of Almahata Sitta meteorite was analyzed by a team of undergraduate students at the departments of Chemistry and Physics at University of Khartoum, as an attempt to obtain chemical information to classify the meteorite chemically. Sample #27 was analyzed using Atomic Absorption spectroscopy for the elemental analysis², the results shows various amounts of cobalt, calcium, magnesium, zinc, manganese, and high percentage of iron up to (44.6%). And for the identification of amino acids the high performance liquid chromatography is used to detect trace amounts of six different amino acids.

Also a fragment sample#43 was analyzed using X-ray diffraction technique in order to obtain structural information and the XRD data was studied using WINFIT and POWDERCELL software. The estimated crystallite size of different phases of the sample was found to range between 20 nm to 60 nm.

The plasma formation technique is used to detect the light elements in Almahata Sitta meteorite. This was investigated by using high voltage pluses which produce high pluses current. The electrons beam was focused by external magnetic field to defuse electron on target surface geometry with pulse 177 ns. Brightness of plasma emission and spectroscopic analysis show different content of some meteorite samples. The spectroscopic absorption pattern indicates the existence of some light elements in Almahata Sitta meteorite such as (Cu, N, F, I, Kr) in addition to Ne and Pm.

**References:** [1] SHADDAD M. H. et al. 2010. *Meteoritics & Planetary Science* 45:1553–1556. [2] Skoog D. A. and West D. M. 1980. *Principles of instrumental analysis*, second edition.
MAKING THE MOON FROM THE EARTH'S MANTLE—THE GEOCHEMICAL PERSPECTIVE

M. M. Meier¹, A. Reufer², W. Benz² and R. Wieler¹. ¹Earth Sciences, ETH Zurich, Switzerland. E-mail: meier@erdw.ethz.ch. ²Physics Department, University of Berne, Switzerland.

**Introduction:** The formation of the Moon from a circum-terrestrial disk ejected in a collision between the proto-Earth and a roughly Mars-sized impactor (a.k.a. the giant impact or GI hypothesis [1]) is widely accepted today. However, the close geochemical similarity of lunar rocks and Earth's mantle has posed a serious problem for the hypothesis, as simulations of the GI consistently [1] had the Moon forming predominantly (~80%) from impactor material. In view of large isotopic and elemental inhomogeneities in different solar system objects, e.g., in O, Si, Cr and W isotopic ratios, or Mg/Cr and Hf/W elemental ratios, it seems unlikely that the impactor had a composition closely matching that of the Earth. Similarity in W isotopes and Hf/W ratios further implies that proto-Earth and impactor must have formed their cores at the same time. A post-impact re-equilibration model has been proposed to account for the similarity in O isotopes [2], but so far, no complete, quantitative model of this exists [3]. It therefore remains unclear at this stage whether this re-equilibration model can satisfy all geochemical constraints. Recently, a modified GI scenario with new parameters was presented, where a steep, high-velocity collision yields an iron-depleted disk mainly derived from terrestrial (or "target") material and a reasonable Earth-Moon angular momentum [4]. Here, we compare the predictions from this model with constraints from the above-mentioned geochemical systems.

**Results and Discussion:** The geochemical similarity of Earth and Moon can be expressed in a factor $\delta f_T = (M_{\text{target}} / M_{\text{total}})_{\text{disk}} / (M_{\text{target}} / M_{\text{total}})_{\text{post-impact Earth}} - 1$, reflecting the depletion of terrestrial material in the Moon-forming disk. Since both Earth and Moon incorporate significant amounts of impactor material, it is this difference in relative contributions that has to approach zero to satisfy geochemical constraints. The exact value of $\delta f_T$ depends on the composition of proto-Earth and impactor, but also on their initial mass and the mass lost from the system. Typical values for $\delta f_T$ in the "canonical" model [1] do not exceed -66%, incompatible with geochemical constraints. For an impactor that resembles Mars in O, Si and Cr isotopic composition, $\delta f_T$ should exceed -5%, -59% and -45%, respectively. A typical outcome of the new high-velocity GI model is a $\delta f_T$ of -34%, compatible with these constraints except for O. However, O may be a special case, as there are at least three different solar system materials (Earth, E- and CI-chondrites) that plot on the terrestrial fractionation line, possibly because they all sample the same reservoir, from which also the impactor could have been derived. Alternatively, post-impact re-equilibration as proposed by [2][3] may have erased any remaining isotopic differences.

**Conclusions:** The high-velocity GI model [4] can satisfy the known geochemical constraints. A possible exception is O, if the impactor is as different in O isotope composition as Mars. In this case, post-impact re-equilibration [2] is still needed, for which our model constitutes a new starting point, as the high-velocity GI model also predicts a much hotter early Earth. This may have implications for geophysics and the formation of the atmosphere.

**References:** [1] Canup et al. 2001. *Nature* 412:708. [2] Pahlevan and Stevenson. 2007. *Earth and Planetary Science Letters* 262:438. [3] Pahlevan et al. 2011. *Earth and Planetary Science Letters* 301:433. [4] Reufer et al. 2011. Abstract #1136. 42nd Lunar and Planetary Science Conference.

THERMAL HISTORY OF VIG1623-5 FUN CAI: EXPERIMENTAL MODELING

R. A. Mendybaev and F. M. Richter. University of Chicago, IL, USA. E-mail: ramendyb@uchicago.edu.

**Introduction:** Vig 1623-5 is a FUN CAI composed mainly of 30-170 µm grains of forsterite (Fo) enclosed within pyroxene (Px) and/or äkermanitic (Åk₈₉) melilite (Mel) and surrounded by incomplete gehlenitic (Åk₁₆₋₆₀) melilite mantle [1]. The inclusion is characterized by extremely large mass-dependent fractionations of magnesium, silicon and oxygen isotopes [1, 2]. Earlier we showed [3] that isothermal evaporation of FUN1 melt (Mg₂SiO₄ plus several wt% of CaO and Al₂O₃) at 1900 °C results in residues with chemical and isotopic compositions of Mg, Si and O close to those in Vig1623-5. In this paper we present results of cooling experiments aimed at reproducing the texture of 1623-5.

**Experiments:** Most experiments have been conducted in a vacuum furnace using FUN1 and FUN2 melt compositions; the latter is a product of evaporation of FUN1 and has a composition close to 1623-5 (see [3] and references therein). A one-atmosphere furnace was used in experiments when FUN2 melt was cooled from temperatures below 1600 °C in H₂.

**Results:** Cooling FUN1 from 1850 °C in vacuum at 30-120 °C h⁻¹ results in residues composed of Fo and Px. Contrary to the crystallization in a closed system, fast cooling (~120 °C h⁻¹) leads in formation of massive Fo (up to 1 mm in size) while slow cooling (~30 °C h⁻¹) results in "quench-like" textures. This is due to the fact that as Mg and Si evaporate, the liquidus of the melt decreases, but when the evaporation rate becomes sufficiently slow, instant nucleation and crystallization of Fo begins. These experiments show that although evaporation of Mg and Si from FUN1 melt results in residues appropriately enriched in heavy Mg and Si isotopes [3], the texture and mineral compositions of the residues are distinctly different from that of Vig1623-5.

Cooling of FUN2 melt in vacuum from ~1650 °C at 30-60 °C h⁻¹ results in Fo (up to 200-300 µm) embedded in Px and Mel. Small Sp grains and some An are also observed in the residues. The compositions of forsterite (up to 1.5 wt% CaO) and the Ca-, Al-rich pyroxene in the residues are close to that of Vig1623-5 [1]. Melilite from these experiments, however, is more gehlenitic (~Åk₇₅) than the Åk₈₉ found in the main mass of Vig1623-5.

We showed that when Type B melt is exposed to H₂ melilite mantle forms [4]. Similar experiments are underway using FUN2 melt. Run products from the vacuum experiments described above were reheated into one atm H₂ and cooled from T ≤ 1600 °C. We have not yet succeed in forming a melilite mantle, but cooling from 1425 °C at 10 and 25 °C h⁻¹ results in 10-50 µm thick forsterite-free mantle (5-10 µm Sp grains are enclosed into Mel and Px). The expectation based on the phase diagram is that slower cooling (or slightly higher starting temperatures) will evaporate sufficient Mg and Si so that Px would convert into Mel followed by dissolution of Sp in the gehlenitic Mel. The last step in this process was observed in evaporation experiments with Type B-like compositions [4].

**Conclusions:** Formation of texture similar to Vig1623-5 FUN CAI seems to require multiple heating and cooling events.

**References:** [1] Davis et al. 1991. *Geochimica et Cosmochimica Acta* 55:621-637. [2] Marin-Carbonne et al. 2011. 74th Meteoritics Society Meeting. [3] Mendybaev et al. 2010. Abstract #2725. 41st Lunar and Planetary Science Conference. [4] Mendybaev et al. 2006. *Geochimica et Cosmochimica Acta* 70:2622-2642.
5145

μ-FTIR, μ-RAMAN SPECTROSCOPIES AND SEM-EDX ANALYSES OF GRAINS OF THE PARIS METEORITE

S. Merouane¹, Z. Djouadi¹, R. Brunetto¹, J. Borg¹ and P. Dumas². ¹IAS, UMR 8617 Orsay, France. ²Synchrotron SOLEIL, Saint-Aubin, France. E-mail: sihane.merouane@ias.u-psud.fr.

**Introduction:** We investigated a few fragments from the Paris meteorite through mid- and far-infrared micro-spectroscopy (2.5–100 µm), Raman spectroscopy and SEM-EDX analyses to get information on their mineralogy and composition (organics and elementary). These data allow us to compare Paris (at least these fragments if representative of the meteorite) with other carbonaceous chondrites in particular, and, more generally, with other primitive extraterrestrial objects present in our collections (e.g., IDPs for Interplanetary Dust Particles).

**Experiments:** Few micron-sized fragments among the 13 mg received from the MNHN (Paris, France) were crushed into a Diamond compression cell [1]. Infrared spectra were obtained on the SMIS2 beamline of the SOLEIL synchrotron (France) with a NicPlan microscope attached to a Fourier transform infrared spectrometer (FTIR). Raman spectra were acquired with a spectrometer DXR from Thermo Fisher with a 532 nm laser and EDX analyses were performed with a SEM Hitachi 3600N and an EDX spectrometer ThermoNoran System SIX at IEF (Orsay, France).

**Results and Discussion:** Mineralogy Infrared spectra are dominated by phyllosilicates and carbonates, implying a significant aqueous alteration, most probably on the parent body. Far-infrared and Raman spectra reveal the presence of magnetite and traces of sulfides, possibly troilite. Only one inclusion of crystalline anhydrous silicate (mixture of magnesium-rich olivine and pyroxene) was found in the analyzed fragments.

*Organic Matter* The study by IR spectroscopy of the aliphatic organic matter present in the meteorite leads to a CH₂/CH₃ ratio of 2.3 ± 0.2, which is in good agreement with the reported values in Orgueil (CI) and Murchison (CM) [2]. The aromatic carbon (detected by Raman spectroscopy) has a quite ordered structure indicating that these grains have suffered a weak thermal alteration and seem to be less primitive than IDPs [3].

*Elementary Composition* Some elementary ratios were extracted from the EDX data analyses. The obtained values Mg/Si = 0.88, Fe/Si = 0.95, S/Si = 0.42, Al/Si = 0.091, compatible with the chondritic values reported by Anders and Grevesse [4], in the range of CI and CM elementary ratios [5, 6]. However, Ni/Si = 0.078 and Ca/Si = 0.25 are higher than for the CI and CM chondrites. Our analytical techniques show an excellent correlation between the spatial distribution of the carbonates (found by IR spectroscopy) and element Ca (as shown by EDX mapping) implying that major part of the Ca is incorporated in the carbonates.

**Conclusion:** The combination of different analytical techniques enables both qualitative and quantitative studies of the samples. Microspectroscopy is also a useful tool to determine the distribution of the different components in micron-sized grains.

The different analyses of the few fragments indicate that they have undergone a significant aqueous alteration and a weak thermal alteration. Both their mineralogical and organic composition, as well as their aromatic component, classify these grains of the Paris meteorite close to the CI or CM meteorite-type.

**References:** [1] Grossemy F. 2008. Ph.D. thesis [2] Matrajt G. et al. 2005. *Astronomy & Astrophysics* 433:979–995 [3] Raynal P. I. 2003. Ph.D. thesis. [4] Anders E. and Grevesse N. 1989. *Geochimica et Cosmochimica Acta* 53:197–214. [5] Brownlee D. E. 1978. *IAU Colloq.* 52:134–150 [6] Hudson G. J. et al. 1981. *Science* 211:383–386.

5146

QUANTITATIVE STUDY OF THE 10 µM SILICATE BAND IN INTERPLANETARY DUST PARTICLES

S. Merouane¹, Z. Djouadi¹, R. Brunetto¹, J. Borg¹ and P. Dumas². ¹IAS, UMR 8617, Orsay, France. ²Synchrotron SOLEIL, Saint-Aubin, France, E-mail: sihane.merouane@ias.u-psud.fr.

**Introduction:** Interplanetary dust particles (IDPs) are particles originating from an asteroidal or cometary parent body. Their mineralogical composition can give clues on their parent bodies' characteristics and thus on the formation of the solar system [1].

In this study we analyze a few IDPs of the NASA collection with complementary analytical techniques (Infrared (IR) micro-spectroscopy, Raman spectroscopy and SEM-EDX analyses) to get information on their mineralogy, their organics and their elementary composition.

**Experiment:** The studied IDPs were crushed into a Diamond'EXPress compression cell following the protocol of [2, 3]. IR micro-spectroscopy from mid to far (2.5–100 µm) infrared is performed on the SMIS2 beamline of the SOLEIL synchrotron (France) using a NicPlan microscope attached to a Fourier Transform infrared spectrometer (FTIR). Raman spectra were acquired at SOLEIL with a spectrometer DXR from Thermo Fisher with a 532 nm laser and SEM-EDX analyses were performed with a SEM Hitachi 3600N and an EDX spectrometer. ThermoNoran System SIX at "l'Institut d'Electronique Fondamentale" (IEF, Orsay, France).

