Origin and history of ureilitic material in the solar ...web.colby.edu/tldunn/files/2018/06/Ruofei_reading.pdfchondrites were implanted by low-velocity collisions. (5) Recently, the

Origin and history of ureilitic material in the solar system: The view from

asteroid 2008 TC3 and the Almahata Sitta meteorite

Cyrena Anne GOODRICH1*, William K. HARTMANN1, David P. O’BRIEN1,Stuart J. WEIDENSCHILLING2, Lionel WILSON2, Patrick MICHEL3, and Martin JUTZI4

1Planetary Science Institute, 1700 E. Ft. Lowell, Tucson, Arizona 85719, USA2Lancaster Environment Centre, Lancaster University, Lancaster LA1 4YQ, UK

3Observatoire de la Cote d’Azur, UMR 6202 Cassiop�ee, CNRS, BP 4229, 06304 Nice Cedex 4, France4Physics Institute, Space Research and Planetary Sciences, Center for Space and Habitability, University of Bern, Bern,

Switzerland*Corresponding author. E-mail: [email protected]

(Received 14 February 2014; revision accepted 27 October 2014)

Abstract–Asteroid 2008 TC3 (approximately 4 m diameter) was tracked and studied in spacefor approximately 19 h before it impacted Earth’s atmosphere, shattering at 44–36 kmaltitude. The recovered samples (>680 individual rocks) comprise the meteorite AlmahataSitta (AhS). Approximately 50–70% of these are ureilites (ultramafic achondrites). The restare chondrites, mainly enstatite, ordinary, and Rumuruti types. The goal of this work is tounderstand how fragments of so many different types of parent bodies became mixed in thesame asteroid. Almahata Sitta has been classified as a polymict ureilite with an anomalouslyhigh component of foreign clasts. However, we calculate that the mass of fallen materialwas ≤0.1% of the pre-atmospheric mass of the asteroid. Based on published data for thereflectance spectrum of the asteroid and laboratory spectra of the samples, we infer that thelost material was mostly ureilitic. Therefore, 2008 TC3 probably contained only a fewpercent nonureilitic materials, similar to other polymict ureilites except less wellconsolidated. From available data for the AhS meteorite fragments, we conclude that 2008TC3 samples essentially the same range of types of ureilitic and nonureilitic materials asother polymict ureilites. We therefore suggest that the immediate parent of 2008 TC3 wasthe immediate parent of all ureilitic material sampled on Earth. We trace critical stages inthe evolution of that material through solar system history. Based on various types of newmodeling and re-evaluation of published data, we propose the following scenario. (1) Theureilite parent body (UPB) accreted 0.5–0.6 Ma after formation of calcium-aluminum-richinclusions (CAI), beyond the ice line (outer asteroid belt). Differentiation beganapproximately 1 Ma after CAI. (2) The UPB was catastrophically disrupted by a majorimpact approximately 5 Ma after CAI, with selective subsets of the fragments reassemblinginto daughter bodies. (3) Either the UPB (before breakup), or one of its daughters (afterbreakup), migrated to the inner belt due to scattering by massive embryos. (4) One daughter(after forming in or migrating to the inner belt) became the parent of 2008 TC3. Itdeveloped a regolith, mostly ≥3.8 Ga ago. Clasts of enstatite, ordinary, and Rumuruti-typechondrites were implanted by low-velocity collisions. (5) Recently, the daughter wasdisrupted. Fragments were injected or drifted into Earth-crossing orbits. 2008 TC3 comesfrom outer layers of regolith, other polymict ureilites from deeper regolith, and main groupureilites from the interior of this body. In contrast to other models that have beenproposed, this model invokes a stochastic history to explain the unique diversity of foreignmaterials in 2008 TC3 and other polymict ureilites.

Meteoritics & Planetary Science 50, Nr 4, 782–809 (2015)

doi: 10.1111/maps.12401

782© The Meteoritical Society, 2014.

INTRODUCTION

Nobody has ever figured out ureilites, not even me.

— Mike Drake (October, 2008).

On October 6, 2008, a small (approximately 4 mdiameter) asteroid subsequently named 2008 TC3 wasdiscovered and recognized to be on a trajectory to hitEarth (Jenniskens et al. 2009; Kwok 2009; Shaddadet al. 2010). In the approximately 19 h following thediscovery, a number of astronomical observations of theasteroid were made. Its orbit was determined veryprecisely (McGaha et al. 2008), lightcurve data wereobtained (Scheirich et al. 2010; Kozubal et al. 2011),and a reflectance spectrum in the 550–1000 nmwavelength interval was measured (Jenniskens et al.2009, 2010).

Asteroid 2008 TC3 hit the Earth’s atmosphere overthe desert of northern Sudan in the early hours ofOctober 7, 2008. The fireball created by the impact wasfirst detected at approximately 65 km altitude. A fewseconds later, the asteroid shattered into fragments in aseries of explosions. The main detonation occurred atapproximately 37 km altitude, which is much higherthan fireballs of typical ordinary chondrite compositionand therefore suggests that the asteroid had relativelylow bulk strength (Shaddad et al. 2010; Popova et al.2011). Dust clouds resulting from the explosions wereobserved between approximately 44 and 35 km(Borovi�cka and Charv�at 2009; Jenniskens et al. 2009).No significant dust cloud formation was observed belowabout 33 km, and no meteorites were expected tosurvive (Jenniskens et al. 2009). Nevertheless, organizedsearch campaigns in the predicted fall area led torecovery of >680 individual meteorite fragments, eachapproximately 0.2–400 g in mass (Bischoff et al. 2010a,2010b; Jenniskens and Shaddad 2010; Shaddad et al.2010; Horstmann and Bischoff 2014). The meteorite fallwas named Almahata Sitta (AhS), after a local placename. AhS is the first meteorite to be recovered froman asteroid that was tracked and studied in space beforehitting the Earth.

The AhS sample collection is unusual because itincludes a wide variety of different meteorite types.Although none of the samples studied so far showcontacts between these different types, analyses of short-lived cosmogenic radionuclides and low weatheringgrades strongly suggest that they all belong to the samefall (Bischoff et al. 2010a, 2010b; Kohout et al. 2010),and therefore resided together in the asteroid before ithit Earth. Based on the 110 AhS samples that havebeen studied petrographically and chemically,approximately 69% of the AhS samples (approximately79% by mass) belong to the ureilite group of

achondrites and 31% (21% by mass) are diverse typesof chondrites (Horstmann and Bischoff [2014] andreferences therein). Estimates based on hand-specimenproperties of a larger collection of approximately 600samples suggest up to approximately 50% chondriticsamples (Kohout et al. 2010; Shaddad et al. 2010).Among the chondritic samples, enstatite chondrites(EC) are the most abundant, including both EL (lowiron), and EH (high iron) types, as well as E-brecciasand impact melt rocks (Bischoff et al. 2010a, 2010b;Horstmann et al. 2012; Horstmann and Bischoff 2014).Ordinary chondrites (OC) of both the L (low iron) andH (high iron) groups have also been found, as well asone carbonaceous chondrite (CC) of CB (Bencubbin)type, and one Rumuruti (R)-type chondrite (Bischoffet al. 2010a, 2010b, 2012; Horstmann et al. 2010, 2012;Zolensky et al. 2010; Horstmann and Bischoff 2014).Never before have individual meteorite samples of sucha wide range of types been recovered from a single fall.

Almahata Sitta has been classified as a polymictureilite with an anomalously high component of foreignmaterials (Horstmann and Bischoff 2014). The diversityof these foreign materials is remarkable. They includesamples of nearly every major chondrite group,representing at least five different parent asteroids and awide range of chemical and isotopic environments in theearly solar system. Our goal in this work is tounderstand when and where a ureilite parent bodyformed in the solar system, how fragments from thatbody became mixed with fragments of many other typesof meteorite parent bodies, and why this process ofmixing did not happen to, or at least is not apparent in,most other meteorites.

Using published observations of the asteroid, themeteorite strewn field, and the AhS samples, we inferthat asteroid 2008 TC3, in contrast to AhS, containedonly a small component (a few percent or less) of diverseforeign materials, similar to other polymict ureilites. Weidentify common properties among the ureilitic materialsin AhS, the collection of main group ureilites, and theureilitic materials in other polymict ureilites, andbetween the foreign materials in AhS and other polymictureilites. Based on these commonalities, we hypothesizethat the immediate parent of 2008 TC3 was also theimmediate parent of all ureilitic material that has beensampled by meteorites on Earth, and that using 2008TC3/Almahata Sitta as a guide to the end point we cantrace the history of this material through many stages ofsolar system history.

Based on various types of evidence and modeling,we infer the following scenario (Fig. 1) (1) early (≤1 Maafter CAI) accretion of a ureilite parent body (UPB) inthe outer asteroid belt, with differentiation beginningsoon thereafter; (2) catastrophic impact disruption (at

Origin and history of ureilitic material 783

approximately 5 Ma after CAI), leading to formation ofureilite daughter bodies (UDBs) from subsamples of theUPB; (3) perturbational migration, of either the UPBbefore disruption or one or more UDBs afterdisruption, into the inner asteroid belt; (4) formation ofa regolith, containing clasts of diverse foreignimpactors, on a UDB; and (5) subsequent collisionalevolution and breakup of a UDB to deliver all knownureilitic material (main group ureilites, polymictureilites, and Almahata Sitta) to Earth. In the finalsection of the paper, we compare this model for theformation of 2008 TC3 with other models that havebeen discussed in the literature (e.g., Herrin et al. 2010;Bischoff et al. 2010b; Gayon-Markt et al. 2012; Meieret al. 2012; Horstmann and Bischoff 2014).

METHODS

In this paper, we develop a new model for theorigin and history of ureilitic material in the solarsystem (presented in preliminary form in Hartmannet al. 2011), informed by the discovery of asteroid 2008TC3 and the fall of Almahata Sitta. The investigationscarried out in developing this model include (1) review

and synthesis of various types of published data and useof these data to formulate new hypotheses, (2) newmodeling of several types (including petrologic, thermal,dynamical, and collisional), and (3) evaluation ofvarious aspects of this model compared with others thathave been proposed. Each of the major sections of thepaper begins with a brief statement of the main goaland methods used in that section.