**Procedure:** The present work focuses on the mineralogical composition of the grains, and more particularly on the silicate signature, given in IR by the 10 µm feature, corresponding to the well-known stretching mode of the Si-O bonds. We are developing an analytical procedure consisting on a "deconvolution" of this band in order to quantify its different contributors (olivine, pyroxene, as well as other possible species, either crystalline or amorphous). These contributors are identified by combining the obtained results with the different analytical techniques (IR, Raman and SEM-EDX).

Silicate-type spectra were first acquired for standard minerals (olivine, pyroxene). We used crystalline minerals crushed in a diamond cell (as the same manner as the IDPs) and amorphous thin films synthesized by electron beam evaporation from the crystalline silicates [4] on a KBr or CsI substrate. Using a linear combination of the standard spectra (amorphous and crystalline), we are able to fit the 10 µm band of the IDPs.

We will present the first results obtained on the validation of our analytical procedure on synthetic samples and its application on IDPs. This work is still under improvement but this technique seems promising to infer the quantitative mineralogical composition of IDPs and thus give clues on the processes undergone by their parent bodies.

**References:** [1] Sandford S. A. and Bradley J. P. 1989. *Icarus* 82:146–166. [2] Grossemy F. 2008. Ph.D. thesis, France. [3] Brunetto R. et al. 2011. *Icarus* 212:896–910. [4] Djouadi Z. et al. 2005. *Astronomy & Astrophysics* 440:179–184.
5471

FIRST ANALYSIS OF ALL XENON ISOTOPES IN GENESIS SOLAR WIND AloS COLLECTOR

A. Meshik¹, O. Pravdivtseva¹, C. M. Hohenberg¹, J. H. Allton² and D. S. Burnett³. ¹Washington University, Physics CB1105, Saint Louis, MO 63130. E-mail: am@physics.wustl.edu. ²NASA, Johnson Space Center, Houston. ³Geology 100-23, CalTech, Pasadena.

**Introduction:** Recently we made successful measurements of all solar wind (SW) krypton isotopes in Genesis Al-collectors [1]. Here we report results of 20 independent analysis of Xe extracted from Aluminum on Sapphire (AloS) under four different experimental conditions.

**Experimental:** Separation of the solar wind heavy noble gases, captured in Genesis collectors, from surface and possible interface contaminations is not trivial. One way of dealing with these contaminations is depth profiling. But due to low SW-Xe abundance it requires an extremely sensitive mass spectrometer, such as RELAX [2]. In this work we developed a different approach, extracting total Xe and Kr and making a blank correction on the basis of measured Kr/Xe ratios which are different in the SW and the blank. The later turned out to be not atmospheric and was determined by minimizing the spread of calculated SW-Xe fluencies between different extracted areas.

**Results:** SW fluence and preliminary isotopic composition of Xe extracted from AloS at four different conditions are shown in the table below, where α – angle between incident laser beam and the plane normal to the surface of Al-film; n – number of independent laser extractions.

**Conclusion:** Our preliminary Genesis Xe results are generally agree with SW-Xe obtained from the lunar soil [3,4]. Some subtle differences need to be further investigated. We plan to verify our data treatment and possibly develop some rejection criteria since in several cases the data scatter exceed statistically expected deviations.

**Acknowledgment:** Supported by NASA grant NNX07AM76G.

**References:** [1] Meshik A. P. et al. 2011. Abstract #2703. 42nd Lunar and Planetary Science Conference. [2] Crowther S. A. and Gilmour J. D. 2011. Abstract #1969. 42nd Lunar and Planetary Science Conference. [3] Wieler R. and Baur H. 1994. *Meteoritics* 29:570–580. [4] Pepin R. O. et al. 1995. *Geochimica et Cosmochimica Acta* 59:4997–5022.

<table><thead><tr><th></th><th>α = π/8<br>n = 4</th><th>α = π/4<br>n = 6</th><th>α = π/2<br>n = 6</th><th>α = π/4<br>n = 4</th><th>71501<br>lunar soil [3,4]</th></tr></thead><tbody><tr><td><sup>132</sup>Xe<sub>SW</sub></td><td>1.23</td><td>1.01</td><td>1.16</td><td>1.18</td><td></td></tr><tr><td>10<sup>6</sup>at cm<sup>2</sup></td><td></td><td></td><td></td><td></td><td></td></tr><tr><td><sup>136</sup>Xe/<sup>132</sup>Xe</td><td>.2966 ± .0028</td><td>.2987 ± .0025</td><td>.3002 ± .0015</td><td>.3013 ± .0007</td><td>.2985 ± .0009</td></tr><tr><td><sup>134</sup>Xe/<sup>132</sup>Xe</td><td>.3713 ± .0029</td><td>.3699 ± .0028</td><td>.3749 ± .0018</td><td>.3724 ± .0008</td><td>.3666 ± .0015</td></tr><tr><td><sup>131</sup>Xe/<sup>132</sup>Xe</td><td>.8135 ± .0083</td><td>.8297 ± .0047</td><td>.8190 ± .0026</td><td>.8302 ± .0016</td><td>.8272 ± .0028</td></tr><tr><td><sup>130</sup>Xe/<sup>132</sup>Xe</td><td>.1613 ± .0020</td><td>.1641 ± .0017</td><td>.1614 ± .0010</td><td>.1612 ± .0005</td><td>.1661 ± .0009</td></tr><tr><td><sup>129</sup>Xe/<sup>132</sup>Xe</td><td>1.0413 ± .0051</td><td>1.0473 ± .0025</td><td>1.0408 ± .0014</td><td>1.0358 ± .0015</td><td>1.0420 ± .0070</td></tr><tr><td><sup>128</sup>Xe/<sup>132</sup>Xe</td><td>.0829 ± .0014</td><td>.0833 ± .0012</td><td>.0818 ± .0007</td><td>.0846 ± .0003</td><td>.0848 ± .0009</td></tr><tr><td><sup>126</sup>Xe/<sup>132</sup>Xe</td><td>.00405 ± .00029</td><td>.00389 ± .00027</td><td>.00371 ± .00017</td><td>.00438 ± .00013</td><td>.00423 ± .00014</td></tr><tr><td><sup>124</sup>Xe/<sup>132</sup>Xe</td><td>.00520 ± .00036</td><td>.00510 ± .00030</td><td>.00460 ± .00018</td><td>.00484 ± .00009</td><td>.00488 ± .00012</td></tr></tbody></table>

5178

ULTRARAPID CHONDRITE FORMATION BY HOT CHONDRULE ACCRETION? EVIDENCE FROM UOCs

K. Metzler. Institut für Planetologie, Wilhelm-Klemm-Strasse 10, 48149 Münster, Germany. E-mail: knut.metzler@uni-muenster.de.

**Introduction:** Most unequilibrated ordinary chondrites (UOCs) appear to be processed rocks, affected by postaccretional mechanical processes that led to abrasion, fragmentation, relocation and mixing of their constituents. This can be concluded from the brecciation textures of many UOCs [1], and the existence of components with aberrant chemical compositions [2] and shock stages [3]. Additionally, many obviously unbrecciated UOCs may represent “hidden” breccias [4], consisting of mixtures of components with different provenience and history. Hence, the search for mechanically unaltered relicts of the very first stages of chondrite accretion is a crucial task. In my view, those relicts are actually present in UOCs. They have been sporadically described in the literature as units with close-fit textures of mutually indented chondrules and are interpreted as the result of collisions between hot and thus plastic chondrules [e.g., 5–9].

**Results:** I found that these textures are restricted to a specific type of chondritic rock which occurs in several UOCs as lithic clasts up to 10 cm in size. This rock type consists of a mixture of deformed and undeformed chondrules and is characterized by low abundances of interchondrule matrix, low abundances of distinct chondrule fragments, and restricted variations of chondrules sizes. In order to distinguish this rock type from other chondritic lithologies I propose to name it “cluster chondrite.” In order to quantify the above textures I measured the degree of deformation for 188 chondrules in three cluster chondrite clasts from three brecciated UOCs (NWA 869, L3–6; NWA 1756, LL3.10; NWA 5205, LL3.2). This parameter is defined as the ratio of the measured chondrule perimeter and the calculated circle perimeter of the corresponding chondrule area in a thin section. Despite considerable data scatter a characteristic inverse correlation between the apparent size (area) of chondrules and their degree of deformation is observed. The meaning of this finding remains to be explained. For comparison, 53 chondrules from a “regular” UOC (NWA 4572, LL3.6) have been measured in the same way and in this case the described inverse correlation is missing.

**Conclusions:** Clasts of cluster chondrites seem to represent remnants of primary accretionary rocks of unknown original dimensions. They seemingly formed by concentration, collision and instantaneous compaction of hot and plastic chondrules within hours to a few days after chondrule forming events. Chondrule deformations seem to result from chondrite accretion, not from postaccretional processes, as advocated by [10]. The fact that clasts of these rocks occur in many UOCs may indicate that ultrarapid chondrite accretion immediately after chondrule formation has been a widespread process in the protoplanetary disk.

**References:** [1] Bischoff A. et al. 2006. In *Meteorites and the early solar system* II:679–712. [2] Scott E. R. D. 1984. Smithsonian Contributions to the Earth Sciences 26:73–93. [3] Stöffler D. 2011. Personal communication. [4] Romstedt J. and Metzler K. *Meteoritics* 29:523–524. [5] Hutchison R. and Bevan A. W. R. 1983. In *Chondrules and their origins* 162–179. [6] Taylor G. J. et al. 1983. In *Chondrules and their origins* 262–278. [7] Holmén B. A. and Wood J. A. 1986. *Meteoritics* 21:399. [8] Hutchison R. 1996. 27th Lunar and Planetary Science Conference. 579. [9] Zanda B. 2004. Earth and Planetary Science Letters 224:1–17. [10] Rubin A. E. and Brearley A. J. 1996. *Icarus* 124:86–96.
## PRESOLAR SILICON CARBIDE X GRAINS, EXPLOSIVE HYDROGEN BURNING, AND THE EVOLUTION OF MASSIVE STARS

B. S. Meyer¹, M. J. Bojazi¹, L.-S. The¹ and M. F. El Eid². ¹Department of Physics and Astronomy, Clemson University, Clemson, SC 29634-0978, USA. ²Department of Physics, American University of Beirut, Beirut, Lebanon.

**Introduction:** The nitrogen-15 abundance in presolar silicon carbide (SiC) X grains remains a puzzle. These grains presumably formed in the outflows from exploding massive stars (supernovae). Detailed supernova yields [1] fail to explain the correlation between aluminum and nitrogen isotopic signatures in the presolar SiC X grains [2].

**Nitrogen-15 Production:** To help understand the nitrogen isotopic abundances in presolar SiC X grains, we explore ¹⁵N nucleosynthesis in massive stars. Nitrogen-15 is produced as the supernova shock passes through the helium-rich layers [3] and heats them up. Nitrogen-14 captures an alpha particle to make ¹⁸F. The ¹⁸F can then undergo a (n,α) reaction to produce ¹⁵N. It can also be produced by explosive hydrogen burning when the shock first encounters the hydrogen-rich outer layers of the star. In these shock-heated layers, abundant ¹⁴N captures a proton to become ¹⁵O. After the shocked layers expand and cool, the ¹⁵O decays to ¹⁵N. The amount of ¹⁵N produced depends on the post-shock temperature.

**Presupernova Stellar Evolution:** Most supernova models do not show evidence of strong explosive hydrogen burning; however, the 25 solar mass model of [1] does show a strong spike in the ¹⁵N abundance due to this burning. The carbon and nitrogen signatures in presolar SiC X grains can be explained if material from the ¹⁵N spike is included in mixing models to account for the grains' origins [2]. The 25 solar mass model in [1] shows an ¹⁵N spike whereas, for example, the 15 solar mass models does not because the pre-shock temperature at the inner edge of the hydrogen shell in the 25 solar model is roughly 5 × 10⁷ K while it is about 2 × 10⁷ K for the 15 solar mass model. Since the post shock temperature is a multiple of the pre-shock temperature (about 2 for a Mach 2 shock and 4 for a Mach 3 shock), the 25 solar mass model can achieve a higher post-shock temperature than the 15 solar mass model for the same speed shock.

Our presupernova models of a 25 solar mass star show significant differences in the pre-shock temperature at the inner H-shell depending on the treatment of convection and mass loss [4]. If we include mass loss and use the Schwarzschild convection criterion, the temperature at the inner edge of the hydrogen shell is 2.5 × 10⁷ K. If we include semi-convection using the Ledoux criterion, the inner-H shell temperature is ~3.5 × 10⁷ K. If we use the Schwarzschild convection criterion and do not include mass loss, the temperature is ~2 × 10⁷ K. These differences mean that, if indeed ¹⁵N-spike material is needed to explain the nitrogen composition of SiC X grains, presolar SiC X grains may help constrain mixing and mass loss in presupernova stellar evolution.

**References:** [1] Rauscher T. et al. 2002. *The Astrophysical Journal* 576:323–348. [2] Lin Y. et al. 2010. *The Astrophysical Journal* 709:1157–1173. [3] Meyer B. S. and Bojazi M. J. 2011. Abstract #2376. 42nd Lunar and Planetary Science Conference. [4] El Eid M. F. et al. 2009. *Space Science Reviews* 147:1–29.

## PROCESS OF THE SUEVITE GENESIS OF THE RIES CRATER, GERMANY

C. Meyer and M. Jébrak. Université du Québec à Montréal, Canada.
E-mail: cornelia.meyer@utweiler.de.

**Introduction:** The processes of formation and transport of particles in suevite during impact crater formation remain poorly understood and were investigated at the 15-Ma-old, 25 km wide Ries crater, Germany. The suevite of the Ries crater occurs in three different geological settings: (1) Crater suevite in the central crater cavity inside the inner ring; (2) Outer suevite on top of the continuous ejecta blanket; and (3) Dikes in the crater basement and in displaced megablocks. For suevite genesis, the following processes have been discussed to-date in literature: (1) Fall-back of material into the crater and its periphery upon collapse of an ejecta plume [1], and (2) Horizontal transport of ejected material, akin to (a) an impact melt flow [2], (b) a pyroclastic flow [3], (c) initiated by phreatomagmatic explosion [4].

**Methods:** In order to differentiate between these emplacement modes we analyzed at different locations inside and outside the crater: (1) The three-dimensional shape and orientation of the suevite particles, (2) The stereometric parameters of melt particles and lithic clasts (grain shape, grain size distribution, and content of particles), and (3) The chemical composition of the melt particles compared to their corresponding matrix.

**Results:** The suevite can be distinguished in the following subunits: (1) Suevite with similar melt and matrix composition, which can be considered as original impact melt; (2) Strongly altered suevite, with similar melt and matrix composition; (3) Suevite containing K-poor and K-rich melt with either silicate matrix similar in chemistry to the melt, or non filled pore spaces; and (4) Suevite with different melt and matrix composition.

**Discussion:** Five stages can be distinguished for the formation and deposition of the Ries suevite: (1) An early ejecta plume, mainly composed of projectile and sediments, is formed during the excavation stage and is deposited (<2 m thick) inside and outside the crater on top of the Bunte Breccia after crater formation, (2) After a hiatus, the interaction of water with the impact melt lake induce a phreatomagmatic explosion of the upper melt lake, (3) A basal, non erosive pyroclastic surge, initiated by the explosion, moves radially outward and is deposited inside and outside the crater, (4) At the same time an ejection column rises. After its collapse, a density stratified pyroclastic flow is initiated which propagates to the crater margin and partially covers the first deposited suevite layer, (5) The collapse of the ejection column creates contemporaneously an upward rising ejecta plume, where accretionary lapilli are formed. Finally, these lapilli, together with the residual plume material, are deposited in the inner crater (a) into an existing water layer, or (b) reworked by incoming water in the crater at a later time.

**Acknowledgment:** This work was prepared at the Museum für Naturkunde, Berlin, and was supported by the Deutsche Forschungsgemeinschaft (RE 528/4-1).

**References:** [1] Stöffler D. 1977. *Geologica Bavaria* 75:443–458. [2] Osinski G. R. et al. 2004. *Meteoritics & Planetary Science* 39:1655–1683. [3] Newsom H. E. 1986. *Journal of Geophysical Research* 91:E239–E251. [4] Artemieva N. A. et al. 2009. Abstract #1526. 40th Lunar and Planetary Science Conference.
ELECTRON BACK-SCATTER DIFFRACTION STUDY OF IRON METAL IN ALMAHATA SITTA UREILITE

T. Mikouchi¹, C. A. Goodrich², V. H. Hoffmann³,⁴, M. E. Zolensky⁵ and K. Sugiyama⁶. ¹Department of Earth and Planetary Science, University of Tokyo, Tokyo 113-0033, Japan. ²Planetary Science Institute, Tucson, AZ 85719, USA. ³Department of Geoscience, University of Tuebingen, 72076 Tuebingen, Germany. ⁴Department of Geography-Environmental Science, University of Munich, 80333 Muenchen, Germany. ⁵NASA-JSC, Houston, TX 77058, USA. ⁶Institute of Materials Research, Tohoku University, Miyagi 980-0812, Japan. E-mail: mikouchi@eps.s.u-tokyo.ac.jp.

Almahata Sitta (AS) is a polymict breccia mainly composed of ureilite lithologies showing variable mineralogy [1]. These ureilite lithologies can be texturally divided into two types. One is a compact and coarse-grained, and composed of large olivine and pigeonite (reaching mm-size) with graphite and metal, and is similar to typical ureilites. The other is porous and consists of fine-grained (10–20 µm in size) olivine and pyroxene with small amounts of metal and graphite. Fe metal is commonly observed in both types and its unique mineralogy can offer important clues to understand the ureilite formation history [2, 3]. Here we report electron back-scatter diffraction (EBSD) analysis of Fe metal and associated opaque phases in three ureilite samples of AS.

The sample #1 (AS#1) is fine-grained similar to the sample #7 (AS#7) that was first described after the fall [1]. We observed both large vein and small scattered Fe metal grains in these two samples by EBSD. Almost all Fe metal grains are bcc kamacite (α-Fe) although they show variable grain sizes (5–100 µm). The sample #44 (AS#44) is a typical coarse-grained sample and contains vein metal (approximately 500 µm long) at grain boundaries between silicates [2]. These areas show a large diversity of phase assemblages and textures [2, 4]. One common assemblage is a crystalline intergrowth of two phases with apparent contrast in BEI in spite of no obvious compositional difference by EMPA [2]. The sample prepared by colloidal-silica polishing for EBSD work shows clear relief surface in these areas corresponding to the BEI contrast. In other intergrowths, elongated laths (up to 30 µm long) are present and are inferred from BEI and EMPA to be cohenite (Fe₃C). EBSD analysis showed that the obtained Kikuchi bands match with the cohenite structure. In these intergrowths, Fe metal is homogeneous in BEI, but EBSD shows that small amounts of γ-Fe (fcc) were present as well as kamacite.

The coexistence of two Fe metal phases (α and γ) and cohenite in one of the AS#44 metal intergrowths is important and similar occurrence can be found in commercial iron alloy [5]. γ-Fe (“austenite”) is stable at high temperature and upon cooling it decomposes into α-Fe (“ferrite”) and Fe₃C (“cementite”) at low temperature [5]. The small amount of γ-Fe in AS#44 may be relict because of fast cooling by shock melting [2]. Alternatively, shock remelting of pre-existing metal + graphite formed a eutectic-looking texture. Because the carbon content of these areas is far too high to have been dissolved in austenite at high temperature, we prefer the latter hypothesis. In contrast, Fe metal in AS#1 and #7 does not contain cohenite and γ-iron. This is probably due to their different thermal history and lower C content.

**References:** [1] Zolensky M. E. et al. 2010. *Meteoritics & Planetary Science* 45:1618–1637. [2] Goodrich C. A. et al. 2010. *Meteoritics & Planetary Science* 45:A66. [3] Hoffmann V. H. et al. 2011. *Meteoritics & Planetary Science* 46: in press. [4] Ross A. et al. 2011. Abstract #2720. 42nd Lunar and Planetary Science Conference. [5] Smith W. F. and Hashemi J. 2006. *Foundations of materials science and engineering*, McGraw-Hill. pp. 1056.

HYDROCODE SIMULATIONS OF BINARY ASTEROID IMPACTS

K. Miljkovic, S. Mannick, G. S. Collins and P. A. Bland. Earth Science & Engineering, Imperial College London, UK. E-mail: k.miljkovic@impreial.ac.uk.

**Introduction:** Observational data [1, 2] and numerical modeling [3, 4] show that 15% of the near-Earth asteroid (NEA) population are binary. However, the observed crater populations on Earth [5] and Venus [6] indicate that the number of doublet impact craters is much less than 15%. This discrepancy could originate from the assumption that may imply that not all binary asteroids create a clearly separated doublet crater, many could form a single or a complex shaped crater. In this work we use numerical modeling to quantify the threshold for the formation of distinct doublet craters on Earth and compare the complex crater morphology formed by a closely separated binary asteroid impact with that formed by oblique single asteroid impact.

**Background:** Identification of craters formed by binary asteroid impact on Earth is not simple because the craters are obscured by erosion and the presence of oceans. Conversely, on less active planetary bodies identifying doublet craters is hampered by a saturation of craters [6]. A rare example of a terrestrial doublet crater is Clearwater East and West that is estimated to have formed 290 ± 20 Myr ago [7]. The low probability of two craters forming so close together in space and time suggests they were formed at the same time by a binary asteroid. Clearwater East and West craters are 32 km and 22 km in diameter, respectively. Applying impact scaling laws [8] to the size of these craters we estimate that the impacting binary asteroid comprised a 3.3 km diameter primary and 2.1 km diameter secondary. The two craters are separated by 28.5 km, which corresponds to 11 times the primary asteroid diameter (Dp).

**Method:** The iSALE-3D hydrocode [9] was used to simulate binary asteroid impacts on Earth, using a granite material model as a representative for Earth’s surface and the asteroids. In the simulations, the asteroid diameters were: (i) Dp = 1200 m for the primary and Ds = 600 m for the secondary asteroid and (ii) Dp = Ds=600 m. The separation between the binary components (L) was varied from 1.25 Dp to 10 Dp. These values are representative of Ds/Dp and L for binary asteroids in NEA region, based on current observations [1].

**Results:** Preliminary results show that impact of a closely separated binary asteroid forms an elongated crater shape. For case (i) a single crater forms with a tear drop shape for $L < 8$ Dp; a doublet crater is formed if $L > 8$ Dp. For case (ii) a single, elliptical crater is formed for $L < 6$ Dp, if $6 < L < 8$ Dp the elliptical crater starts to show two off-centre peaks and only for $L > 10$ Dp does the impact form a doublet.

**References:** [1] Pravec P. et al. 2006. *Icarus* 181:63–93. [2] Margot J. L. et al. 2002. *Science* 296:1445–1448. [3] Richardson D. C. et al. 1998. *Icarus* 134:47–76. [4] Bottke W. F. and Melosh H. J. 1996. *Icarus* 124:327–391. [5] Melosh H. J. and Stansberry J. A. 1991. *Icarus* 94:171–179. [6] Cook C. M. et al. 2003 *Icarus* 165:90–100. [7] Earth Impact database (http://www.passc.net/EarthImpactDatabase/index.html). [8] Schmidt R. M. and Housen K. R. 1987. *International Journal of Impact Engineering* 5:543–560. [9] Elbeshausen D. et al. 2009 *Icarus* 204:716–731.
5484

Mg AND Cr ISOTOPIC COMPOSITIONS OF INDIVIDUAL
CHONDRULES FROM PRIMITIVE CHONDRITES

C. A. Miller, R. W. Carlson and C. M. O'D. Alexander. ¹Department of Terrestrial Magnetism, Carnegie Institution of Washington, 5241 Broad Branch Road NW, Washington, D.C. 20015, USA. E-mail: cmiller@dtm.ciw.edu.

Temporal constraints on the evolution of the early solar system are
provided by high-precision long-lived (absolute; Pb-Pb; [1]) and short-
lived (relative; e.g., Al-Mg, Mn-Cr; [2, 3]) radiochronometers. Short-lived
chronometers are particularly informative as their geochemical variety
permits evaluation of a variety of processes.

A series of high-precision Pb-Pb and Al-Mg ages indicate that, with the exception of the CBs, chondrules formed over an interval of ~1.8 Myr [2, 4]. A correlation between older chondrule ages and lower concentrations of volatile elements such as Si has also been suggested [2]. Such an extensive period of chondrule formation is difficult to reconcile with models for a turbulent protoplanetary disk that is thoroughly mixed on 10<sup>4</sup>-10<sup>6</sup> year time scales [5]. If the Al-Mg age spread is an accurate indicator of chondrule formation ages, then the conditions under which chondrules formed must be revisited; if not, then the variability may be due to mixing of older-refractory and younger-volatile reservoirs within chondrules, or of Al-Mg "resetting."

Evaluating individual chondrules for both Mg and Cr isotope
composition is a way of evaluating these hypotheses. Like Al-Mg, the
Mn-Cr chronometer is active during the interval of chondrule formation.
Relative age differences between chondrules that are consistent across
both chronometers would support a lengthy interval of chondrule
formation. Also, while the parent element is more refractory than the
daughter within the Al-Mg system, the situation is reversed for Mn-Cr [6].
Therefore, if chondrule age variability results from the mixing of
refractory and volatile reservoirs, individual chondrules with relatively old
Al-Mg ages should have relatively younger Mn-Cr ages.

Techniques for the purification of Mg and Cr have been developed
and isotope data from will be presented for individual chondrules from
primitive carbonaceous and ordinary chondrites.

**References:** [1] Amelin Y. et al. 2002. *Science* 297:1678-1683. [2] Mostefaoui S. et al. 2002. *Meteoritics & Planetary Science* 37:421-438. [3] Shukolyukov A. and Lugmair G. W. 2006. *Earth and Planetary Science Letters* 250:200-213. [4] Kurahashi E. et al. 2008. *Geochimica et Cosmochimica Acta* 72:3865-3882. [5] Alexander C. M. O'D. 2005. *Chondrites and the protoplanetary disk* pp. 972-1002. [6] Lodders K. 2003. *The Astrophysical Journal* 591:1220-1247.

5481

$^{26}$Al-$^{26}$Mg CHRONOMETRY OF CAIs FROM CV3 CHONDRITES

R. K. Mishra and M. Chaussidon. Centre de Recherches Pétrographiques et Géochimiques, CRPG - Nancy Université-CNRS, UPR 2300, 15 rue Notre Dame des Pauvres, BP 20, 54501 Vandoeuvre les Nancy, France. E-mail: ritesh@crpg.cnrs-nancy.fr.

**Introduction:** The advancements in analytical capabilities for meteoritic studies have enabled greater precision and detection of smaller isotopic variations for understanding phenomena, processes and their time scales hitherto unexplored. These improvements in analytical techniques open the possibility of utilizing the long-trusted work horse of $^{26}$Al-$^{26}$Mg systematics to test a few critical assumptions [1–3]. Calcium, aluminum-rich inclusions (CAIs) represent the first formed solids of Solar system and have Pb-Pb absolute ages of ~4568 Ma [4, 5]. High Precision Al-Mg systematics of CAIs can allow to constrain the physico-chemical conditions and their evolution during the earliest stage of birth of solar system [6, 7]. CAIs of different types in Vigarano and Efremovka (CV 3.1–3.4) were analyzed in the present study to: (1) Estimate the spatio-temporal scale of homogeneity (heterogeneity) of isotopic compositions, (2) Elucidate existence (or non existence) of a particular epoch leading to formation of a particular (or all) type of CAIs, (3) Existence of epochs or continuum of high temperature events leading to constraint on their causal's.

Results from $^{26}$Al-$^{26}$Mg analysis of three CAIs and one AOA from Efremovka and nine CAIs from Vigarano are presented here. The analyzed CAIs include typical type Bs, fine-grained spinel rich, and igneous ones. Their Mg isotopes show that they formed as early as 70 kyr to as late as 2.5 Myr after the birth of solar system (Fig. 1). Their isotopic studies along with bulk Al-Mg ratio can be used to decipher their complex evolution from condensation to melting and re-melting.