THE ALMAHATA SITTA METEORITE—AN

ANOMALOUS POLYMICT UREILITE?

In this section, we first review key properties ofureilites. Then we use literature data for the AhSsamples to compare them with ureilites and reevaluatethe classification of the Almahata Sitta meteorite as ananomalous polymict ureilite.

Brief Description of Ureilites

Ureilites are the second largest group of achondrites(>350 individual meteorites are currently classified asureilites in the Meteoritical Bulletin Database),comprising approximately 95% main group ureilites and

UPB[a] Accretion & differentiationin outer belt( 1 Ma after CAI)

Major impactCatastrophic disruption:shattering and dispersal

Reaccumulation ofdaughter bodies (UDB)from select sample of UPB

[d] Migration ofa UDB toinner belt?

[e] Developmentof diverse regolithon the UDB ( 3.8 Ga)

R

EL

EH

OC

[f] Recent breakup of UDB, drift of fragmentsinto Earth-crossing orbits( 50 Ma ago)2008 TC3 Other

PolymictUreilites

Main Group Ureilites

[b] Migration ofUPB toinner belt?

CB

[c] Catastrophic disruption and reassembly (~5 Ma after CAI)

≤

≤

≥

Fig. 1. Stages in the history of ureilitic material in the solar system, as inferred in this paper. a) The UPB accreted 0.5–6 Maafter CAI, most likely in the outer asteroid belt. Shortly thereafter, the asteroid was heated by decay of 26Al and partiallydifferentiated. A hypothetical resulting compositional structure is shown (colors corresponding to ureilites of different Fo values).b) The UPB may have migrated to the inner asteroid belt after it formed. c) A major impact at approximately 5 Ma after CAIcatastrophically disrupted the UPB, resulting in shattering and dispersal, followed by gravitational reaccumulation to formoffspring bodies of select material. d) An alternative to step (b) is that a daughter body migrated to the inner belt aftercatastrophic disruption of the UPB. e) A deep regolith, composed of mixed ureilitic lithologies and diverse “foreign” clastsdeveloped on the UDB. f) In recent times, this body broke up and pieces of it drifted into Earth-crossing orbits to become maingroup ureilites (the interior), normal polymict ureilites (deep regolith), and 2008 TC3 (shallow regolith).

784 C. A. Goodrich et al.

approximately 5% polymict ureilites. Main groupureilites are coarse-grained ultramafic rocks consistingprincipally of olivine and pyroxenes (dominantly low-Catypes). They are characterized by high abundances (upto approximately 7 wt%) of carbon (dominantlygraphite), with metal and sulfides as accessory phases(see reviews in Goodrich 1992; Mittlefehldt et al. 1998;Goodrich et al. 2004; Krot et al. 2013). Their olivinegrains typically have reduced rims (up to approximately10 lm wide) and veins formed by local reactions withadjacent graphite. Silicate mineral compositions (exceptfor reduction rims) are very homogeneous within eachureilite, and record high equilibration temperatures of1200–1300 °C. However, among samples, they show alarge variation in Fo (=mg# = molar MgO/[MgO+FeO])of olivine from approximately 74 to 95, which cannotbe explained by igneous fractionation (Mittlefehldt et al.1998). A notable feature is that although thedistribution of samples between Fo 74 and Fo 95 isnearly continuous, there is a large peak in thedistribution around Fo 79–81 (Fig. 2). Ureilites alsoshow variation in pyroxene composition (i.e., Wocontent), and ureilites have been subdivided intopigeonite (Wo 5–13)-, orthopyroxene (Wo ≤5)-, andaugite-bearing types (Goodrich et al. 2004). Among thelow-Ca pyroxene ureilites, abundance of pyroxeneincreases (from approximately 10 to 60 vol%), andpyroxene type changes from pigeonite to orthopyroxene,with increasing Fo (Singletary and Grove 2003;Goodrich et al. 2007). In terms of oxygen isotope

compositions, ureilites differ from any other group ofachondrites in showing D17O heterogeneity, with valuesthat range from approximately �0.2 to �2.5 along theslope approximately 1 CV chondrite (CCAM) line(Fig. 3). Oxygen isotope compositions are alsocorrelated with Fo (decreasing D17O with increasingFo). Ureilites are generally interpreted to representresidues from approximately 15–30% partial melting ofa carbon-rich asteroid (Boynton et al. 1976; Takeda1987; Warren and Kallemeyn 1992; Scott et al. 1993;Goodrich 1999; Goodrich et al. 2007).

Polymict ureilites are fragmental and/or regolithbreccias, consisting of ≥95% lithic and mineral clastsof materials that are indistinguishable from main groupureilites, except possibly in being more shocked (seereviews in Goodrich et al. 2004; Downes et al. 2008).Both within single samples and among samples, theseclasts span the same range and distribution of ureilitetypes (Fo values, with correlated pyroxene type andoxygen isotope composition) as the main group ureilitecollection on the whole (Fig. 2). The remainder of thematerial consists of a few % feldspathic clasts thathave been shown by oxygen isotopes to be related toureilites, plus a few % nonindigenous clasts. Thelatter include fragments of ordinary, enstatite, andRumuruti-type chondrites; fragments of angrites; and

70 75 80 85 90 95 1000

10

20

30

40

50

60

Fo of olivine

234 main group ureilites

Mineral clasts in other polymict ureilites Coarse-grained ureilitic samples of AhS

Fig. 2. Histogram of olivine core compositions (Fo) for 234main group ureilites (data from many sources), compared with38 coarse-grained ureilitic samples of Almahata Sitta(tabulated in Horstmann and Bischoff 2014) and 324 olivineclasts in other polymict ureilites (Downes et al. 2008).

-8 -6 -4 -2 0 2 4 6 8 10 12-8

-6

-4

-2

0

2

4

6

CVCO

CR

EC

E-M

R

CM

OCH

LLL

PCMYR

CCAM

main group ureilitesureilitic samples in AhS

δ18OSMOW (‰)

δ17O

SMO

W(‰)

Fig. 3. Oxygen isotope compositions of main group ureilites(Clayton and Mayeda 1996) and 18 ureilitic samples fromAlmahata Sitta (Bischoff et al. 2010b; Rumble et al. 2010),compared with compositional fields of various chondritegroups (Clayton et al. 1984, 1991; Clayton and Mayeda 1999;Bischoff et al. 2011). Ureilites show significant heterogeneityin D17O (=d17O � 0.52 9 d18O), unlike any other achondrites.Black circle = bulk Earth composition from Eiler (2001).CCAM = carbonaceous chondrite anhydrous mineral line(Clayton and Mayeda 1999); E-M = Earth-Moon; YR =proposed array of primordial oxygen isotope reservoirs fromYoung and Russell (1998); PCM = primitive chondrite mineralline from Ushikubo et al. (2012).


dark clasts that resemble CC-matrix materials (Jaquesand Fitzgerald 1982; Prinz et al. 1986, 1987, 1988;Ikeda et al. 2000; Cohen et al. 2004; Kita et al. 2004;Downes et al. 2008). Solar wind implanted gases arepresent in the three polymict ureilites that have beenanalyzed for noble gases, distinguishing them asregolith breccias (Bischoff et al. 2006). Because allpolymict ureilites are petrographically similar, they areinferred to have all formed in the same surface ornear-surface regolith environment (Goodrich et al.2004; Downes et al. 2008).

The Ureilitic Almahata Sitta Samples

The classification of Almahata Sitta as ananomalous polymict ureilite was based on petrographicand chemical studies of a single sample, AhS #7, whichis an unusually fine-grained, porous ureilitic sampleconsisting of granoblastic domains of pyroxene andolivine (Jenniskens et al. 2009). The presence of olivine-rich and pyroxene-rich domains led to the perception ofmultiple lithologies, i.e., the polymict nature of thesample (Jenniskens et al. 2009). However, more detailedstudies (Zolensky et al. 2010) showed that both domainsare very similar to the mosaicized textures found inhighly shocked main group ureilites (Berkley et al. 1980;Takeda et al. 1985; Berkley 1986; Saito and Takeda1990; Ogata et al. 1991; Warren and Rubin 2010).Thus, it seems likely that these domains representoriginally larger single grains of olivine and pyroxene ina typical unshocked main group ureilite (a similarconclusion was reached by Horstmann and Bischoff2014). In this case, AhS #7 is not in itself a polymictureilite, which is important because it underscores theobservation that all AhS samples studied so far aremono-lithologic.

Approximately 75 ureilitic samples of AhS havenow been studied. Based on the main data setsavailable (Bischoff et al. 2010b; Goodrich et al. 2010a;Herrin et al. 2010; Mikouchi et al. 2010; Warren andRubin 2010; Zolensky et al. 2010; Horstmann et al.2012; Hutchins and Agee 2012; Bischoff et al. 2013),and the recent review by Horstmann and Bischoff(2014), the AhS ureilite samples can be divided intofour types (1) 38 coarse-grained, compact olivine-pyroxene samples showing a range of Fo contents,olivine/pyroxene ratios, pyroxene types, and grain sizes;(2) 27 fine-grained, highly porous samples (of whichAhS #7 is an example) which exhibit granoblasticmosaic textures and variations in grain size, and arecommonly highly reduced; (3) three metal+sulfide-richsamples with enclosed ureilitic silicate portions; and (4)one feldspar-rich, igneous-textured sample of andesiticcomposition.