**References:** [1] Villeneuve J. et al. 2009. *Science* 325:985-988. [2] Jacobsen B. et al. 2008. *Earth and Planetary Science Letters* 272:353-364. [3] Davis et al. 2010. Abstract #2496. 41st Lunar and Planetary Science Conference. [4] Bouvier A. and Wadhwa M. 2010. *Nature geosciences* 3:637-641. [5] Amelin Y. et al. 2010. *Earth and Planetary Science Letters* 300:343-350. [6] MacPherson G. J. et al. 2010. *The Astrophysical Journal* 711:L117-12. [7] Kita et al. 2010. Abstract #2154. 41st Lunar and Planetary Science Conference.

<figure><img src="image_22.png" /><figcaption>Fig. 1. δ $^{26}$Mg$_{0}$ sVs $^{26}$Al/$^{27}$Al plot of early solar system solids.</figcaption></figure>
TOWARD A REGOLITH MATURITY INDEX FOR HOWARDITES

D. W. Mittlefehldt¹, J. A. Cartwright², J. S. Herrin³ and K. N. Johnson⁴. ¹NASA/Johnson Space Center, Houston, TX, USA. E-mail: david.w.mittlefehldt@nasa.gov. ²Max-Planck-Institut für Chemie, Mainz, Germany. ³ESGC/Johnson Space Center, Houston, TX, USA. ⁴Arizona State University, Tempe, AZ, USA.

The Dawn spacecraft has just arrived at asteroid 4 Vesta, parent of the howardite, eucrite and diogenite (HED) meteorites [1], to begin a yearlong surface study from orbit [2]. As Dawn will view a debris-covered surface, understanding the formation and mixing processes for the debris layer will strongly aid surface data interpretations. Howardites are polymict breccias mainly composed of clasts derived from basaltic (eucritic) and orthopyroxenitic (diogenitic) parent materials [3]. Some howardites are poorly reworked (fragmental howardites) whilst others have been extensively gardened in an active regolith (regolithic howardites) [4]. The latter may represent an ancient, well-mixed regolith, whilst the former may be from more recent ejecta deposits [4]. Due to environmental differences, regolith development on Vesta differs in detail from that on the Moon [4-6].

We have been developing petrological criteria to apply to howardite thin sections to determine their relative regolithic maturity, which we are fine-tuning with comparison to noble gas data [7, 8]. Whilst we previously emphasized the abundance of “reworked” clasts (fragmental and impact-melt breccia clasts), this is an imperfect criterion: one howardite with abundant reworked clasts (EET 99408) shows no evidence of solar wind Ne (SW-Ne), yet, two of our alleged fragmental howardites have clear SW-Ne signatures (LEW 85313, MET 00423) [7, 8].

We are now investigating the diversity in minor and trace element contents of low-Ca pyroxene clasts in howardites as a measure of regolith grade, and will begin analyses of such grains within reworked clasts. Our hypothesis is that regolithic howardites (and the breccia clasts they contain) will show greater diversity because they sampled more diverse diogenitic plutons than fragmental howardites, which formed from ejecta from only a few impacts [e.g., 4]. Our initial LA-ICP-MS work showed ranges in trace element diversity in low-Ca pyroxenes (estimated from the standard error of the mean of analyses), where those howardites considered of medium to high regolithic grade showed greater diversity [9]. Our EMPA results (from a larger howardite suite) show an overall greater diversity in our putative medium to high regolithic grade howardites, though there are exceptions. The greatest diversity is found for paired howardites GRO 95574 and GRO 95581, which were not considered regolithic in our initial study. We will continue investigating avenues to determine regolith maturity in thin section, factoring in bulk rock compositional data, and will coordinate these studies with noble gas results.

**References:** [1] Drake M. J. 2001. *Meteoritics & Planetary Science* 36:501–513. [2] Russell C. T. et al. 2006. *Advances in Space Research* 38:2043–2048. [3] Mittlefehldt D. W. et al. 1998. Reviews in Mineralogy vol. 36, chapter 4. [4] Warren P. H. et al. 2009. *Geochimica et Cosmochimica Acta* 73:5918–5943. [5] Housen K. R. and Wilkening L. L. 1982. *Annual Review of Earth and Planetary Sciences* 10:355–376. [6] Mittlefehldt D. W. et al. 2011. Abstract #2569. 42nd Lunar and Planetary Science Conference. [7] Cartwright J. A. et al. 2011. Abstract #2655. 42nd Lunar and Planetary Science Conference. [8] Cartwright J. A. et al. 2011. Abstract #5042. This Conference. [9] Johnson K. N. et al. 2011. Abstract #2073. 42nd Lunar and Planetary Science Conference.

IN SITU TRACE ELEMENT MEASUREMENTS ON RODA AND THE ORIGIN OF DIOGENITES

D. W. Mittlefehldt¹ and J. S. Herrin². ¹NASA/Johnson Space Center, Houston, TX, USA. E-mail: david.w.mittlefehldt@nasa.gov. ²ESGC/Johnson Space Center, Houston, TX, USA.

The origin of diogenites remains poorly understood. A recent model interprets many diogenites to have been formed from melts that were derived by remelting initial magma ocean cumulates, and these penultimate parent melts were then contaminated by melts derived from remelting of the basaltic (eucritic) crust to form the ultimate diogenite parent melts [1] (hereafter the remelting model). This is a very complicated petrogenesis that has profound implications for the geological evolution of 4 Vesta if correct. This model was developed based on trace element analyses of bulk rock samples that had been leached in acids to remove phosphates; the compositions of the residues were interpreted to be close to those of cumulus orthopyroxenes ± plagioclase, chromite and olivine [1]. In situ measurements of phases in diogenites can be used to test this model. We have begun a campaign of laser ablation ICP-MS of orthopyroxene grains in diogenites for this purpose. Here we report our first results on one diogenite, Roda.

We have determined a suite of trace lithophile elements on nine, mm-sized pyroxene grains separated from Roda that have previously been studied [2, 3]. A key observation supporting the remelting model is the very low Eu/Eu* of leached residues; values too low to represent orthopyroxene that crystallized from melts with chondritic Sm/Eu and Gd/Eu [1]. (Eu* = Eu interpolated from REE diagrams.) Crustal remelts have low Sm/Eu and Gd/Eu, and orthopyroxenes that crystallized from parent melts contaminated by them would have very low Eu/Eu* [1].

Roda grains have Eu/Eu* of 0.243 to 0.026; the latter a value lower than any measured on bulk diogenite leached residues (0.041) [1]. There is a general negative correlation between Eu/Eu* and some incompatible elements (Zr, Nb, Hf), but not others (LREE). This appears inconsistent with the remelting model as it would suggest an evolving parent melt with La decreasing as Zr increased and Eu/Eu* decreased. Grain R-15 includes trace element-rich trapped melt phases [2, 3]. This grain has the highest Eu/Eu* and LREE contents, indicating that the trapped melt had a high Eu/Eu*. Thus, our first data on one diogenite do not provide support for the remelting model [1].

Roda is unusual in that its orthopyroxene grains show wide ranges in trace element contents [4]. Previous in situ REE analyses of grain R-15 did not reveal evidence for subsolidus equilibration with trace element-rich trapped melt phases, and led to the suggestion that Roda may be polymict, with different grains representing different lithologies of diverse compositions [3]. Thus, based on our results on Roda, it is perhaps premature to abandon the remelting model. In situ measurements on a suite of diogenites is planned to further address this issue.

**References:** [1] Barrat J.-A. et al. 2010. *Geochimica et Cosmochimica Acta* 74:6218–6231. [2] Mittlefehldt D. W. 1994. *Geochimica et Cosmochimica Acta* 58:1537–1552. [3] Fowler G. W. et al. 1997. Abstract #1494. 28th Lunar and Planetary Science Conference. [4] Fowler G. W. et al. 1995. *Geochimica et Cosmochimica Acta* 59:3071–3084.
CARBON-BEARING IMPACT-RELATED MATERIALS OF BURIED
CRATER BROKEN AT TAKAMATSU, JAPAN

Y. Miura. Yamaguchi University and Al. I. Cuza University. E-mail: moonyas@hotmail.com.

**Introduction:** Impact-related evidences at buried and broken impact crater structure in Takamatsu, Shikoku, Japan are reported mainly by gravity anomaly and shocked-related materials on the Cretaceous Ryoke granite [1–10]. Together with impact-related materials on the granitic basement rock, anomalous materials with carbon-bearing materials which are not usually included in the granitic rocks are found at drilled materials. Main purposes of the paper are to elucidate carbon and calcium-bearing nano-particles at the Takamatsu, Shikoku, Japan [9, 10].

**Anomalous Carbon and Calcium Contents:** The Ryoke granitic rock at Takamatsu has few carbon and calcium contents in bulk XRF analyses. However, drilled sample at 950 m in depth contains nano-grains with high carbon (0.001–0.01 mm) and CaSiCO₃ in composition (in 100 nm size). Carbon-rich grains in Takamatsu show very weak nano-diamond-like peak in the Raman spectra.

**Formation Process of Carbon-Bearing Grains:** Main source of carbon and calcium cannot be found in the granitic rock as follows:

1) Carbon sources mainly found at crater bottom (ca. 950 m in depth) are considered to be foreign rocks of sedimentary limestone (from sea-sediments) or carbon dioxides gas during impact reaction probably on sea-water impact.

2) Calcium sources to form nano-grains of CaSiCO₃ in composition are also considered to be sedimentary limestone stored in the underground due to similar origins of carbon sources.

**Summary:** The present study is summarized as follows:

1) Very high carbon contents at Takamatsu in Japan are obtained in drilled sample of 950 m in depth, which indicates the sources from sedimentary sea-limestone and/or carbon dioxides gas during impact reaction on sea-impact.

2) Calcium sources with similar origin of carbon to form nano-grains of CaSiCO₃ in composition are considered to be sedimentary limestone existed in the underground.

**References:** [1] Kono Y. et al. 1994. *Report of Kaken B (Japan)* 36 pp. [2] Miura et al. 1995. *Meteoritics* (USA) 30:552. [3] Miura et al. 1999. Abstract #1127. 30th Lunar and Planetary Science Conference. [4] Miura et al. 2000. Abstract #3152. 31st Lunar and Planetary Science Conference. [5] Miura Y. et al. 2000. Abstract #4132. 31st Lunar and Planetary Science Conference. [6] Miura Y. et al. 2001. Abstract #1981. 32nd Lunar and Planetary Science Conference. [7] Miura Y. 2002. Abstract #1231. 33rd Lunar and Planetary Science Conference. [8] Miura Y. et al. 2004. Abstract #2094. 35th Lunar and Planetary Science Conference. [9] Miura Y. 2005. Abstract #2394. 34th Lunar and Planetary Science Conference. [10] Miura Y. et al. 2006. Abstract #1239. 37th Lunar and Planetary Science Conference. [11] Miura Y. 2009. Abstract #2565. 40th Lunar and Planetary Science Conference.

FIRST EVIDENCE FOR NATURAL DISSOCIATION OF OLIVINE
TO SILICATE-PEROVSKITE AND MAGNESIOWÜSTITE IN A
SHOCKED MARTIAN METEORITE DAG 735

M. Miyahara¹, E. Ohtani¹, S. Ozawa¹, M. Kimura², A. El Goresy³, T. Sakai¹, T. Nagase¹, K. Hiraga⁴, N. Hirao⁵ and Y. Ohishi⁵. ¹Tohoku University. E-mail: miyahara@m.tohoku.ac.jp. ²Ibaraki University. ³Universität Bayreuth. ⁴IMR. ⁵Spring-8.

**Introduction:** High-pressure and -temperature experiments expect that olivine dissociates to silicate-perovskite (Pv) + magnesiowüstite (Mw) at ~24 GPa (e.g., [1]). However, the dissociation of olivine at high-pressure and -temperature condition has never been reported from any natural samples so far. Therefore, the dissociation mechanism of olivine in natural samples is still unclear. Olivine is one of major constituents of the Martian meteorite. Some Martian meteorites record high-pressure conditions of ~24 GPa or even inferred higher pressures during natural dynamic events (e.g., [2]). It is likely that olivine in such Martian meteorites dissociates into Pv + Mw because the pressure conditions are beyond the stable pressure field of olivine, wadsleyite and ringwoodite. Accordingly, we carefully scanned olivine grains in the Martian meteorite, Dar al Gani (DaG) 735 with FE-SEM and FIB-TEM techniques, and here, we report the first evidence for olivine dissociation at high-pressure and -temperature conditions.

**Results and Discussion:** DaG 735 is an olivine-phyric shergottite and contains a shock-melt vein. The dissociation texture of olivine (Fa₃₄₄₁) was observed in olivine grains adjacent to the shock-melt vein. FE-SEM images show that the dissociation textures can be divided into two settings; 1) fine-grained granular (< ~100 nm) part and 2) coarse-grained (< ~1.0 µm) granular part. Fine-grained granular textures exist always adjacent to the shock-melt vein. The abundance of coarse-grained granular texture increases in dominance with increasing a distance from the shock-melt vein wall to the interior of the olivine grains. Raman spectra corresponding to pyroxene-glass were obtained from both dissociation settings. FIB-TEM studies show that the fine-grained granular parts consist of equigranular Mw and vitrified-Pv. The coarse-grained granular parts include fragments with alternating lamellar intergrowth. The alternating lamellae consist of Mw and vitrified-Pv.

The dissociation mechanisms of olivine or ringwoodite to Mw + Pv were previously studied using high-pressure generating devices (e.g., [3]). Micro-textural evolution from lamellar into equigranular, similar to the textural sequence encountered here, was previously observed with increasing temperature but at the same pressure conditions. The micro-textural evolution observed in DaG 735 probably resulted from the decreasing thermal gradient but at the same pressure magnitude into the olivine grain during a dynamic event. Olivine dissociation requires the pressure condition of ~24 GPa or more. In addition, a kinetic barrier exists around 700°C for the dissociation of olivine [4]; i.e., the dissociation does not take place below ~700°C even at higher overstepped pressure. Accordingly, pressure condition recorded in the shock-melt vein of DaG 735 would be beyond ~24 GPa, and temperature condition should be > 700°C at least.

**References:** [1] Ito E. and Takahashi E. 1989. *Journal of Geophysical Research* 94:10637–10646. [2] Stöffler D. et al. 1986. *Geochimica et Cosmochimica Acta* 50:889–903. [3] Poirier J. P. et al. 1986. *Nature* 321:603–605. [4] Wang Y. et al. 1997. *Science* 275:510–513.
## DAG 1054: A NEW ANOMALOUS UREILITE FROM DAR AL GANI

V. Moggi-Cecchi¹, S. Caporali¹,³ and G. Pratesi². ¹Museo di Scienze Planetarie, Provincia di Prato, Via Galcianese, 20/h, I-59100 Prato, Italy E-mail: v.moggi@pratoricerche.it. ²Dipartimento di Scienze della Terra, Università di Firenze, Via La Pira, 4, I-50121, Florence, Italy. ³Dipartimento Chimica, Università di Firenze, Via della Lastruccia, 3, I-50019, Sesto Fiorentino, Italy.

**Introduction:** DaG 1054 is a new meteorite found in 1999 by Romano Serra during a scientific expedition in the Dar al Gani rocky desert area of Libya. The meteorite consists of a single stone weighing 28.4 g, with no traces of fusion crust. A cut surface on one side reveals an achondritic texture. The type specimen (8 g), a thick section and a polished thin sections are on deposit at Museo di Scienze Planetarie, Prato, Italy (inv. N° 5138). The Museo del Cielo e della Terra (OAM), Bologna, Italy holds the main mass [1].

**Description:** The sample has a very fine-grained, homogeneous unrebrecciated and mosaicized texture mainly consisting of very small (up to 30 µm) rounded and non-zoned olivine crystals and rare orthopyroxene (Fig. 1a). Carbon-rich material, including very small diamond grains, occurs at grain boundaries but no distinct aggregates of carbonaceous material are visible. Opaque phases are formed almost exclusively by Fe-silicides of non-stoichometric suessitic composition, dispersed in small to very small (up to 100 µm across) grains or sub parallel veins, frequently altered to oxides (Fig. 1b). Very rare sulfide grains (mainly troilite) can be occasionally found. EMPA analyses performed on olivine and pyroxene provided the following results: Low-Ca pyroxene: $Fs_{12,2-20,9}Wo_{0.7-3,0}$, $Cr_2O_3$ = 0.85-0.92 wt%; Olivine ($Fa_{11,0-17,1}$, CaO = 0.29-0.31 wt%, $Cr_2O_3$ = 0.80-0.85 wt%). SEM-EDS analyses performed on the suessitic phase proved it to have a variable, non-stoichometric composition, with 2 main clusters at $Fe_{78}Si_{21}$ at % (Cr = 0.68, Ni = 0.49) and at $Fe_{74}Si_{23}$ at. % (Cr = 1.34, Ni = 1.54). Troilite contains up to 4.3 wt% Cr. The mosaicized structure suggests a high shock stage; weathering grade is low.

**Conclusions:** Although several features like the fine grained and mosaicized texture are similar to those displayed by polymict ureilites like Haverö, Y-74154, ALH 81101, and Almahata Sitta, the thin section contains no coarse-grained clasts and appears homogeneous and unrebrecciated. The lack of zoning of olivine and the presence of suessite of variable composition, also reported by [2–5] for other ureilites, suggest a closer relationship with both NWA 1241 and the fine-grained clast MS-161 of Almahata Sitta [6]. These anomalous features indicate that this meteorite could be a fine-grained monomict clast of a larger polymict body.

**References:** [1] Weisberg M. K. et al. 2011. *Meteriotics & Planetary Science* 46:3, in press. [2] Ikeda Y. 2007. *Polar Science* 1:45. [3] Herrin J. S. et al. 2007. Abstract #2404. 38th Lunar and Planetary Science Conference. [4] Ross A. J. et al. 2010. Abstract #2361. 41st Lunar and Planetary Science Conference. [5] Smith C. L. et al. 2010. *Meteriotics & Planetary Science* 45:Suppl., A192. [6] Bishoff A. et al. 2010. *Meteriotics & Planetary Science* 45:10–11, 1638.

<figure><img src="image_23.png" /></figure>

## COMETARY NOBLE GASES TRAPPED BY THE STARDUST AEROGEL

R. K. Mohapatra¹, U. Ott² and I. D. Clark¹. ¹University of Ottawa, Ottawa K1N 6N5, Canada. ²Max-Planck-Institut für Chemie, D-55020 Mainz, Germany. E-mail: uli.ott@mpic.de.

**Introduction:** The Stardust mission used monolithic silica aerogel for collecting particles from comet Wild 2 [1] because of its remarkable porosity. During its flight through the comet's coma, Stardust encountered also cometary gas, however. We report here preliminary observations from laboratory experiments and modeling of cometary gas flux at the Stardust aerogel.

**Adsorption and Desorption:** Results from our experiments on a sample of “pre-flight” aerogel indicate efficient gas adsorption and desorption properties. An aerogel cold trap, presumably cleaner in terms of hydrocarbon background, may be ideal for processing gas samples in simultaneous nitrogen and noble gas mass spectrometry.

**Implantation of Cometary Gas:** We have also explored the possibility of sampling cometary gas introduced into the Stardust aerogel by implantation. The encounter velocity of 6.1 km sec⁻¹ [2] corresponds to implantation energy of roughly 0.2 eV per nucleon, which does not support trapping of the light noble gases. We have used the TRIM code [3] to investigate trapping efficiency by implantation for the heavy noble gases- Ar, Kr, and Xe. Incident cometary argon at the Stardust aerogel was modeled by assuming Ar/O as measured in Hale-Bopp [4] and oxygen inferred from water production by Wild 2 [2, 5]. Abundances of Kr and Xe were derived by assuming solar like Ar/Kr/Xe. Results are compiled in Fig. 1. For an aerogel density of 0.066 g cm⁻³, determined on our pre-flight aerogel, implanted Ar, Kr, and Xe may reside in the upper 4.5–32 nm region of the flown aerogel; it would be sited accordingly deeper for a density <0.01 g cm⁻³ characteristic for the impact surface [1].

**Conclusions:** Preliminary results from laboratory experiments and modeling of cometary gas suggest near-surface samples of the Stardust aerogel should be searched for possibly implanted cometary gases.

**Acknowledgments:** We thank M. Zolensky for discussions and the Stardust Allocation committee for providing aerogel samples for our tests.

**References:** [1] Brownlee et al. 2006. *Science* 314:1711–1716. [2] http://stardust.jpl.nasa.gov/science/wild2.html. [3] Ziegler J. F. et al. 2010. *Nuclear Instruments and Methods in Physics Research B* 268:1818–1823. [4] Stern S. A. et al. 2000. *The Astrophysical Journal Letters* 544:L169–L172. [5] de Val-Borro M. et al. 2010. *Astronomy & Astrophysics* 521:L50.

<figure><img src="image_24.png" /><figcaption>Fig.1. Model assessment of expected cometary gas in aerogel.</figcaption></figure>
## REVISITING FIRES AT THE K-Pg BOUNDARY

J. Morgan¹, N. Artemieva²,³, C. Belcher⁴, T. Goldin⁵ and Elisabetta Pierazzo². ¹Imperial College London, London, UK. E-mail: j.morgan@imperial.ac.uk. ²Planetary Science Institute, Tucson, AZ, USA. ³Institute for Dynamics of Geospheres, Moscow, Russia. ⁴University of Vienna, Vienna, Austria. ⁵University of Edinburgh, Edinburgh, UK.

**Introduction:** The discovery of large amounts of soot in K-Pg boundary clays led to the hypothesis that there were major wildfires at the K-Pg boundary [1, 2]. Impact generated fires could be ignited in two ways: through the fireball for distances of up to 3000 km [3] and globally through the re-entry of hypervelocity ejecta [4]. Models of the latter have assumed a particular velocity and mass flow for the re-entering ejecta [4, 5]. In our new simulations we model the entire process, starting with the ejection of material from Chicxulub and ending with these ejecta re-entering the Earth's atmosphere. We have also reviewed the evidence that has been used to argue against the global wildfire hypothesis [6, 7], including the lack of charcoal and relative abundance of uncharred material in K-Pg boundary sites in North America.

**Method:** We simulate the Chicxulub impact with the 3D hydrocode SOVA [8] complemented by the ANEOS equation of state for geological materials [9]. With SOVA, we model (1) the impact and initial ejection of material and (2) the ballistic continuation of ejecta on a spherical earth until the ejecta arrives at the top of the atmosphere [10]. We then use the 2D two-phase fluid flow code KFIX-LPL to model the atmospheric re-entry of ejecta [5] and the thermal pulse delivered to the Earth's surface. We have modeled the ejection of material for two impact angles, 45° and 60°, and calculated the associated thermal pulse at 2000–2500 km and 7000–8000 km from Chicxulub. These distances correspond to the southern USA and Europe—locations where there is a wealth of observational data.

**Results:** When sedimentation rates and dilution of charcoal with other ejecta material are taken into account, the K-Pg boundary record in North America may be consistent with local wildfires. Some fires at, or shortly after, impact appear to be required to explain the large mass of soot in the global K-Pg layer. The thermal pulse delivered to the surface depends on radiating temperature of the ejecta and the opacity distribution of the atmosphere and thus is strongly dependent on re-entry velocity, spherule size and mass flow of ejecta, and these change with angle of impact and direction (downrange, uprange, etc.). The thermal pulse (≥ 20 kW m⁻²) in the downrange direction at sites up to a few thousand kilometers from Chicxulub may be sufficient to ignite fires, but is less severe (< 20 kW m⁻²) for distal sites and in other directions. Our results suggest that ejecta re-entry ignited some wildfires, but did not instantaneously ignite fires globally.

**References:** [1] Wolbach W. et al. 1985. *Science* 230:167–170. [2] Wolbach W. et al. 1990. Geological Society of America Special Paper 247:391–400. [3] Shuvalov V. and Artemieva N. 2002. Geological Society of America Special Paper 356. pp. 695–703. [4] Melosh H. et al. 1990. *Nature* 343:251–254. [5] Goldin T. and Melosh H. 2009. *Geology* 37:1135–1138. [6] Belcher C. et al. 2005. *Geological Society of London* 162:591–602. [7] Belcher C. et al. 2009. *Proceedings of the National Academy of Sciences* 106:4112–4117. [8] Shuvalov V. 1999. *Shock Waves* 9:381–390. [9] Thompson S. L. and Lauson H. S. 1972. *SC-RR-61 0714*. Sandia Nat. Lab. 119 p. [10] Artemieva N. and Morgan J. 2009. *Icarus* 201:768–780.

## SIMULATION OF GAS-GRAIN COLLISIONS: A MECHANISM FOR CHONDRULE FORMATION

A. Morlok¹, Y. C. Sutton², M. Patel¹, N. St. J. Braithwaite² and M. M. Grady¹,³. ¹PSSRI, The Open University, Walton Hall, Milton Keynes, MK7 6AA, UK. E-mail: a.morlok@open.ac.uk. ²Department of Physics & Astronomy, The Open University, Walton Hall, Milton Keynes, MK7 6AA, UK. ³The Natural History Museum, Cromwell Road, SW7 5BD, London, UK.

**Introduction:** The exact process of chondrule formation is still not known. The most common scenarios are formation in a protoplanetary nebular environment: e.g., X-wind model [1] or nebular shock fronts, where precursor grains are heated by collisions with gas [2]. Most chondrules probably have been heated to 1500–1600 °C, followed by a cooling rate of usually between 10 and 1000 °C h⁻¹ [3]. In this part of our study about the formation of chondrules by shock processes, we try to simulate the formation of chondrules by gas-grain collisions in a nebular environment.

Earlier, related studies usually did not result in chondrule-like objects [4]. In a first series of experiments using plasma arcs, we produced successfully spherical silicate droplets in a simple set-up under atmospheric conditions. These chondrule-like objects already show similarities to actual chondrules [5]. In a second series, the experiments take place in an environment closer to the earlier solar nebula, at low pressure and a reducing environment.

**Techniques:** A plasma arc is generated at atmospheric pressure using a low voltage solid state switching circuit to drive a Tesla coil (natural frequency ~325 kHz). The gas temperature was determined in an earlier experiment through fitting of modeled spectra to observed emission spectra: At 25W, the gas temperature is in the region of >2100°K.

The plasma generator is placed in an environmental chamber at the PSSRI/Open University. With the help of a remote controlled dispenser, starting material (mixture of olivine, feldspar, metal, and carbon) is dropped through the arc) under varying atmospheres (air and N2) at changing pressures of 5 × 10⁻¹, 1 × 10⁻¹, 4.8 × 10⁻², 5 × 10⁻³ bar.

The run products are recovered below the arc and analyzed using SEM/EDX for their structure and chemical composition.

**First Results:** The first run did show interaction of the material with the plasma arc. However, preliminary screening of the run products with SEM/EDX indicates only few melt droplets, which are difficult to identify in the original material [5]. Here we will present the results of a new series of experiments with varying starting compositions.

**References:** [1] Shu et al. 2001 *The Astrophysical Journal* 548:1029–1050. [2] Desch and Connolly Jr. 2002. *Meteoritics & Planetary Science* 37:183–208. [3] Zanda B. 2004. *Earth and Planetary Science Letters* 224:1–24. [4] Güttler et al. 2008. *Icarus* 195:5094–5510. [5] Morlok et al. 2011. *Abstract #1081*. 42nd Lunar and Planetary Science conference.
THERMAL HISTORIES OF CHONDRULES: AN ASSESSMENT OF THE EFFECT OF A SIZE DISTRIBUTION OF PRECURSOR PARTICLES

M. A. Morris¹ and S. J. Desch¹. ¹School of Earth and Space Exploration, Arizona State University, AZ, USA. E-mail: melissa.a.morris@asu.edu.