Any of the coarse-grained ureilitic AhS samples thathave been studied could be individually classified as amain group ureilite in terms of mineralogy, textures,mineral compositional, and oxygen isotope properties.Based on the data compiled by Horstmann and Bischoff(2014) and the original descriptions referenced therein,these samples span nearly the same range of Fo values(olivine core Fo 78–96) as the main group ureilitecollection on the whole and ureilitic clasts in otherpolymict ureilites (Fig. 2). In addition, they show thesame variety of ureilite types (olivine-pigeonite, olivine-orthopyroxene, and augite-bearing), with characteristicfeatures of each of these types (e.g., AhS #44 is atypical olivine-pigeonite ureilite of Fo 79 and AhS #15is a typical augite-bearing ureilite of Fo 87). Over halfof these samples conform to the correlation of low-Capyroxene composition with Fo that is seen among maingroup ureilites (for the others, insufficient information isavailable to assess this). Studied sample sizes are mostlytoo small to determine whether pyroxene abundance iscorrelated with Fo as in main group ureilites. Oxygenisotope compositions of ureilitic samples of AhSconform to the established ureilite trend (Fig. 3), and,in keeping with their range of Fo, span the same rangeof D17O values as main group ureilites and ureiliticclasts in other polymict ureilites (Rumble et al. 2010).

Horstmann and Bischoff (2014) describe thetextures of coarse-grained AhS samples MS-156, MS-202, and MS-171 as unique. However, the unusualtexture of reduced olivine that they describe in MS-156and MS-202 is similar to that of highly reduced olivinein main group ureilite ALHA82130/82106 (Takeda et al.1989) and the mosaicized texture of MS-171 is simlar totextures in highly shocked areas of LAP 03587 andLAR 04315 (Warren and Rubin 2010). Thus, thesesamples are not unique compared to previously knownureilites.

The fine-grained AhS ureilite samples have receivedconsiderable attention for their apparently anomalousproperties, in particular their high porosity. However,as pointed out by Mikouchi et al. (2010) and Warrenand Rubin (2010), most of the features of these ureilitesare observed in highly shocked main group ureilites orhighly shocked areas of heterogeneously shocked maingroup ureilites (Goodrich et al. 2013). Some of thefeatures of the AhS samples, such as their porosity andthe structural states of their pyroxenes, have beenstudied in much greater detail (e.g., Herrin et al. 2010;Mikouchi et al. 2010; Zolensky et al. 2010) thananalogous features in main group ureilites, which maycontribute misleadingly to the impression that they arenew. And although the high abundance of these samplesin the AhS ureilite collection (approximately equal tothe coarse-grained samples) is notable, it seems likely


that similar highly shocked samples are more abundantamong main group ureilites than the literature suggests,because such samples tend to be ignored by researchersinterested mainly in primary petrogenesis. Furthermore,highly shocked ureilite clasts are often very abundant inother polymict ureilites (Goodrich et al. 2004; Ikedaet al. 2000).

The three metal+sulfide-rich AhS samples havesignificantly higher proportions of metal+sulfide thanobserved in any main group ureilite (Goodrich et al.2013), and their metal shows distinct zonation patternsand compositions (Horstmann et al. 2011; Ross et al.2011a, 2011b). Horstmann and Bischoff (2014)emphasize the unique textures of these samples.However, similar textures have been observed on asmall scale in highly shock-smelted pyroxenes in somemain group ureilites (Goodrich et al. 2013).Furthermore, metal-sulfide-rich ureilitic clasts, few ofwhich have been studied in any detail, have beenobserved in other polymict ureilites (e.g., Ikeda et al.2000). Thus, such samples may not be unusual as clastsin a polymict ureilite.

Finally, the single andesite sample from AhS, MS-MU-011, which has been shown by oxygen isotopes tobe related to ureilites (Bischoff et al. 2013), clearlyrepresents a rock type not known among main groupureilites. However, it is very similar in texture,mineralogy and mineral compositions, and chemicalcomposition to “the albitic lithology,” the mostabundant population of felspathic clasts in otherpolymict ureilites (Ikeda et al. 2000; Cohen et al. 2004;Goodrich and Wilson 2014). Thus, this sample toorepresents a previously known clast type in polymictureilites.

In summary, based on available data, the ureiliticsamples in AhS span nearly the same range of Fovalues as the main group ureilite collection and ureiliticclasts in other polymict ureilites (Fig. 2), and show thesame lithologic types. In petrographic details,particularly features attributable to shock, theycorrespond even more closely to the range and varietyof ureilitic materials seen in other polymict ureilites.Thus, we argue that these samples represent essentiallythe same selection of ureilitic material as that seenamong main group ureilites and in other polymictureilites, suggesting a common history for AhS and allother ureilites.

The Nonureilitic Almahata Sitta Samples

Approximately 30% of the AhS samples studied arenonureilitic, comprising various types of chondrites (seeHorstmann and Bischoff [2014] and references therein).Enstatite chondrites are the most abundant (27 of 35

samples) including a variety of EL and EH (petrologictypes 3, 3/4, 6, impact melt rocks and breccias), plusfive samples of ordinary chondrites (both H and L ofpetrologic types 4–6), and one Rumuruti-type and onecarbonaceous (CB) chondrite. We note that there issignificant overlap between this variety of materials andthe nonureilitic clasts in other polymict ureilites. Thelatter include fragments of ordinary chondrites (Jaquesand Fitzgerald 1982; Prinz et al. 1986, 1987, 1988; Ikedaet al. 2000; Kita et al. 2004), enstatite chondrites (Kitaet al. 2006; Downes et al. 2008) and R chondrites(Ikeda et al. 2000; Cohen et al. 2004; Downes et al.2008; Ross et al. 2010; Goodrich and Wilson 2014),although their relative abundances are poorly knowndue to the low overall abundance (<1%) of nonureiliticmaterials. Some apparent differences between thenonureilitic populations in Almahata Sitta and in theother polymict ureilites are (1) materials like the CB inAhS (Bischoff et al. 2012) have not been reported inother polymict ureilites, (2) materials like the “darkclasts” (CC-matrix-like materials) that are common inother polymict ureilites (Ikeda et al. 2003; Goodrichet al. 2004) have not been found among the AhSsamples, and (3) foreign achondritic materials like theangrite clasts in other polymict ureilites (Prinz et al.1986; Goodrich et al. 2004; Kita et al. 2004) have notbeen found among the AhS samples. It is not clear howsignificant these differences are, considering that onlyapproximately 15% of the total AhS sample collectionhas been studied in detail, and that the nonureiliticmaterials in other polymict ureilites have not beenthoroughly characterized. What is clear, however, isthat both populations are exceptional in the greatvariety of meteorite types they contain. In contrast, inmost other meteoritic breccias the foreign clastpopulation is dominated by CC matrix-like dark clasts(Bischoff et al. 2006). This important similarity suggestsa common history for AhS and other polymict ureilites.

Is Almahata Sitta an Anomalous Polymict Ureilite?

We have shown in the previous two subsectionsthat the variety of types of materials among theAlmahata Sitta samples that have been studied is verysimilar to that in other polymict ureilites (in contrast tothe view of Horstmann and Bischoff [2014] that thereare significant differences). Based on this observation,Almahata Sitta (the collection of samples, though notany of the individual samples) could be classified as atypical (nonanomalous) polymict ureilite. In otherrespects, however, Almahata Sitta is clearly anomalous.First, as emphasized by Bischoff et al. (2010b) andHorstmann and Bischoff (2014), AhS has a significantlyhigher fraction of nonureilitic materials than other


polymict ureilites (approximately 30% versus ≤1%).Second, if we regard the individual Almahata Sittasamples as clasts in a breccia, then there is a significantdifference in clast size between Almahata Sitta andother polymict ureilites. The AhS ureilitic samples rangeroughly from 0.5 to 6 cm in size (estimates based onmasses reported by Shaddad et al. 2010). In contrast, inother polymict ureilites, main group ureilite clasts rangeonly rarely up to the lower end of that range; most areon the order of 109 smaller (Goodrich et al. [2004] andreferences therein). Finally, it is worth underscoring themost obvious difference between AhS and otherpolymict ureilites, which is the circumstance that itsclasts landed on Earth as individual fragments ratherthan cemented together. In the next section we addressthe key question of what held these clasts together inasteroid 2008 TC3.

ASTEROID 2008 TC3—AN ANOMALOUS

POLYMICT UREILITE?

In this section we use published observations of2008 TC3 and the fall of Almahata Sitta to evaluate thenature of the asteroid and whether the meteorite is arepresentative sample of the asteroid. AlthoughAlmahata Sitta (a collection of monolithologic samples)may be anomalous as a polymict ureilite, asteroid 2008TC3 need not be.

Studies of the meteorite strewn field suggest that thetotal fallen mass was approximately 40 kg (Shaddadet al. 2010). Combining this with estimates of the pre-atmospheric mass of the asteroid, which range fromapproximately 42 to 83 t (Borovi�cka and Charv�at 2009;Jenniskens et al. 2009; Welten et al. 2010), we calculatethat the meteorites on the ground account for ≤0.1% ofthe entry mass. In other words, ≥99.9% of the mass ofthe asteroid was lost in the atmosphere, probably as thematerial comprising the observed dust clouds. Based onMeteosat observations of the atmospheric entry of 2008TC3, Borovi�cka and Charv�at (2009) estimated that thenonsublimated dust particles in these clouds weremicron-sized, suggesting that the bulk of the asteroidwas extremely fine-grained. In addition, estimates of thebulk density of the asteroid from abundances ofcosmogenic radionuclides in the meteorites are in therange 1.2–1.6 g cm�3, which is significantly lower thandensities of any of the meteorites and indicates a bulkporosity of 50 � 7% (Welten et al. 2010). Thus, we inferthat 2008 TC3 consisted mostly of highly porous, weaklyconsolidated matrix material, with only a small fractionof isolated, larger sized fragments of denser, well-consolidated rocks that became the fallen meteorites.The composition of this matrix material is unknown, butthe asteroid reflectance spectrum provides a clue.