**Introduction:** Chondrules represent one of the best probes of the physical conditions and processes acting in the early solar nebula. Chondrule formation models are assessed on their ability to match the meteoritic evidence, especially constraints on their thermal histories. The model most consistent with chondrule thermal histories appears to be passage through shock waves in the solar nebula [1, 2]. These models meet all known constraints on thermal histories during chondrule formation, except for the restriction on the amount of time chondrules remain at >1300 K prior to achieving peak temperature, which [3] estimate is <30 min, based on the lack of isotopic fractionation of sulfur. In the shock model of [4], chondrules are “preheated” to temperatures above 1300 K in the preshock region, remaining there for over 4 h. Heating in the preshock region is controlled mainly by opacity, which is greatly affected by the sizes of solids, with smaller particles contributing more. Previous models have typically included chondrule precursors of only one size [4–7]. We feel it is now necessary to include a simple size distribution of precursor particles in the shock model in order to assess the effects on the thermal histories of chondrules, especially in the preshock region.

**Methods:** We used the shock code of [4] to carry out our calculations of chondrule thermal histories with the addition of “microchondrules” of radius, *a* = 10 µm, mixed with 300 µm chondrule precursors and submicron-sized dust. This improves on the models of [4, 5], which assumed a simple monodispersion of 300 µm precursors, as well as micron-sized dust. The dust, considered coupled to the gas, evaporates completely and nearly instantaneously at 1500 K [4]. Chondrules and microchondrules evaporate at identical, energy-limited rates when their temperatures exceed their liquidus temperature of approximately 1820 K. Vapor evaporated from the chondrule precursors is added back to the gas.

**Results and Discussion:** We considered two cases: one in which all large solids are in chondrules, and one in which half of the mass of chondrules is replaced by microchondrules. A key result from these studies is that microchondrules are present in the preshock region, where they significantly enhance the opacity (by a factor of 15), but they evaporate in the postshock region, and therefore do not contribute to the opacity or cooling rate there. Cooling rates through the crystallization range are raised only slightly over the canonical case of [4], from approximately 10–20 to 20–40 K h⁻¹, still well within the range consistent with the meteoritic constraints on thermal histories. The time spent at T > 1300 K before the shock, however, is reduced to <30 min.

**Conclusion:** The inclusion of a simple size distribution of chondrules, containing “microchondrules” of ≤10 µm in radius, leads the shock model to match the last unmet constraint on chondrule formation.

**References:** [1] Desch S. J., Morris M. A., Connolly H. C., Jr., and Boss A. P. 2011. *Meteoritics & Planetary Science*, submitted. [2] Desch S. J., Morris M. A., Connolly H. C., Jr., and Boss A. P. 2010. *The Astrophysical Journal* 725:692. [3] Tachibana S. and Huss G. R. 2005. *Geochimica et Comochimica Acta* 69:3075. [4] Morris M. A. and Desch S. J. 2010. *The Astrophysical Journal* 722:1474. [5] Desch and Connolly, 2002. *Meteoritics & Planetary Science* 37:183. [6] Ciesla F. J. and Hood L. L. 2002. *Icarus* 158:281. [7] Iida A., Nakamoto T., Susa H., and Nakagawa Y. 2001. *Icarus* 153:430.

LABORATORY EXPERIMENTS OF ROTATIONALLY DEPENDANT CATASTROPHIC DISRUPTION

A. J. W. Morris, M. J. Cole, M. C. Price and M. J. Burchell. School of Physical Science, University of Kent, Canterbury, CT2 7NH, UK. E-mail: am661@kent.ac.uk.

**Introduction:** Catastrophic disruption of minor bodies is a major evolutionary process for rocky bodies in the asteroid belt, e.g., see [1] for a general review and [2] for a typical recent paper (concerning the Themis family). Two tools are available for its study, hydrocode simulations and laboratory experimentation. Both offer advantages. In the latter case it is possible to control key parameters and see how they influence the outcome of impact events (albeit at laboratory, i.e., cm scale). Previous studies have paid attention to areas concerning gravity or impact related disruption [2] and the use of hydrocode modeling to ascertain internal faults and their effects on disruption [3] thresholds to mention a few variables usually investigated. But although spin is a feature of asteroids, it is rarely considered in impact experiments. Here we use laboratory impact experiments to study the influence of target spin on the outcome of the impact event. Accordingly, we examine the effects governing hypervelocity impacts on solid cement spheres.

**Method:** The target spheres were made in-house and had (on average) a diameter and mass of 75 mm and 367 g respectively. A rotating target holder was used which had an average rotation rate of 3.44 revs⁻¹. Two datasets were obtained for stationary (0 rev⁻¹) and rotating shots (3.44 revs⁻¹) with 22 impacts in total. Impacts were obtained using a two stage light gas gun [4] firing stainless steel spheres (dia. 1–3 mm) at speeds between 1 and 7.5 km s⁻¹. This gave impact energy densities *Q* (kinetic energy/original target mass) in the range 7–2600 J kg⁻¹. This spanned the regime of cratering to disruption defined by m<sub>f</sub>/m<sub>i</sub> (mass largest fragment/original target mass) being greater than or less than 0.5 respectively, with *Q*\* being the energy density at m<sub>f</sub>/m<sub>i</sub> = 0.5. The m<sub>f</sub>/m<sub>i</sub> results ranged from 0.99 to 0.04 in each of the data sets. For stationary targets it was found that *Q*\* = 1500 J kg⁻¹, and we noted that m<sub>f</sub>/m<sub>i</sub> fell below 0.8 at *Q* = 750 J kg⁻¹. For the rotating targets we found that whilst *Q*\* = 1450 J kg⁻¹, m<sub>f</sub>/m<sub>i</sub> fell below 0.8 at *Q* = 450 J kg⁻¹.

**Discussion:** From the present data we found that the stationary and rotating targets had similar *Q*\* values and thus the level of rotation in these experiments is insufficient to affect the overall disruption threshold. There may however be a difference in behavior at low *Q* values with rotating targets removing more spall material from craters—this is being investigated further.

**References:** [1] Ryan E. V. 2000. *Annual Review of Earth and Planetary Science* 28:367–389. [2] Leliwa-Kopystynski J. et al. 2009. *Meteoritics & Planetary Science* 44:1929–1935. [3] Michel P. et al. 2004. *Icarus* 168:420–432. [4] Stewart S. T. et al. 2009. *The Astrophysical Journal* 691:L133–L137. [5] Burchell M. J. et al. 1999. *Measurement Science & Technology* 10:41–50.
## ESA LUNAR LANDER'S SEARCH FOR VOLATILES

A. D. Morse¹, S. J. Barber¹, J. M. Pillinger¹, S. Sheridan¹, I. P. Wright¹, E. K. Gibson², J. A. Merrifield³, N. R. Waltham⁴, L. J. Waugh⁵ and C. T. Pillinger¹.
¹Planetary & Space Sciences Research Institute, Walton Hall, The Open University, Milton Keynes. MK7 6AA, UK. E-mail: a.d.morse@open.ac.uk. ²Astromaterials Research Office, NASA Johnson Space Center, Houston, TX 77058, USA. ³Fluid Gravity Engineering, 1 West Street, Emsworth, PO10 7DX, UK. ⁴RAL Space, Didcot, Oxfordshire, OX11 0QX, UK. ⁵EADS Astrium Ltd., Gunnels Wood Rd., Stevenage SG1 2AS, UK.

**Introduction:** Following the Apollo era the moon was considered a volatile poor body. Samples collected from the Apollo missions contained only ppm levels of water formed by the interaction of the solar wind with the lunar regolith [1]. However more recent orbiter observations have indicated that water may exist as water ice in cold polar regions buried within craters at concentrations of a few wt% [2]. Infrared images from M3 on Chandrayaan-1 have been interpreted as showing the presence of hydrated surface minerals with the ongoing hydroxyl/water process feeding cold polar traps. This has been supported by observation of ephemeral features termed "space dew" [3]. Meanwhile laboratory studies indicate that water could be present in appreciable quantities in lunar rocks [4] and could also have a cometary source [5].

**Lunar Lander:** The European Space Agency (ESA's Directorate of Human Spaceflight and Operations) have provisionally scheduled a robotic mission to demonstrate key technologies to enable later human exploration. Planned for launch in 2018, the primary aim is for precise automated landing, with hazard avoidance, in zones which are almost constantly illuminated (e.g., at the edge of the Shackleton crater at the lunar south pole). The Lander design requirement for constant solar illumination restricts the landing location, however these sunlit areas are close to permanently shadowed volatile rich regions and the analysis of volatiles is a major science objective of the mission.

**Lunar Volatile and Resources Analysis Package:** The authors have been commissioned by ESA to conduct an evaluation of possible technologies to be included in L-VRAP which aims to demonstrate volatile extraction and determine the volatile inventory of the moon with a view for future In Situ Resource Utilization (ISRU). Surface samples will be collected by a robotic arm with the possibility of a rover to collect more distant samples. The concentration, chemical and accurate isotopic ratios (D/H, ¹²C/¹³C, ¹⁵N/¹⁴N, ¹⁸O/¹⁶O and noble gases) of liberated volatiles will be determined, possibly using similar technology to the Philae comet lander of the Rosetta mission [6]. In addition the chemical and isotopic composition of the transient lunar atmosphere will be monitored [7]. Modeling the effects of contamination from the Lander is an essential part of this study so that these can be recognized and minimized.

**References:** [1] Pillinger C. T. 1979. *Report Progress Physics* 42:897-961. [2] Mitrofanov I. G. et al. 2010. *Science* 330:483-486. [3] McCord T. B. et al. 2010. *Journal Geophysical Research*, submitted. [4] Boyce J. W. et al. 2010. *Nature* 466:466-469. [5] Greenwood J. P. et al. 2011. *Nature Geoscience* 4:79-82. [6] Morse A. D. et al. 2009. Rosetta: ESA's mission to the origin of the solar system, 19.6:669-686. [7] Hofmann J. H. 1975. Lunar Atmospheric Composition Experiment, NASA report.

## SIZE OF THE GROUP IVA IRON METEORITE CORE: CONSTRAINTS FROM THE AGE AND COMPOSITION OF MUONIONALUSTA

N. A. Moskovitz¹ and R. Walker².
¹Carnegie Institution of Washington, DTM, DC, USA. E-mail: nmoskovitz@dtm.ciw.edu. ²University of Maryland, MD, USA.

**Introduction:** The group IVA fractionally crystallized iron meteorites display a diverse range of metallographic cooling rates, ranging from 100 to 6600 K Myr⁻¹ [1, 2]. These have been attributed to their formation in a metallic core, ~150 km in radius that cooled to crystallization without any appreciable insulating mantle. Such an exposed core may have resulted from a hit-and-run collision [3] between two large (~10³ km) proto-planetary bodies. Here we build upon this formation scenario by incorporating several new constraints. These include (1) a recent U-Pb radiometric closure age of 4565.3 Mya (<2.5 Myr after CAIs) for the group IVA iron Muonionalusta [4], (2) new measurements and modeling of highly siderophile element compositions for a suite of IVAs, and (3) consideration of the thermal effects of heating by the decay of the short-lived radionuclide ⁶⁰Fe.

**IVA Fractional Crystallization Model:** We model the fractional crystallization sequence of the IVA system using an approach similar to [5]. With initial S, P, Re, and Os compositions of 3%, 0.1%, 295 ppb, and 3250 ppb respectively we find that compositions similar to Muonionalusta are produced after ~60% fractional crystallization. For an inwardly crystallizing core [6] this suggests formation at 70% the radius, R, of the body.

**Thermal Conduction Model:** To match the wide range of IVA cooling rates we follow [1] by assuming an exposed core that cools without any insulating silicate mantle. We solve the 1D thermal conduction equation [7] for cores with variable amounts of live ⁶⁰Fe, ranging from zero [8] up to a plausible maximum of ⁶⁰Fe/⁵⁶Fe = 4 × 10⁻⁷ [9]. With core size and ⁶⁰Fe abundance as the primary variables in this model, we fit both the range of IVA cooling rates and the crystallization of Muonionalusta at 4565.3 Myr at a radius of 0.7R.

**Results:** Our calculations suggest that the IVA core was 50-110 km in radius after the hit-and-run collision. This range is primarily due to uncertainties in the initial abundance of live ⁶⁰Fe incorporated into the IVA core. This emphasizes the need to better constrain initial ⁶⁰Fe abundances for the IVAs and other iron meteorite parent bodies that may have crystallized within a few half-lives of ⁶⁰Fe after CAIs ($\tau_{1/2}$ = 2.62 Myr). Candidates for such early formation include the IIAB, IIIAB and IVB groups. Lastly, our models define a relationship between cooling rate and closure age. For example, IVAs with the fastest cooling rates (> 1000 K Myr⁻¹) may have absolute ages separated by as little as a few times 10⁵ years after CAI formation. The details of this relationship can be firmly established if old U-Pb ages are confirmed/measured for Muonionalusta and other IVAs.

**References:** [1] Yang J., Goldstein J. I., and Scott E. R. D. 2007. *Nature* 446:888-891. [2] Yang J., Goldstein J. I., and Scott E. R. D. 2008. *Geochimica et Cosmochimica Acta* 72:3042-3061. [3] Asphaug E. et al. 2006. *Nature* 439:155-160. [4] Blichert-Toft J. et al. 2010. *Earth and Planetary Science Letters* 296:469-480. [5] Walker R. J. et al. 2008. *Geochimica et Cosmochimica Acta* 73:5189-5201. [6] Ruzicka A. and Hutson M. 2006. *Meteoritics & Planetary Science* 41:1959-1987. [7] Moskovitz N. and Gaidos E. 2011. *Meteoritics & Planetary Science*, in press. [8] Quitte G. et al. 2010. *The Astrophysical Journal* 720:1215-1224. [9] Mishra R. K. et al. 2010. *The Astrophysical Journal Letters* 714:L217-L221.
THE SEARCH FOR PRESOLAR OXIDES IN PARIS

S. Mostefaoui. Muséum National d'Histoire Naturelle, LMCM UMR CNRS-7202, Case 52, 57 Rue Cuvier 75321 Paris Cedex 05, France.
E-mail: smail@mnhn.fr.

**Introduction:** Paris is considered the less aqueously altered among the CM chondrites [1]. Its material is assumed to be better preserved and therefore represents the best CM candidate for the detection of presolar oxides in the matrix. In this study I present preliminary NanoSIMS results of the O-isotopography of the Paris matrix. The estimation of presolar oxide abundance in comparison with other chondrites will allow to better understanding the alteration and thermal histories of this meteorite group.

**Experiment:** A polished thick section of Paris was observed with optical microscope. Three areas of the fine matrix largely distant from each other and reflecting the alteration lithologies of the meteorite were meticulously selected excluding zones of large grains and/or chondrule fragments. The areas were screened for the presence of presolar O-rich dust using the NanoSIMS. A focussed $Cs^+$ ion beam (<100 nm) of <1 pA was rastered over $9 \times 9 \ \mu\text{m}^2$ areas. Negative secondary ions of the three O-isotopes, $^{28}$Si, and $^{27}$Al$^{16}$O were simultaneously measured in multicollection and $256 \times 256$ pixel image sequences (total integration time of approximately 1 h per image set) were acquired on different locations.

**Results:** Seventy-six NanoSIMS $9 \times 9 \ \mu\text{m}^2$ images, covering a total matrix area of 6156 $\mu\text{m}^2$, were acquired for the search of presolar oxides in Paris. None of the measured areas show a sign of an anomalous O-isotopic hotspot suggesting the absence (at least to the level allowed by the NanoSIMS spatial resolution) of presolar oxides in these areas. This testifies the rarity of such grains in Paris. Their absence prevents determining the precise abundance of presolar oxides in this meteorite. But, assuming the finding of a presolar oxide in the subsequent $9 \times 9 \ \mu\text{m}^2$ area of the matrix with a size typical of presolar oxide grains in chondrites one can estimate an upper limit of the matrix-normalized abundance in Paris not to exceed approximately 10 ppm.

**Discussion:** The present results show that the abundance of presolar oxides in Paris is low. It is definitely lower than the abundances of presolar oxides in the CRs QUE 99177 and MET 00426 (220 and 160 ppm, respectively) [2] and in the Acfer 094 carbonaceous chondrite for which an equivalent matrix area was scanned for 12 presolar oxides identified [3]. In Murchison, which is a CM like Paris, the abundance of presolar oxides (all silicates) was estimated to approximately 3 ppm [4]. Although Paris is predictable to have a higher abundance than Murchison, one need to explore more than 2x the area I explored here in order to confirm it. Despite the Paris material is the less suffering among CM chondrites from aqueous alteration either thermal metamorphism [5] or still the aqueous alteration [4] were responsible for destruction of oxide presolar isotopic signatures in Paris. Alternatively, a preaccretionary selective destruction [6] leading to heterogeneous distributions of presolar oxide grains within carbonaceous chondrites formation environment cannot be excluded. Continuing NanoSIMS measurements are in progress in order to better estimate the abundance of presolar oxide in this chondrite.

**References:** [1] Zanda B. et al. 2010. *Meteoritics & Planetary Science* #5312. [2] Floss C. and Stadermann F. J. 2009. *Geochimica et Cosmochimica Acta* 73:2415. [3] Mostefaoui S. et al. 2004. *The Astrophysical Journal* 613:L149. [4] Nagashima K. et al. 2005. Abstract #1671. 36th Lunar and Planetary Science Conference. [5] Huss G. R. and Lewis R. S. 1995. *Geochimica et Cosmochimica Acta* 59:115. [6] Huss G. R. 2004. *Antarctic Meteorite Research* 17:132.

SYNCHROTRON X-RAY DIFFRACTION AND MÖSSBAUER SPECTROSCOPY STUDIES OF WEATHERING EFFECTS ON ORDINARY CHONDRITES FROM THE ATACAMA DESERT, CHILE

P. Munayco, J. Munayco, R. R. Avillez, M. Valenzuela, J. Gattacceca, P. Rochette and R. B. Scorzelli.

This work reports the results obtained by $ ^{57}\text{Fe} $ Mössbauer spectroscopy (MS) and synchrotron radiation X-ray diffraction (SR-XRD) on 15 meteorite samples of weathered ordinary chondrites (OC), collected in the San Juan area of the Atacama Desert, northern Chile, that continues previous work done [1] with samples of this densest collection area.

MS is a powerful tool to determine the differences in alteration rate between Fe-bearing phases in a weathered meteorite. The analyses of the spectral area of the primary phases in relation to the total amount of oxidation allow the determination of which phases are most susceptible to weathering. Since the abundance of ferric iron is directly related to the level of terrestrial weathering, the decrease in the spectral area of a primary mineral, should be followed by an increase in a ferric component. In this work the MS results are complemented by SR-XRD experiments performed in the D10A-XRD2 line of the National Synchrotron Light Laboratory (LNLS, Brazil). The use of synchrotron radiation amplifies the use of XRD since it provides very intense and sharp lines for very crystalline phases. SR-XRD is particularly suited to weathered meteorites which usually have small amounts of many corrosion phases, together with the primary phases. Further, XRD is not restricted to iron contained phases.

The Mössbauer spectra at RT exhibit two $Fe^{2+}$ doublets, associated to olivine and pyroxene, a third $Fe^{3+}$ doublet, corresponding to superparamagnetic oxides and/or oxyhydroxides, and additionally, four sextets due to magnetically ordered phases – troilite, magnetite, kamacite/ taenite and large particle goethite. At low temperature (4.2 K) the spectra exhibit the same components observed at RT and additionally, a magnetic component (not seen at RT) identified as akaganéite. An increase in the relative areas of goethite and magnetite is also observed and in some samples hematite/maghemite was also detected. The presence of akaganéite and hematite as well as the increase in relative areas, is due to the magnetic splitting at 4.2 K of the superparamagnetic $Fe^{3+}$ doublet observed at RT.

The fit of the SR-XRD results was done using the Bruker TOPAS 4.2® program with the fundamental parameters approach. A standard silicon powder was used to optimize the fundamental parameters. Up to 18 different phases were considered during the Rietveld modeling. Iron was allowed to substitute the magnesium sites of the pyroxene and olivine phases. All the other phases were considered stoichiometric or with constant composition. The goethite amount seems to provide a clear indication for meteoritic weathering. Further the total amount of weathering related phases correlates very well with its Mossbauer equivalent.

**References:** [1] J. Munayco, P. Munayco, E. M. Valenzuela, J. Gattacceca, P. Rochette, R. B. Scorzelli. 2010. Abstract #5035. 73rd Meteoritical Society Meeting, New York, USA.
ELECTRONIC STRUCTURE CALCULATIONS OF WATER ADSORPTION ON FORSTERITE GRAINS: IMPLICATIONS FOR PLANETARY WATER

Krishna Muralidharan¹ and Michael J. Drake². ¹Material Science and Engineering, The University of Arizona, Tucson, AZ, 85721, USA.
²Lunar and Planetary Laboratory, The University of Arizona, Tucson, AZ, 85721, USA.

**Introduction:** Recent computational investigations have clearly indicated that adsorption occurring during the early stages of planetary accretion is a viable delivery source of planetary water [1, 2]. In this study we use rigorous electronic structure and atomistic simulations to examine and accurately characterize the mechanisms of water adsorption on to forsterite grains-a primary planetary building block. Specifically, we focus on examining dissociative adsorption of water. Previous investigations had focused on associative molecular adsorption, which serves as a precursor to the ultimate dissociative adsorption of water.

**Results and Discussion:** Using density functional theory (DFT), the dissociation energy-landscapes of three surface orientations ({100}, {110}, {010}) of forsterite were studied. Of particular interest is the {100} surface, as it contains multiple adsorption sites. The unhydrated {100} surface, which is a primary cleavage plane of olivine, contains three possible dissociation sites (S2, S3, and S4), each characterized by the presence of undercoordinated oxygen and magnesium atoms, thereby lending themselves well to the dissociation of a water molecule. The S2 site is coordinated to three oxygen atoms, while S3 is nested between two surface tetrahedra and is bonded to four oxygen atoms. S4 is also four-coordinated, with three oxygen atoms lying on the surface. In order to calculate the dissociative adsorption energies of the three sites, plane wave-DFT utilizing the Troullier Martin pseudopotential [3] to represent the electron-nuclear interaction, in conjunction with the local density approximation [4] (LDA) for exchange-correlation was used. Amongst the 3 sites, S2 yields the highest dissociation energy, corresponding to $420 \text{ KJ mol}^{-1}$, which can be attributed to the fact that it is the least coordinated cation as well as to the fact that it is located close to the surface and is extremely susceptible to hydroxylation. S3 and S4 also yield very high dissociation energy values of $288 \text{ KJ mol}^{-1}$ and $180 \text{ KJ mol}^{-1}$, respectively. Note that molecular water can associatively bind to the {100} surface very strongly, with energies ranging between $100-150 \text{ KJ mol}^{-1}$. Thus from a computational standpoint, one can assume that on {100} both associative and dissociative adsorption can occur, though based on energetics dissociation is preferred.

**Conclusion:** Using electronic-structure calculations, we have shown that adsorption of water on to forsterite grains is energetically stable at thermodynamic conditions pertaining to the accretion disk. Further, we identify that water dissociation is the dominant adsorption mechanism and conclude that adsorption can be a significant source of terrestrial planetary water. The conclusions of Muralidharan et al. [2] that Earth could acquire many oceans of water by adsorption remain robust. Terrestrial planetary systems around other stars should also contain water.

**References:** [1] King M. E. et. al. 2010. *Earth and Planetary Science Letters* 300:11. [2] Muralidharan K. et al. 2008. *Icarus* 198:400. [3] Troullier N. and Martin J. L. 1991. *Physical Review B* 43:1993. [4] Ceperley D. M. and Alder B. J. 1980. *Physical Review Letters* 45:566.

CONDENSATION KINETICS AND CHEMICAL FRACTIONATION OF SILICATE AND METAL IN A PROTOPLANETARY DISC

H. Nagahara¹ and K. Ozawa². ¹Department Earth Planetary Science, The University Tokyo, Tokyo, Japan. E-mail: hiroko@eps.s.u-tokyo.ac.jp. ²Department Earth Planetary Science, The University Tokyo, Tokyo, Japan.

**Introduction:** Condensation is the first step of planet formation. Although recent finding of CAIs in cometary samples [1, 2] have been thought to represent large scale material transportation from the inner edge to outer region of the disc, it is not clear whether CAIs were formed at the inner edge of the disc or by transient heating in the outer region. Rather, the presence of a large variety of chondrites prefers mixing of dusts in a limited scale. Time scales of condensation and physical evolution of dusts in a range of total pressure, gas cooling time scale, and dust/gas separation conditions are evaluated by using a kinetic condensation model including homogeneous and heterogeneous condensation cases.

**Model:** The model is based on kinetic condensation theory in cooling gas [3-5] with solar composition (H-He-C-N-O-Mg-Al-Si-Fe) developed by our group is applied. The model calculates the changes of dust size, number density, chemical composition, occurrence of different phases with time, where cooling time scale and critical size for dust separation from the residual gas are free parameters. We use experimentally determined kinetic parameters for condensation and evaporation by our group.

**Results and Discussion:** If condensation proceeds without dust separation, the condensed dusts have a very narrow size range and the system finally does not show chemical fractionation, where the grain size of dusts is exclusively controlled by cooling time scale of the system. The grain size distribution and bulk composition of both residual system and separated dusts showed a wide range when dust are separated from the residual gas when they reached a critical size, which is due to repeated homogeneous nucleation. If metal grains with size as small as those we find in GEMS (tens of nanometer in size), the cooling condition is fairly well constrained, which is comparable to the chondrule cooling time scale. Thus, the critical size for dust/gas separation plays an important role for chemical fractionation and dust size distribution. The mode of nucleation (homogeneous or heterogeneous) is found to be the most crucial parameter for chemical fractionation: totally heterogeneous condensation does not explain Mg/Si/Al fractionation among chondrites, suggesting Al2O3 separation in the early stage of condensation in order to explain compositions of chondrites.

**References:** [1] Brownlee D. 2006. *Science* 314:1711-1716. [2] McKeegan K. D. et al. 2006. *Science* 314:1724-1728. [3] Yamamoto T. and Hasegawa H. 1977. *Progress of Theoretical Physics* 58:816. [4] Drain B. T. and Salpeter E. E. 1977. *The Journal of Chemical Physics* 67:2230. [5] Kozasa and Hasegawa. 1989. *Progress of Theoretical Physics* 77:1402.
## SOLAR NOBLE GASES IN ITOKAWA REGOLITH MATERIALS RETURNED BY THE HAYABUSA MISSION

K. Nagao¹, R. Okazaki², T. Nakamura³, Y. N. Miura⁴, T. Osawa⁵, K. Bajo¹, S. Matsuda¹, M. Ebihara⁶, T. R. Ireland⁷, F. Kitajima², H. Naraoka², T. Noguchi⁸, A. Tsuchiyama⁹, M. Uesugi¹², H. Yurimoto¹⁰, M. E. Zolensky¹¹, K. Shirai¹², M. Abe¹², T. Yada¹², Y. Ishibashi¹², A. Fujimura¹², T. Mukai¹², M. Ueno¹², T. Okada¹², M. Yoshikawa¹² and J. Kawaguchi¹². ¹Geochemical Research Center, University Tokyo, Tokyo, Japan. E-mail: nagao@eqchem.s.u-tokyo.ac.jp. ²Department Earth Planetary Science, Kyushu University, Kyushu, Japan. ³Department Earth Planetary Mater. Science, Tohoku University, Tohoku, Japan. ⁴Earthquake Research Institute, University Tokyo, Tokyo, Japan. ⁵Japan Atomic Energy Agency. ⁶Graduate School Science Engineering, Tokyo Metropolitan University, Tokyo, Japan. ⁷Research School Earth Science, Australian National University. ⁸College Science, Ibaraki University, Ibaraki, Japan. ⁹Department Earth Space Science, Osaka University, Osaka, Japan. ¹⁰National Hist. Science, Hokkaido University, Hokkaido, Japan. ¹¹NASA JSC. ¹²JAXA-ISAS.