The reflectance spectrum of 2008 TC3 is mostclosely matched by the spectra of F-class asteroids(Jenniskens et al. 2009, 2010) in the Tholen asteroidtaxonomy (Tholen and Barucci 1989). F-class asteroidsbelong to the broader C group of asteroids that arethought to be identified with carbonaceous chondrites,and Cloutis and Hudon (2004) concluded that the spectraof ureilites are most similar to those of carbonaceouschondrites and C-type asteroids. The F-class has notbeen associated with any specific meteorite type. Alaboratory study of reflectance spectra of 11 AhS samples(of which 10 were ureilitic and one was an ordinarychondrite) found that spectral mixtures dominated bythe ureilite samples could reproduce the main features ofthe asteroid spectrum (Hiroi et al. 2010). This result,albeit limited (as the study did not include the fulldiversity of AhS sample types and the asteroid spectrumdid not extend redward of 1 lm), suggests that ureiliticmaterial dominated at least the surface of asteroid 2008TC3.

If the ≥99.9% of the mass of 2008 TC3 that waslost in Earth’s atmosphere consisted mostly of ureiliticmaterial, then (in contrast to the interpretation ofBischoff et al. 2010b) the proportions of variousmaterials comprising 2008 TC3 were not anomalouscompared with those in other polymict ureilites; i.e., theabundance of nonureilitic material must have been<0.1%, much lower than in AhS. In that case, theprincipal difference between 2008 TC3 and previouslyknown polymict ureilites is in degree of consolidation,i.e., bulk strength. In other polymict ureilites, most ofthe main group ureilite-like material is in approximately0.1 to 1 mm sized clasts in a coherent fragmentaltexture (i.e., sufficiently coherent to survive passagethrough Earth’s atmosphere), whereas 2008 TC3 was soweakly consolidated and porous that it completelydisintegrated at a high altitude in the atmosphere. Somebulk strength was necessary in 2008 TC3, because it wasobserved to be spinning with a 99 s rotation period(Scheirich et al. 2010). Hence, it could not have been arubble pile (held together only by gravitational forces;Richardson et al. 2002). However, it must have beenonly weakly lithified, suggesting that it had beensubjected to only a minimal amount of heating and/orpressure.

The extent of difference in grain size between 2008TC3 and other polymict ureilites is unclear. As discussedabove, the AhS ureilitic samples are on the order of109 larger than most clasts in other polymict ureilites.However, this could be due to preferential survival inthe atmosphere, compounded by preferential collection,of only the largest clasts, and the average clast size inthe asteroid could have been significantly smaller.Furthermore, the grain size of the lost matrix material


is unknown, with the micron-size estimate fromMeteosat observations (Borovi�cka and Charv�at 2009)being only a lower limit. Hiroi et al. (2010) found thatspectral mixes of 0.1–0.5 mm sized powders of the AhSureilite provided better fits to the 2008 TC3 spectrumthan those of fine (<0.1 mm sized) powders, suggestingthat grain sizes similar to the majority of clasts in otherpolymict ureilites dominated the surface of the asteroid.If this is the case, then 2008 TC3 may have been similarto other polymict ureilites in grain size distribution.

In summary, although the Almahata Sitta polymictureilite is anomalous in its high abundance ofnonureilitic material and large clast size, these featuresare at least partly a result of selection effects duringatmospheric passage and collection, and are notrepresentative of the asteroid. We suggest that 2008 TC3

was a polymict ureilite that differed from other knownpolymict ureilites only in being less consolidated.

Combining the results of this and the previoussection, we emphasize that 2008 TC3 samples essentiallythe same selection of materials as other polymictureilites and the main group ureilite collection,suggesting that all known ureilitic materials share acommon history. Based on these observations, wehypothesize that the immediate parent of 2008 TC3 wasalso the immediate parent of all main group andpolymict ureilitic material that has been sampled inrecent times by meteorites on Earth. In the followingfive sections, we infer the history of this material fromaccretion of the ureilite parent body to arrival on Earth(Fig. 1).

ACCRETION OF THE UREILITE PARENT BODY:

WHEN AND WHERE?

All ureilitic materials in our meteorite collection arethought to come from a single original parent body (theUPB), based on a common and distinctive thermalhistory of rapid quenching from high temperatures ofapproximately 1050–1100 °C (Goodrich et al. 2004;Downes et al. 2008; Herrin et al. 2010). In this section,we constrain the time and location of accretion of thisbody. Where the UPB formed is critical to theoverarching goal of this work to explain how thedisparate meteoritic materials in AhS and otherpolymict ureilites came to reside in the same rock.

When Did the UPB Accrete?

The timing of accretion of the UPB can beconstrained from thermal modeling. The presence ofexcess 26Mg in feldspathic clasts in polymict ureilites(Kita et al. 2003; Goodrich et al. 2010b) providesevidence for live 26Al at the time of UPB formation.

This isotope, together with smaller contributions fromother short-lived radionuclides, has been used by severalauthors as the major heat source to model the thermalhistories of asteroids in general. Thermal modeling ofthis type for the UPB (Wilson et al. 2008) led to theconclusion that the high equilibration temperaturesrecorded in ureilites (1200–1300 °C) could only havebeen achieved if the UPB accreted within approximately0.54–0.57 Ma after CAI formation. Here, we revisethese calculations to make them more general,eliminating the model-dependent assumption thatsmelting (reduction of FeO by C, with consequentproduction of CO gas) occurred during melting on theUPB. We also use a wider range of initial temperatures,specifically 150–250 K instead of the earlier 240–260 K.The modeling of Wilson et al. (2008) assumed a size of125 km radius for the UPB, constrained by the smeltingmodel. The new modeling is applicable to asteroid sizesranging from approximately 100 km to 250 km radius,encompassing the range of UPB size estimates in bothsmelting and nonsmelting models (Michel et al. 2013).

The controlling parameters in the thermal modelare the time of accretion, the initial temperature and icecontent (varied from 0 to 25 mass%), and the bulkcomposition (insofar as this controls the 26Al content).For a given set of parameters, the model predicts thetime of onset of silicate melting, the peak temperaturereached as a function of depth in the asteroid, and thetime at each depth at which that peak temperature isreached. We can therefore determine a window ofaccretion times which permit the UPB to achieve thepeak temperatures that are recorded in ureilites. Peaktemperatures also depend on physical and chemicalprocesses during melt migration. Wilson et al. (2008)showed that on asteroidal-sized bodies, magmaseparation from melting source rocks at depth shouldalways be a very efficient process, due to the formationof a hierarchical structure of veins and dikes extractingmelt from grain boundaries, and found that any givenparcel of melt should move from mantle depths to thebase of the crust in a time of order 1 month. Incomputing the temperature structure of the UPB, theytracked this transfer of melt, including the consequencesof the 26Al heat source being preferentially partitionedinto the melt along with the dominant 27Al isotope. Wehave retained this aspect of the model.

Rapid melt migration was enhanced in the Wilsonet al. (2008) model by the increased buoyancy generatedby the formation of CO gas during smelting. However,removing the assumption of smelting only decreases theefficiency of melt migration by a factor of order three(Wilson and Goodrich 2012), so that melt transfer timesare still less than 1 year. To determine the effect of thison the thermal history of the UPB, it is necessary to


compare this time scale for heat advection with the oneinvolved in the redistribution of heat by conduction. Itcan be shown (Wilson et al. 2008; Warren 2011a) that,until significant melt migration begins, asteroids heatedby a source as powerful as 26Al will have an outer shellof thickness approximately 5–10 km in which thetemperature decreases rapidly to the surface valuecontrolled by radiative heat transfer to space. Beneaththis shell there will be a much smaller temperaturegradient, decreasing with increasing depth, having alength scale of order 10 km, and evolving on a timescale of order 1 Ma. Thermal conduction betweenregions with these separations and on this time scale isthe process that ultimately determines the peaktemperature occurring at any depth in the asteroid.Thus, changing the speed at which melt migrationoccurs on the much shorter approximately 1 year timescale has a negligible effect on the large-scale thermalhistory.

Removing the assumption of smelting also changesthe silicate melting sequence. However, the onlysignificant effect is on olivine/pyroxene ratios (Goodrichet al. 2007). As there is little difference in the specificheats of these two minerals, changing their relativeproportions has a negligible effect on the thermalhistory. The abundance of plagioclase in the startingmaterial, and the temperature at which plagioclase isexhausted during melting, which are the only aspects ofthe silicate melting sequence that significantly affect thethermal history, are essentially unchanged.

The results of the new modeling for UPB accretiontime are shown in Fig. 4. They imply a window ofapproximately 0.54–0.62 Ma after CAI, only slightlydifferent from the previous result of Wilson et al.(2008). Silicate melting still begins at approximately1 Ma after CAI. These results for the thermal history ofthe UPB (foreseen by Scott 2006) are consistent withthe <1 Ma accretion time obtained by Bizzarro et al.(2005), and the 0.75 Ma accretion time obtained byHevey and Sanders (2006), based on similar thermalmodeling (i.e., short-lived radionuclides as the heatsource) for differentiated planetesimals in the solarsystem, and also with evidence from the 182Hf-182Wsystem for very early differentiation (<1 Ma after CAI)of the parent asteroids of many iron meteorites (Kleineet al. 2005; Markowski et al. 2006).

Early formation of the UPB (within 1 Ma afterCAI) is also consistent with models for accretion ofthe asteroids from smaller planetesimals. This accretiontime scale is much shorter than that cited by Grimmand McSween (1993), who did not perform actualsimulations of accretion, but used a simplified analyticmodel for the growth rate. Moreover, they started withthe assumption that accretion in the outer belt was too

slow to allow melting by 26Al, and chose parametersdesigned to yield that result. More recent numericalsimulations show that accretion of asteroids from smallplanetesimals can occur on shorter timescales. Ageneral feature of such models (e.g., Weidenschilling2011) is an early phase of rapid runaway growth,producing an “oligarchy” of roughly Moon- to Mars-sized planetary embryos. Perturbations by the largeembryos increase velocities among the residualpopulation of smaller bodies, slowing their accretion.Thus, most of the mass of an asteroid, regardless ofsize, is accreted on a time scale generally comparableto the formation time of the embryos (0.1–1 Ma).Weidenschilling (2011) modeled this growth process fora range of assumed initial planetesimal sizes and foundthat smaller initial planetesimals yield shorter embryogrowth times. The best match to the observed sizedistribution of the asteroids occurs for simulationsstarting with sub-km planetesimals, which yield growthtimescales of order 0.1 Ma. It should be kept in mindthat this is the interval after planetesimal formation,which occurred sometime after formation of CAIs. Asplanetesimals must have formed in significantly lesstime than the half-life of 26Al (approximately 0.7 Ma),accretion of the UPB could have occurred in <0.8 Maafter CAI, consistent with our estimate from thermalmodeling.