**Introduction:** The Hayabusa spacecraft successfully returned regolith materials from the surface of S-type asteroid 25143 Itokawa [1]. The samples were undamaged by the collection on the Itokawa surface, and handled in the clean chambers of the JAXA curation facility under vacuum and nitrogen gas [2], which is essentially different from other extraterrestrial materials. We report noble gas elemental and isotopic compositions of three grains RA-QD02-0015 (#0015), RA-QD02-0053 (#0053), and RA-QD02-0065 (#0065).

**Experimental Procedures:** The samples are blocky and translucent olivine grains with dimensions of approximately 40, 40, and 60 µm [3], and their estimated masses are 0.06, 0.06, and 0.21 µg for #0015, #0053, and #0065, respectively. Noble gases were extracted from each grain by means of stepwise heating at approximately 200 and 300°C, and finally melted for complete extraction using a Nd-YAG CW laser. The extracted noble gases were measured with a modified VG5400/MS-III at the Geochemical Research Center, University of Tokyo.

**Results and Implications:** High concentrations of He, Ne, and Ar, as high as those in lunar soils and IDPs, are released at variable temperatures from each sample. Their isotopic compositions are essentially identical to those of solar wind (SW). This clearly confirms that the grains resided on the surface of Itokawa. Only He is depleted compared with the relative abundances of He, Ne, and Ar in SW. These noble gas characteristics can be explained by repeated implantation and preferential loss of solar He by removal of weathered He-rich rim on the grain surface [4] through friction among regolith grains on Itokawa. Duration of SW implantation to the measured samples is estimated as several hundreds of years. Concentrations of cosmogenic noble gases produced by galactic cosmic rays (GCR) were so low that they could not be detected in excess to the SW isotopic ratios beyond experimental errors. An upper limit of GCR exposure ages for the Hayabusa samples based on the Ne isotopic compositions is approximately 10 Myr. The short exposure age suggests that the regolith materials have been continuously escaping from Itokawa's surface at the relatively high rate of several tens of cm Myr⁻¹. The weak gravitational force at Itokawa's surface (escape velocity approximately 0.2 m s⁻¹) [5] should be a main reason for the short residence time of the Itokawa regolith grains.

**References:** [1] Yano H. et al. 2006. *Science* 312:1350–1353. [2] Okazaki R. et al. 2011. Abstract #1653. 42nd Lunar and Planetary Science Conference. [3] Nakamura T. et al. 2011. *Science* (submitted). [4] Noguchi T. et al. 2011. *Science* (submitted). [5] Miyamoto H. et al. 2007. *Science* 316:1011–1014.

## GLOBAL DISTRIBUTION OF SILICON BY KAGUYA GAMMA-RAY SPECTROMETER

H. Nagaoka¹, N. Hasebe¹, K. J. Kim², Y. Karouji³, S. Kobayashi³, M. Hareyama³, N. Yamashita⁴, O. Gasnault⁴, O. Forni⁴, C. d'Uston⁴, R. C. Reedy⁵, E. Shibamura⁶, M. N. Kobayashi⁷, H. Takeda⁸, and the Kaguya Gamma Ray Spectrometer team. ¹Waseda University, Waseda, Japan. E-mail: hiroshi-nagaoka@asagi.waseda.jp. ²KIGAM. ³JAXA. ⁴IRAP. ⁵PSI. ⁶Saitama Pref. University, Saitama, Japan. ⁷Chiba Institute Tech, Chiba, Japan. ⁸University of Tokyo, Tokyo, Japan.

**Introduction:** The distribution of major and important trace elements on the lunar surface is essential to understand the origin and evolution of the lunar crust. Gamma-ray spectrometers (GRS) carried on previous lunar and planetary missions have played important roles to determine the chemical element abundance of the surface of planetary body [e.g., 1]. A high precision GRS was carried on the first Japan's large-scaled lunar explorer, Kaguya which observed from December 2007 to June 2009 [2]. The GRS mounted on Kaguya (KGRS), which was the first to use a germanium detector for a lunar mission [3], has measured gamma rays emitted from the Moon with the high energy resolution. The distribution maps of three radioactive elements were reported, based on KGRS data analysis (K [4], Th, and U [5]). Preliminary results for Fe, Ti, Ca, and Al have been reported [6–8]. In this work, we present preliminary results of analysis for a thermal neutron capture reaction peak from Si using KGRS data obtained during the low altitude observation.

**Data Processing:** KGRS data obtained in the lower altitude observation (approximately 50 km from February to June, 2009) in the late portion of its mission were used. First, the lunar surface was pixelated by 600 × 600 km equal areas, gamma-ray counts were accumulated in each pixel, and 106 spectra were made. Noises in resultant spectra were filtered out. The Si distribution was surveyed using 3539.5 keV peak, which is well resolved by KGRS spectra. The net counting rates of this peak were determined on the basis of the total peak area method [9]. Neutron correction was applied to the net counting rates using Lunar Prospector (LP) thermal neutron data [10, 11].

**Results:** The 3539.5 keV neutron-capture peak from Si was used to map Si over the entire lunar surface. The counting rates of Si in mare region (Procellarum, Imbrium, Tranquillitatis) seem to be lower, as compared with highlands. The counting rates of Si in the farside highland regions near Hertzsprung were slightly high, as compared with their surrounding highland regions. Ample studies of this result are required in future.

**References:** [1] Prettyman T. H. et al. 2006. *Journal of Geophysical Research* 111:E12007. [2] Kato M. et al. 2010. *Space Science Review* 154:3–19. [3] Hasebe N. et al. 2008. *Earth Planets, and Space* 60:299–312. [4] Kobayashi S. et al. 2010. *Space Science Review* 154:193–218. [5] Yamashita N. et al. 2010. *Geophysics Research Letter* 37:L10201. [6] Gasnault O. et al. 2009. Abstract #2253. 40th Lunar and Planetary Science Conference. [7] Karouji Y. et al. 2010. Abstract #5166. 73rd Annual Meteoritical Society Meeting. [8] Yamashita N. et al. 2011. Abstract #2093. 42nd Planetary Science Conference. [9] Gilmore G. and Hemingway J. D. 1995. *Practical Gamma-ray Spectrometry*. Wiley, Chichester. [10] The Planetary Data System, http://pds.nasa.gov/ [11] Maurice S. et al. 2004. *Journal of Geophysical Research* 109:E07S04.
## OXYGEN-ISOTOPE COMPOSITIONS OF CHONDRULES AND MATRIX GRAINS IN KAKANGARI CHONDRITE

K. Nagashima¹, A. N. Krot¹ and G. R. Huss¹. ¹HIGP/SOEST, University of Hawai'i at Mānoa, USA. E-mail: kazu@higp.hawaii.edu.

**Introduction:** The relationship between chondrules and matrix in primitive chondrites is not well understood. Some workers see matrix as raw material for the chondrules that did not experience chondrule-forming process [e.g., 1], while others see matrix as complementary to chondrules in the sense that material that evaporated from the chondrules during chondrule melting events re-condensed in the matrix [e.g., 2]. Oxygen isotopes can potentially provide important constraints on the relationship between chondrules and matrix and on thermal processing of matrix in the solar nebula. Here we report on an in situ O-isotope study of chondrules and matrix grains in the Kakangari chondrite. The Kakangari chondrules and matrix have similar chemical compositions but slightly different O-isotope compositions: matrix is ¹⁶O-enriched by 2.6‰ compared to the bulk chondrules which plot along the terrestrial fractionation (TF) line [3–5].

**Methods:** Oxygen-isotope compositions were measured with the UH Cameca ims-1280 SIMS. For chondrules, we used ~1 nA primary Cs⁺ beam with 7 µm raster and 2 Faraday cups (FCs) for ¹⁶O and ¹⁸O, and an electron multiplier (EM) for ¹⁷O. Matrix grains were measured with ~15 pA beam focused to ~1 µm using FC-EM-EM for ¹⁶O, ¹⁷O, and ¹⁸O, respectively. Analysis spots were marked by focusing the electron beam on the points of interest to produce a spot that is visible as depleted oxygen signal in scanning ¹⁶O⁻ ion image, allowing us to identify exact locations of small grains in matrix.

**Results and Discussion:** Oxygen-isotope compositions of olivine (Fo~.97) and low-Ca pyroxene (En~.95-.97) in five chondrules (Type I PO and POP) plot near the intersection of the TF line and the Young and Russell line [6] with mean Δ¹⁷O values of +0.2 ± 0.6‰ (2σ) and -0.1 ± 0.7‰, respectively. There are no significant systematic differences among chondrules. The Δ¹⁷O values of chondrule phenocrysts are in good agreement with O-isotope compositions of bulk chondrules, but the δ¹⁸O values are slightly lower than those reported by [3]. One olivine grain in a coarse-grained chondrule rim is ¹⁶O-rich (Δ¹⁷O ~ -23‰). Oxygen-isotope compositions of matrix olivine (Fo~.96-.97) and low-Ca pyroxene (En~.94-.96) grains of ~5–20 µm in size have a bimodal distribution: 12 out of 13 olivine and 4 out of 17 pyroxene grains are similarly ¹⁶O-rich (Δ¹⁷O ~ -23.5 ± 2.9‰), others are similarly ¹⁶O-poor (Δ¹⁷O ~ -0.1 ± 1.7‰). These results indicate the ¹⁶O-enriched bulk matrix composition compared to bulk chondrules [3] could be explained by ubiquitous presence of ¹⁶O-rich grains in matrix. Nearly identical O-isotope compositions of chondrules and ¹⁶O-poor matrix grains may imply that both components formed in the same nebular region. ¹⁶O-rich matrix olivine and pyroxene might be related to AOAs and perhaps formed at a different time and place from Kakangari chondrules, suggesting those grains escaped thermal processing during chondrule formation and were mixed with Kakangari ¹⁶O-poor matrix grains later.

**References:** [1] Huss G. R. 2003. *Geochimica et Cosmochimica Acta* 67:4823–4848. [2] Klerner S. and Palme H. 1999. *Meteoritics & Planetary Science* 34:A64. [3] Prinz M. et al. 1989. 20th Lunar and Planetary Science Conference, pp. 870–871. [4] Brearley A. J. 1989. *Geochimica et Cosmochimica Acta* 53:2395–2411. [5] Weisberg M. K. et al. 1996. *Geochimica et Cosmochimica Acta* 60:4253–4263. [6] Young E. D. and Russell S. S. 1998. *Science* 282:452–455.

## MINERALOGY AND THERMAL HISTROY OF ITOKAWA SURFACE PARTICLES RECOVERED BY HAYABUSA MISSION

T. Nakamura¹, T. Noguchi², M. Tanaka³, M. E. Zolensky⁴, M. Kimura², A. Tsuchiyama⁵, A. Nakato¹, T. Ogami¹, H. Ishida¹, M. Uesugi⁶, T. Yada⁶, K. Shirai⁶, A. Fujimura⁶, R. Okazaki⁸, S. Sandford⁷, S. Wakita¹, Y. Ishibashi⁶, M. Abe⁶, T. Okada⁶, M. Ueno⁶, T. Mukai⁶, M. Yoshikawa⁶ and J. Kawaguchi⁶. ¹Department Earth & Planet. Mat. Sci., Tohoku University, Tohoku, Japan. E-mail: tomoki@m.tohoku.ac.jp. ²The College of Science at Ibaraki University, Ibaraki, Japan. ³WEBRAM,SPring-8, National Institute for Maths Science. ⁴NASA JSC. ⁵Department Earth & Space Science, Osaka University, Osaka, Japan. ⁶JAXA-ISAS. ⁷NASA ARC. ⁸Department Earth & Planetary Science, Kyushu University, Kyushu, Japan.

**Introduction:** Surface dust particles of asteroid Itokawa were successfully collected by spacecraft Hayabusa [1], but are very small similar in size range to the particles recovered by Stardust mission. Analysis of these tiny particles required multidiscipline edge-of-the-art techniques. We have developed a sequential experimental procedure consisting of synchrotron-based analysis, scanning and transmission electron microscopy, and secondary ion mass spectrometry [2].

**Experiments:** Approximately forty particles of about 30–180 µm in size were attached on the top of thin carbon fibers. Layers of epoxy resin with several-micron thick covered entire surface of the particle to minimize exposure to atmosphere. The particles were analyzed at beamline 3A of KEK and BL15XU of SPring-8 for X-ray diffraction. The particles were embedded in epoxy resin and processed by ultramicrotomy and/or polishing. Polished sections of the particles were analyzed by scanning electron microscopes for chemical composition and EBSD. Areas of interest in some particles were lifted up and thinned by FIB for observation by a transmission electron microscope.

**Results and Discussion:** Most particles are composed mainly of relatively coarse (typically from 10 to 50 µm in size) and well-crystalline silicates such as olivine, low- and high-Ca pyroxene, and plagioclase. Some particles contain chromite, chlorapatite, merrillite, and troilite up to 20 µm in size. Small inclusions (up to 10 µm) of taenite, kamacite, troilite, and chromite occur within coarse silicates in many particles. Compositions of silicates are relatively homogeneous. The common occurrence of triple junctions suggests that most particles have been thermally annealed sufficiently long for the silicates to have recrystallized. Olivine approximately Fa₂₉, low-Ca pyroxene approximately FS₂₂, and kamacite Co (approximately 10 wt%) and Ni (approximately 4 wt%) concentrations of Itokawa particles are within the range of LL chondrites [3]. The presence of coarsened plagioclase (typically 20–50 µm in size) and diopside (typically 10–30 µm) suggests that most of the particles have experienced high-temperature metamorphism with a metamorphic grade from 5 to 6. Metamorphic temperatures were estimated based on two-pyroxene geothermometry to be approximately 800 °C, which is slightly lower than LL6 temperatures [4].

**References:** [1] Nakamura T. et al. 2011. Abstract #1776. 42nd Lunar and Planetary Science Conference. [2] Tsuchiyama A. et al. 2011. Abstract #1777. 42nd Lunar and Planetary Science Conference. [3] Rubin A. E. *Geochimica et Cosmochimica Acta* 54:1217–1232. [4] Nakamura Y. and Motomura Y. 1999. *Meteoritics & Planetary Science* 34:763–772.