Fig. 4. Results from new thermal modeling for accretion times(in Ma after CAI formation time) that allow the UPB to reachpeak equilibration temperatures of 1200–1300 °C, as afunction of the mass fraction of water ice in the asteroid atthe time of accretion, for three values of the initialtemperature of the body. Results are for an asteroid of100 km radius but would not differ significantly for asteroidsup to approximately 250 km radius.


Thus, the UPB may have accreted and begun todifferentiate before the formation of most chondrules,currently dated at 1–3 Ma after CAI (Kleine et al. 2005;Connelly et al. 2008; Kita and Ushikubo 2012; Kitaet al. 2013), and therefore before accretion of mostchondrite parent bodies. As discussed in the nextsection, this scenario could have significant implicationsfor the chemical and isotopic compositions of ureiliteprecursor materials, and therefore for where the UPBaccreted.

Where Did the UPB Accrete?

To infer the location of accretion of the UPB (interms of heliocentric distance) we first discussconstraints on the compositions of ureilite precursormaterials compared with the compositions of knownchondrite types. We then review two possible sources ofinformation about where the various chondrite typesformed (1) evidence from asteroid spectroscopy for thecurrent variation in asteroid (and linked meteorite)types in the asteroid belt and (2) evidence fromchondrites and their components for variations inchemical and isotopic composition in the nebula. Wethen use these two types of information to evaluatewhat the compositions of ureilite precursor materials tellus about where the UPB accreted.

Constraints on the Chemical and IsotopicCompositions of Ureilite Precursor Materials

Ureilites are known for their carbon-richcompositions, which have been suggested to indicate agenetic link to carbonaceous chondrites (e.g., Mason1962; Mueller 1969; Vdovykin 1970). Ureilite carboncontents (Fig. 5) are typically in the range 1–7 wt% andaverage approximately 3.2 wt% (Hudon et al. 2004;Warren 2011b). To infer the carbon contents of ureiliteprecursor materials, we need to take into account theigneous processing they have experienced. The loss of15–30% silicate melt that they experienced could haveresulted in either a decrease or an increase in carboncontents in the residues, depending on whether melt wasremoved in batch (Warren and Kallemeyn 1992) orfractional (Wilson et al. 2008; Wilson and Goodrich2012) processes. In batch melting, during whichsignificant volumes of melt would have formed beforeseparating from the residues, graphite (the high-temperature polymorph of carbon which would haveformed from low-T forms during progressive heating)would be expected to be incorporated into the meltfraction due to its low density, and thus physicallyextracted along with the melt. In this case, the carboncontents of the residues (i.e., ureilites) would be lowerthan those of their precursor materials. However,

Wilson et al. (2008) argued from physical modeling thatfractional melt extraction in a network of thin veins wasmore likely on the UPB. In this case, the graphitecrystals in ureilites, which are typically mm-sized, wereprobably too large to have been extracted because muchof the melt network involved veins no wider thanapproximately 30 lm. This inference is supported bythe absence of carbon in feldspathic clasts (presumed tohave crystallized from UPB melts) in polymict ureilites(Cohen et al. 2004). Thus, it is likely that carboncontents of the ureilites are higher than those of theirprecursor materials by a factor of approximately 1.25(assuming 20% melt loss). On the other hand, ureilitesmust have lost some fraction of carbon (estimated atapproximately 0.5 wt% by Warren 2011b) in the formof CO gas generated during late reduction (Mittlefehldtet al. 1998). Considering these two effects together,ureilite precursor materials would have hadapproximately 0.5–5.5 wt% carbon. These carboncontents are similar to those of the most carbon-richchondrites, which are the CI, CM, CR and ungroupedTagish Lake carbonaceous chondrites (Fig. 5). Carbonisotopes are consistent with a link to such carbonaceouschondrites. The d13C compositions of ureilites are in therange approximately 0 to �11 & (Grady et al. 1985a,

0 1 2 3 4 5 6 7 8

CB

R

carbon content (wt.%)

CKCOCV

CM

CR

CITagish Lake

Ungrouped ChondritesOC

EC

Ureilites

Fig. 5. Carbon contents of bulk ureilites compared with thoseof various types of chondrites. Data from Wiik (1972) andsources therein, Gibson (1976), Kerridge (1985) and sourcestherein, Grady et al. (1986), Jarosewich (1990), Grady andPillinger (1986, 1993), Pearson et al. (2006), and Bischoff et al.(2011). Dashed line with arrows for ureilites shows range ofureilite carbon contents from Hudon et al. (2004). CI, CM,CR, CV, CO, CK, and CB = various groups belonging tothe carbonaceous chondrite class. OC = ordinary chondriteclass. EC = enstatite chondrite class. R = Rumuruti groupchondrites.


1985b; Grady and Pillinger 1987, 1993; Hudon et al.2004), similar to those of CI and CM and distinct fromthose of all other chondrite groups (see compilation inKrot et al. 2013).

On the other hand, lithophile element (magnesiumsilicate forming) compositions of ureilites indicate thatureilite precursor materials were similar to OC or EC.For example, the chondrite classes are distinguishedby significant differences in Si/Mg ratio, increasing inthe order CC?OC&R?EC (Fig. 6). These differenceshave important effects on the partial melting processand the mineralogy (particularly pyroxene/olivineratios) of resulting residues. Goodrich (1999) modeledpartial melting of a range of chondritic compositionsin the system Ol-Plag-Opx-Wo (Longhi 1991), andconcluded that ureilite precursor materials must havehad Si/Mg ratios at least as high as those of OC, i.e.,higher than those of CC, in order to yield residueswith pyroxene/olivine ratios as high as those inureilites. Warren (2011b) reached a similar conclusionbased on mass balance modeling of partial melting ofpossible ureilite precursor materials. The chondriteclasses also differ in Mg/Mn ratios, with most CCgroups having significantly higher values than OC andR-chondrites (Fig. 6). Ureilites have near-constant Mg/Mn ratios between those of CM and CO. However,taking into account the effect of approximately 15–30% partial melting, which would have resulted in anincrease in these ratios (Goodrich and Delaney 2000),the Mg/Mn ratios of ureilite precursors would havebeen similar to those of OC and R-chondrites(Fig. 6).

Furthermore, it has recently become apparent thatCr, Ti, and Ni isotope compositions of ureilites suggestderivation from OC or EC-like materials (Warren2011b). A compilation of high-precision isotope data forthese elements (Shukolyukov and Lugmair 2006;Trinquier et al. 2007, 2009; Leya et al. 2008; Quitt�eet al. 2010; Dauphas et al. 2010) shows that planetarymaterials fall broadly into two groups (Warren 2011b).Carbonaceous chondrites of all types form one group,which is characterized by positive e54Cr, e62Ni, ande50Ti, i.e., enrichments in the neutron-rich isotopes. Theordinary chondrites, enstatite chondrites, and variousdifferentiated meteorites (angrites, Martian meteorites,main group pallasites, HEDs, mesosiderites, and someirons) form the second group, characterized by negativee54Cr, e62Ni, and e50Ti, i.e., deficits of the neutron-richisotopes. Although only a few ureilites have beenanalyzed for these isotopes, their compositions (whichwould not have been changed during igneousprocessing) fall clearly into the second group and thussuggest that ureilite precursors were not CC-like(Warren 2011b).

At the same time, the oxygen isotope compositionsof ureilites, which plot along the slope approximately 1CCAM line defined by the components of CVchondrites (Fig. 3), remain a strong indication thatureilites are genetically related to carbonaceouschondrites (Clayton and Mayeda 1988). The slopeapproximately 1 ureilite trend is remarkable, because allother groups of differentiated solar system materialsplot along slope approximately 1/2 trends (e.g., theEarth-Moon system, Martian meteorites, HEDs,brachinites, acapulcoites, angrites, aubrites), reflectingboth homogenization of primordial oxygen isotope(D17O) variation and mass-dependent fractionationduring igneous processing. However, Warren andKallemeyn (1992) and Scott et al. (1993) argued thatprimordial oxygen isotope (D17O) heterogeneity mightbe preserved if ureilites were residues rather thancumulates. Wilson et al. (2008) showed that in rapid,fractional melt extraction on the UPB, migrating meltswould not have had time to equilibrate oxygen isotopeswith the residual matrix (i.e., ureilites) through whichthey passed. Thus, ureilite oxygen isotope compositions(unlike those of other achondrites) may directly preservethose of their precursor materials and it is reasonable toconclude that those precursors were related in some wayto carbonaceous chondrites. Although derivation fromknown CCs is ruled out by the lithophile element andCr, Ti, and Ni isotopic properties of ureilites discussedabove, the oxygen isotope properties of ureilites remaincritical in showing that the UPB experienced some

1.4 1.5 1.6 1.7 1.8 1.9

60

70

80

90

100

110

120

130

140

R

CM

CO

EC

OC

CI H L LL R EC CV CM CO

wt.

MgO

/MnO

wt. SiO2/MgO

CV

ureiliteprecursors

Fig. 6. Variation in Si/Mg and Mg/Mn ratios among differentchondrite groups. Taking into account the partial meltingexperienced by ureilites, ureilite precursor materials areinferred to have had Si/Mg and Mg/Mn ratios similar to thoseof OC or R-chondrites, rather than CC. Data from Jarosewich(1990) and Bischoff et al. (2011).


process in common with CC but not OC or EC. Wesuggest that this process was aqueous alteration.

Oxygen isotope analyses of the components ofprimitive (type 3) carbonaceous chondrites (Young andRussell 1998; Young et al. 1999; Ushikubo et al. 2012;Tenner et al. 2011a, 2011b) show slope approximately 1arrays that are displaced to lower d18O relative to theaqueously altered CC, i.e., to the left of the CCAM lineon a 3-oxygen isotope diagram (Fig. 3). These studiessuggest that primordial oxygen isotope reservoirs ofsolids in the solar system formed a single slopeapproximately 1 array, which passed near OC and R,and to the left of the CCAM line (Fig. 3). Theendmembers of this array are proposed to have been16O-poor water ice and 16O-rich solid silicates. Therehave been two different proposals for the location ofthis line—the YR line of Young and Russell (1998) andthe PCM (primitive chondrule mineral) line ofUshikubo et al. (2012), but both sets of authors arguethat water-rock interactions, i.e., aqueous alteration onparent asteroids, was responsible for shifting the isotopecompositions of most CC toward the CCAM line. Ifthis hypothesis is correct, it points strongly towardaccretion of water ice and consequent pre-igneousaqueous alteration on the ureilite parent body (Kitaet al. 2014), and tells us that the UPB accreted beyondthe ice line.

Variation in Asteroid Spectral Types (and LinkedMeteorite Types) with Heliocentric Distance

The two main groups of asteroids are S- andC-types, which dominate the inner (semimajor axisa < 2.8 AU) and outer (a > 2.8 AU) asteroid belt,respectively (Gradie and Tedesco 1982; Moth’e-Dinizet al. 2003; DeMeo and Carry 2013). S-type asteroidscan be linked to the ordinary chondrite meteoritesthrough their spectra (Binzel et al. 1996). This issupported by analytical results linking samples of theS-type asteroid Itokawa (returned by the Hayabusaspacecraft) to the LL group of ordinary chondrites(Nakamura et al. 2011). C-type asteroids are generallyassociated with carbonaceous chondrites because theyboth have low albedos and relatively featureless spectra.Higher resolution spectra show that the 0.7 lm featureseen in CM carbonaceous chondrites and indicative ofFe3+ is present in roughly half of C-type asteroids (e.g.,Burbine 1998; Burbine et al. 2002). Carbonaceouschondrites, particularly the CI and CM classes, containsignificantly more water than ordinary chondrites,suggesting that the outer asteroid belt has higherabundances of the most volatile elements comparedwith the inner belt. This is supported by the observationof cometary activity on some asteroids in the outermain belt (Hsieh and Jewitt 2006), the direct detection

of water ice on the surface of C-type asteroid Themis(a ~ 3.1 AU) (Campins et al. 2010; Rivkin and Emery2010), and recent observations of water vapor ventingfrom G-type (closely related to C-type) asteroid Ceres ata = 2.77 AU (K€uppers et al. 2014).

In addition to the S- and C-type asteroids, there areseveral other significant classes. The far outer asteroidbelt, especially the Hilda region around 4AU and theTrojan asteroids of Jupiter, contain primarily P-andD-type asteroids, which have low albedo and relativelyfeatureless, red spectra. The very innermost asteroidbelt, especially the high-inclination Hungaria region,contains many E-type asteroids, which have beenspectrally linked to the aubrite (enstatite achondrite)meteorites by their spectra and high albedos (Gaffeyet al. 1992). The spectrum and albedo of asteroid 21Lutetia (with a ~ 2.4 AU), which was observed by theRosetta spacecraft, are consistent with those of theenstatite chondrites (Vernazza et al. 2009, 2011).Although there is not yet a clear link between theenstatite chondrites and an entire class of asteroids,aubrites are likely derived from enstatite chondrite-typeprecursors (Keil 2010), which suggests that theinnermost part of the asteroid belt was once dominatedby EC-type material.

The origin of the radial variations in the asteroidbelt is still a matter of debate. It has generally beenthought that the observed compositional variationrecords the primordial chemical and thermal state of thesolar nebula throughout the 2–4 AU region (Gradie andTedesco 1982), although some early work suggested thatmaterial had been implanted into the asteroid belt fromthe terrestrial planet zone (Wetherill 1977). More recentwork supports the idea of implantation, showing thatobjects from the terrestrial zone (possibly the parentbodies of iron meteorites and enstatite meteorites) couldhave been scattered outward into the belt (Bottke et al.2006). In addition, more volatile-rich bodies might havebeen scattered inward and implanted in the outerasteroid belt during early-stage giant planet migration(e.g., Walsh et al. 2011) and/or during later dynamicalinstabilities (e.g., Levison et al. 2009). However, even ifthe different asteroid classes originally formed over awider range of radial distances than they presentlyencompass, their arrangement is probably still indicativeof the order (if not the absolute locations) of theirvariation as a function of heliocentric distance in theearly nebula.

Variations in Chemical and Isotopic Compositions ofChondrites

Variations in chemical and isotopic compositionsamong the approximately 15 chondrite groups and theircomponents are enormously complex, and how these


variations might be related to location of accretion isunclear (Brearley and Jones 1998; Scott and Krot 2013).In general, chondrites are comprised of four types ofpetrographic components (1) CAIs (Ca-Al-richinclusions) and AOAs (amoeboid olivine aggregates); (2)chondrules; (3) matrix; and (4) Fe,Ni metal (Brearleyand Jones 1998; Scott and Krot 2013). The chondriteclasses are distinguished by significant differences in theproportions of these components. Carbonaceouschondrites are rich in CAIs, AOAs, and matrix, andpoor in chondrules and metal, compared with both OCand EC. Some differences in bulk chemical compositionamong the chondrite classes correspond to differences inproportions of these components. For example, CC areenriched in refractory elements compared with OC andEC, consistent with CAI and AOAs being the mainrepository of these elements in chondrites. Likewise,matrix is the main repository of highly volatile elements(such as carbon) and organics, accounting for higherabundances of these in CC. It is also rich in FeO,which, along with low abundance of metal, accounts forthe higher oxidation states of CC compared with OCand EC. However, differences in magnesium silicate-forming lithophile elements (e.g., Si/Mg ratio) andmoderately volatile elements (e.g., Mn) are manifestedlargely in chondrules (Jones 2012; Scott and Krot 2013).Each chondrite group, and each class, appears to havesampled a unique, localized reservoir of chondrules(Jones 2012). And although CAIs in all chondrites maybe derived from a common reservoir (Krot et al. 2009),the proportion of them in each group (or at least class)is unique. Moreover, it is not possible to interpret mostof the variations between groups, or even classes, interms of systematic chemical trends (Brearley and Jones1998). In particular, if the mapping of asteroid types tochondrite classes described above is correct, then thechondrite classes vary in the order EC?OC?CC withheliocentric distance. Increasing oxidation state (Rubinand Wasson 1995) and increasing abundance of highlyvolatile elements might be explained by this correlation,particularly if CI, with the highest concentrations ofcarbonaceous, volatile-rich material, is considered theoutermost group (Scott and Krot 2013). However, mostother properties, including abundances of refractory andmoderately volatile elements, cannot. Thus, exceptpossibly for abundances of highly volatile components,chondrites do not show any general geochemical trendsthat could be used to infer the location of UPBaccretion.

Turning to isotopic properties, the origin of oxygenisotope anomalies in solar system materials is stilluncertain (Yurimoto et al. 2007). Of the variousmechanisms that have been proposed, the scenario ofphotochemical self-shielding of CO in the outer

protoplanetary disk (Lyons and Young 2005) is thoughtto be one of the most viable (Yurimoto et al. 2007).This model suggests that in the inner protoplanetarydisk, solid silicates (initially with the 16O-richcomposition inferred for the Sun, D17OSMOW = �50 &)reacted with 17,18O-rich CO+H2O gas generated by driftof water ice inwards across the ice line, and evolved tothe relatively small range of compositions seen in OCand EC and most planetary materials. In contrast, inthe outer disk, solid silicates retained their 16O-richcompositions and water ice retained its 17,18O-richcomposition (D17OSMOW >100&), due to lowtemperatures. The key feature of this model that isrelevant to ureilites is that it implies that bodiesaccreted in the outer disk would contain both 16O-richsilicates and 17,18O-rich ice, giving them the potential todevelop heterogeneous negative D17O values like thoseof ureilites and CC during parent body processing,whereas in the inner disk bodies would accrete withmuch more homogeneous 17,18O-rich compositions.

The origin of the 54Cr, 50Ti, and 62Niheterogeneities in the solar system, and the markedbimodality between CC and nearly all other meteoritetypes in terms of these compositions, is not known.Trinquier et al. (2009) suggested that the protosolarmolecular cloud was originally well-mixed and CC-likein these isotopes, and the observed heterogeneities aroselater due to selective loss of presolar silicates by thermalprocessing in the inner solar system. Leya et al. (2008)suggested that CC acquired their distinct mix ofnucleosynthetic components as a result of formationboth further out in the protoplanetary disk and laterthan other meteorites. Dauphas et al. (2010) identifiedthe carrier of 54Cr anomalies in Orgueil (CI) asnanoparticles (most likely spinel), and suggested thatplanetary materials incorporated different amounts ofthese particles due to late injection by a supernova. Thevery early accretion time (<1 Ma after CAI) of the UPBsupports the hypothesis that the 54Cr and 50Ti excessesin CC were a late addition to the solar system. If, as weinfer from oxygen isotopes, the UPB accreted in thesame general region as CC (i.e., beyond the ice line) butat an earlier time, then the absence of 54Cr and 50Tiexcesses in ureilites suggests that these isotopes wereintroduced to the CC formation region only later.

Synthesis: Where Did the UPB Form?In summary, key lithophile element properties and

Cr, Ti, and Ni isotope compositions of ureilites suggestOC- or EC-like precursors, whereas carbon contents,carbon isotopes and oxygen isotopes suggest CC-likeprecursors. This appears to present a conundrum bycomparison to the evidence from asteroid spectroscopy,which suggests that EC and OC have dominated the


inner asteroid belt and CC the outer belt since early insolar system history. However, based on the evidencefrom chondrites, we suggest that the carbon and oxygenisotope constraints are a more reliable indicator oflocation of formation than the lithophile elements orthe Cr, Ti and Ni isotopes, for the following reasons.

If the UPB accreted before the formation of knownchondrules, as our thermal modeling indicates, thenthere is a serious problem with using comparisons tochondrites to assess where the UPB accreted. Theproperties of chondrules (which are a major componentof most chondrites) are a complex function not only ofthe compositions of chondrule precursor materials, butalso of the conditions and mechanisms of chondruleformation (Jones 2012; Scott and Krot 2013). Localchemical fractionations that occurred during chondruleformation may well have obscured pre-existing chemicalvariations with heliocentric distance. Lithophile elementproperties such as Si/Mg ratio and abundances ofmoderately volatile elements, which are predominantlymanifested in chondrules in chondrites, could have beensignificantly altered by chondrule-forming processes.Thus, comparing these properties of ureilite precursormaterials with those of chondrites may provide littleinformation about conditions before chondruleformation.

On the other hand, because the carbon inchondrites is associated with matrix, the high carboncontents of ureilite precursor materials can be taken asa reliable indicator of accretion in the outer part of thedisk (like CI, CM, and CR). Even more compellingly, inthe context of recent hypotheses for the origin andevolution of oxygen isotopes in the solar system, theCC-like oxygen isotope compositions of ureilites implythat the UPB accreted with a significant amount ofwater ice, and therefore beyond the ice line. Thus, weargue that the UPB accreted in the outer asteroid belt(or outer disk?), in the same general region as CC butat an earlier time (before the neutron-rich isotopes ofCr, Ti, and Ni were introduced to the solar system).

CATASTROPHIC DISRUPTION OF THE UPB AND

FORMATION OF A FAMILY OF OFFSPRING

In this section, we review evidence and argumentsin the literature for early catastrophic disruption of theUPB. Then we present results of new modeling of theformation of a family of offspring in the aftermath ofthis breakup. These events are a critical step leading tothe ultimate formation of 2008 TC3 and Almahata Sittain our model.

All main group ureilites show petrologic evidence ofextremely rapid cooling (approximately 1–20 °C h�1),accompanied by a drop in pressure, through the range

approximately 1050–600 °C (Mittlefehldt et al. 1998;Herrin et al. 2010). This distinctive thermal history (notseen in other groups of achondrites) is commonlyinterpreted to result from excavation during a majorimpact or catastrophic disruption of the UPB while itwas still hot (Takeda 1987; Warren and Kallemeyn1992; Goodrich et al. 2004; Downes et al. 2008; Herrinet al. 2010). An age of approximately 5 Ma after CAIfrom 26Al-26Mg and 53Mn-53Cr age dating of clasts inpolymict ureilites has been suggested to record thisevent (Goodrich et al. 2010b). We hypothesize that thisearly excavation was a key factor in the ultimateformation of 2008 TC3 (Fig. 1c).

Ureilites have various young cosmic-ray exposureages (approximately 0.5–47 Ma; Eugster 2003; Rai et al.2003), which indicate that they were not immediatelydelivered to Earth by this early excavation event. Thus,the excavated material must have been storedsomewhere. This conclusion is supported by theexistence of polymict ureilites (the UPB is unlikely tohave a well-developed regolith 5 Ma after CAI), and theobservation that the clasts they contain show the samedistribution of Fo values (and same thermal history) asmain group ureilites (Fig. 2). In particular, this suggeststhat the regolith represented by polymict ureilitesdeveloped on the same selection of primary ureiliticmaterials as our collection as a whole. There are,however, two possibilities for where this could haveoccurred. If the excavation was caused by asubcatastrophic impact, ureilitic material could havereaccreted back onto the parent body, forming an outerrubbly layer of rearranged materials (Warren andKallemeyn 1989) that subsequently evolved into a finergrained regolith. Alternatively, if the excavation wascaused by complete parent body breakup, then the bodyon which the regolith developed is likely to have been asecond-generation “daughter” body formed in theaftermath of the collision (Michel et al. 2001). Thesealternatives have been discussed in the literature, withcompelling arguments favoring the offspring body(Goodrich et al. 2004; Downes et al. 2008; Herrin et al.2010). This interpretation has the important implicationthat all ureilitic material in our collection may bederived from a single ureilite daughter body (UDB),rather than directly from the original UPB.

Catastrophic disruptions of large asteroids (largeenough that they have significant gravity) involve nearlycomplete shattering of the target into sub-km-sizedpieces, initial dispersal of the resulting fragments, andthen gravitational reaccumulation of subsets of thosefragments to form a family of offspring (Michel et al.2001, 2002, 2004). Numerical modeling has shown thatthe few largest offspring consist of materials derivedfrom various well-defined, restricted regions within the


parent (Michel et al. 2002, 2004). Thus, if all ureilitesare derived from a reassembled daughter body, thedistinctive histogram of ureilite Fo values (Fig. 2),which is unlikely to be a result of any petrologic process(Mittlefehldt et al. 1998), may represent a highly biasedsample of the compositional variation on the UPB(Goodrich et al. 2004; Herrin et al. 2010). Furthermore,that sample is not necessarily from the outermost layers(which would be the most accessible regions if we werederiving ureilites from an intact UPB), but could befrom some deep interior region of the UPB.

Modeling of catastrophic disruption and reassemblyalong the lines of Michel et al. (2001, 2002, 2004) cantest the degree to which reassembled bodies (offspring)are select samples of the parent, and also constrain thespecific regions in the parent from which these samplesare derived. This type of investigation was begun byMichel et al. (2004), but that work considered onlysolid parent asteroids, whereas the UPB is thought tohave been partly molten at the time of catastrophicdisruption (Warren and Kallemeyn 1989; Warren 2012;Wilson and Goodrich 2012). Therefore, we have begunan investigation into catastrophic disruption andreassembly of asteroids in various stages of melting(Michel et al. 2013). In the first step of this work, weselected one parent body size (125 km radius) andconsidered four cases—fully solid, fully molten, halfmolten by mass (melt contained in an outer layer), andsolid except for a thin molten layer (8 km thick)centered at 10 km depth. The latter case wouldcorrespond closely to the state of the UPB at about5 Ma after CAI in the fractional melt extraction model(Wilson et al. 2008; Wilson and Goodrich 2012). Ineach case, the parent body is impacted at 5 km s�1 byan 84 km diameter projectile at a 45 angle, resulting incatastrophic disruption. We track the origin (depth inthe parent) of every fragment and then examine thedistribution of the reaccumulated fragments that makeup the three largest offspring. To make the duration ofcomputer runs tractable, the fragments were modeled asapproximately 1 km in size. This is larger than the tensof meters size inferred for UPB fragments from coolingrates (Herrin et al. 2010). However, tests using particlesizes down to a few tens of meters showed that thisresolution limit has no significant effect on the outcome(Michel et al. 2001). A more detailed description of themodeling and results are presented in Michel et al.(2014). To briefly summarize, the results show animportant dependence on degree of melting. In the fullymolten case, the largest offspring all derive their massfrom essentially the complete range of depths in theparent, with little bias. The other three cases all showedsignificant selective sampling. For example, for theparent that was half molten, the three largest offspring

are each dominated by materials derived from arelatively narrow range (approximately 40 km) ofdepths (centered at approximately 40, 60, or 80 kmdepth). For the parent with only a thin molten layer,which may correspond most closely to the UPB, thethree largest offspring are each dominated by materialsderived from somewhat broader ranges of depth(centered at approximately 30, 35, or 60 km depth).These results show clearly that offspring bodies may behighly biased samples of the interior of the parent withthe specific depths that they represent being stronglydependent on the internal structure of the parent at thetime of the disruption. Modeling of this type mayultimately be useful in testing competing models for theigneous petrogenesis of ureilites (e.g., Michel et al. 2013,2014). In the context of the large-scale model presentedhere, the important conclusion is that all ureiliticmaterial that we have ever seen may come from onebiased sample of the UPB.

MIGRATION OF EITHER THE UPB OR A UDB

FROM THE OUTER TO THE INNER BELT

In this section, we present arguments for earlymigration of either the UPB or one of its daughterbodies from the outer to the inner asteroid belt. Then,we report new quantitative modeling of scattering bymassive embryos in the primordial asteroid belt to testwhether this mechanism could have caused the requiredmigration.

As discussed above, the bulk of the evidence pointsto the ureilite parent body (UPB) accreting beyond theice line, thus beyond where the enstatite and ordinarychondrites formed, and being catastrophically disruptedat approximately 5 Ma after CAI. The fact that AhScontained a significant quantity of EC and OC material(indeed, from multiple groups in each class) suggeststhat one of the ureilite daughter bodies (UDBs) fromthis breakup, which would become the source body forall known ureilites to fall on Earth, must have migratedto the inner part of the asteroid belt where it couldaccumulate the EC and OC material (Fig. 1d).Alternatively, the UPB itself could have migratedinwards prior to breakup (Fig. 1b). Especially in thelatter case, additional selection effects are necessary sothat only one daughter body survives to the presentday, or at least that only one of the daughter bodies issuitably located and has a suitable collisional history todeliver meteorites to Earth.

While migrations of this magnitude would bedifficult to achieve in the current dynamicalenvironment of the asteroid belt, processes occurring inthe early solar system were more amenable to suchhistories. Our current understanding of the history of


the asteroid belt suggests that this region originallycontained much more mass, probably exceeding oneEarth mass. Lunar- to Mars-sized planetary embryoswere able to accrete in this region. These embryosexcited the primordial asteroids and scattered them intoresonances with Jupiter and Saturn on a time scale oftens of Ma. The embryos themselves were similarlyremoved from the belt on the same time scale (Wetherill1992; Chambers and Wetherill 2001; Petit et al. 2001;O’Brien et al. 2007). Following this initial phase ofdepletion, there was likely a secondary depletion event(by a factor of a few to approximately 10) as the outerplanets migrated to their final locations, causing secularresonances to sweep through the belt (Gomes 1997;Gomes et al. 2005; Minton and Malhotra 2009;Morbidelli et al. 2010). Finally, over the rest of solarsystem history, the larger asteroids were depleted by afactor of approximately 2–3 by slow chaotic diffusion(Minton and Malhotra 2010), while smaller ones wereground away by collisions. This sequence of eventsmeans that even if many “daughter” bodies wereformed during the breakup of the UPB, only one or asmall number of them may have survived to the presentday.

There are (at least) two mechanisms that could havecaused migration of the UPB or its daughter bodies. (i)Some bodies in the outer belt may have drifted inwardby gas drag and become temporarily trapped in mean-motion resonances with Jupiter, especially the 3:2 and2:1 resonances near 4 AU and 3.2 AU, respectively(Weidenschilling et al. 2001). The relative effectivenessof the 2:1 resonance depends on Jupiter’s eccentricity,although the 3:2 resonance is effective even at near-zerovalues of Jupiter’s eccentricity. The eccentricities of thetrapped bodies would grow and their velocities wouldincrease relative to the nebular gas until gas dragbecame strong enough to pull them out of resonance.Gas drag would then cause their semimajor axes andeccentricities to decay, leaving them in nearly circularorbits in the inner belt. This mechanism requires thepresence of nebular gas. As the solar nebula likelydissipated within approximately 3–10 Ma (Haisch et al.2001), it would have been most effective during the firstfew Ma of solar system history, probably before thebreakup of the UPB at approximately 5 Ma. It alsowould only be effective after Jupiter formed (which mayhave taken 1 Ma or more), leaving a fairly narrowwindow of time for it to operate. Recent simulations oforbital evolution by Hood and Weidenschilling (2012)suggest that drift into resonances via gas drag wouldnot have been effective in the early dynamicalenvironment of the asteroid belt when large embryoswere present. Perturbations and gravitational scatteringby embryos cause orbits of asteroid-sized bodies to

evolve by random walk, rather than systematic orbitaldecay. Consequently, trapping into resonances is lesseffective than implied by the results of Weidenschillinget al. (2001).

(ii) Scattering by encounters with massive embryoscould change the semimajor axis of a body in theprimordial asteroid belt by approximately 1 AU ormore. This would have occurred before the Main Beltwas depleted of embryos, within a few tens of Ma afterCAI. Multiple scattering events would probably beinvolved, as lunar- to Mars-mass embryos are too smallto cause such a change of orbit by a single encounter.Diffusion by multiple encounters would tend to keepeccentricities and inclinations fairly low, or at least notgreater than the average values for a typical embryo,and damping by gas drag would also be effective duringthe first few Ma. This mechanism could potentiallytransport the UPB into the inner Main Belt prior to itsbreakup at 5 Ma, or if it broke up in the outer MainBelt, could transport one or more of its daughter bodiesinward after 5 Ma. Mechanisms (i) and (ii) are notmutually exclusive, and they may have operated ondifferent bodies, depending on sizes (as gas drag is moreeffective on smaller bodies), original orbits, andstochastic effects. However, we consider mechanism (ii)to be more plausible because, as mentioned above,resonance trapping is not very effective. Wedemonstrate below that gravitational scattering isadequate to produce the inferred migration of at leastthe UPB.

We have carried out quantitative modeling ofmechanism (ii) by simulating accretion of embryos fromsmall planetesimals, and integrating their orbits withasteroid-sized test particles. The initial stage of thesimulation uses the accretion code of Weidenschilling(2011), starting with a swarm of planetesimals between1.5 and 4 AU, with initial diameter 0.1 km. This size ischosen because it yields the best fit to the postaccretionsize distribution of asteroids inferred by Morbidelliet al. (2009), although our conclusions do not dependstrongly on this choice. Rapid runaway growth ofMoon- to Mars-sized embryos occurs, and the numberof embryos is essentially complete at model time50,000 yr. At this time, 100 test particles with radius100 km are introduced, evenly spaced in semimajor axis(note that the times mentioned here are from the timeat which planetesimals formed, not from formation ofCAIs). Orbits of embryos and test bodies are integratedwith the Symba symplectic integrator, includingdamping terms due to drag of nebular gas and collisionswith the background population of smaller bodies. Thegas density is approximately 10�10 g cm�3, typical of aminimum-mass solar nebula. Gas drag tends to dampeccentricities and inclinations, and causes secular decay


of semimajor axes. However, for the asteroid-sized testbodies, the effects of gas drag are negligible comparedwith gravitational scattering and collisional damping.Figure 7a shows the orbital eccentricities versussemimajor axes of the embryos and test bodies at 1 Ma.At this time, there are 157 embryos with total mass 1.75Earth masses. Their eccentricities are mostly in therange 0–0.1 due to mutual stirring by the embryos;Jupiter is not included, as it is assumed to have formedlater. Test bodies diffuse from their starting orbits dueto the embryos’ perturbations, as shown in Fig. 7b. The

mean change in semimajor axis is a few tenths of anAU, but a few bodies have migrated over distancesmore than 1 AU from their starting locations, bothinward and outward. Figure 7c shows the migration ofthe test body with the greatest inward displacement; itssemimajor axis has decreased from 3.51 to 2.36 AU in1 Ma. During this time, its eccentricity has varied, buthas typically been <0.05. Collisions with the backgroundpopulation of smaller bodies resulted in net accretion,with a gain of approximately 20% in mass, andapproximately 6 km in radius. The source of the

Fig. 7. a) Eccentricity versus semimajor axis for the ensemble of planetary embryos (open circles) and asteroid-sized test bodies(dots) at model time 1 Ma after the start of accretion from 0.1 km sized planetesimals. b) Diffusion of the orbits of the testbodies: histogram of difference in semimajor axis at model time 1 Ma from the starting value. c) Time evolution of thesemimajor axis of the test body that experienced the largest inward drift. Its starting distance was 3.51 AU; at 1 Ma its distanceis 2.36 AU. d) Amount of mass accreted by the test body in (c) from bodies in radial zones of semimajor axis during itsmigration. Triangle shows its initial mass and distance. e) Mean velocity of bodies impacting the test body versus semimajor axisof the impactors. Most impacts during this first Ma of evolution are at <1 km s�1.


accreted mass ranges from approximately 2.2 to 3.9 AU(Fig. 7d), and the mean impact velocity is <1 km s�1

(Fig. 7e). Most of this mass was acquired outside 3 AU,and would have had a composition similar to that ofthe bulk of the starting body. The small fraction ofmaterial accreted in the inner belt would notsignificantly change the composition of the body, andwould likely be lost if it experienced a large disruptiveimpact. Note that this material does not represent theforeign material in AhS, which would have beenacquired after the breakup of the UPB, several Malater.

While this body’s history is not claimed to be anexact analog of the UPB, it illustrates some relevantfeatures of this migration mechanism. It is possible fora small fraction (approximately 1%) of asteroid-sizedbodies to be moved from the outer to inner belt (or viceversa) by embryo scattering while preserving theobserved radial stratification of the asteroid belt as awhole. Eccentricities remain fairly low during this stage,typical of the rest of the population, and impactvelocities are too low to disrupt a body ofapproximately 100 km radius until the belt is stirred bythe formation of Jupiter. Migration timescales ofapproximately 1 Ma are possible with plausiblenumbers and masses of embryos. This time scalesuggests that the UPB could have been transferred fromthe outer to inner belt before its disruption at 5 Ma.

The breakup of the UPB could potentially producea large number of ureilitic daughter bodies in the innerasteroid belt. As the evidence suggests that only onedaughter body is the source of AhS and all otherureilites, additional selection effects are necessary. Thesecould arise from several sources. The dynamicaldepletion of the asteroid belt would be ongoing at thetime of the UPB breakup 5 Ma after CAI, removingsome (or possibly all but one) of the larger daughterbodies. Then, if there are multiple daughter bodiesremaining, only one of them may be suitably locatednear a resonance and/or has experienced a recentenough breakup event that could deliver a significantquantity of material to near-Earth space.

This argument does not rule out the alternativepossibility that the UPB was disrupted in the outer belt,and that one or a few of its daughter bodies migratedinto the inner belt by this embryo scattering mechanism.That scenario needs to be assessed in a morequantitative manner. This would require long-termorbital integrations of multiple daughter bodiesextending over the time scale for clearing embryos fromthe belt (tens of Ma). However, the timing and impactvelocity necessary to disrupt the few hundred km scaleUPB requires that the disruption occurred after Jupiterformed and exited the belt. Scattering from the outer

belt into an orbit with low eccentricity and inclinationin the inner belt would be more difficult in thisenvironment. The orbit of 2008 TC3, and presumablythe body that was its immediate source, had a lowinclination (Jenniskens et al. 2010). Thus, earlymigration of the UPB before disruption seems morelikely.

FORMATION OF A DEEP, DIVERSE REGOLITH

ON THE UDB

In this section we argue that once the ureilitedaughter body (UDB) was in the inner asteroid beltregion (either because it formed there or because itmigrated there), its surface layers would have evolved toincorporate the foreign clasts found in Almahata Sitta.

As modeled above, the UDB was likely formedfrom gravitationally reassembled fragments followingthe breakup of the original UPB, and may have beenessentially a ureilitic rubble pile when it was formed.Both empirical and theoretical evidence suggest thatthrough further collisional processing, small solarsystem bodies such as the UDB can rapidly develop asubstantial layer of regolith on their surfaces, consistingof compacted and variously consolidated fine-grainedmatrix containing rocky fragments.

Spacecraft have provided close-up views of asteroidsurfaces. Figure 8a shows a portion of the surface ofthe approximately 35 km long asteroid Eros, having atypical regolith surface with a fine-grained matrixsurrounding larger rock fragments. Faintly visible oldcraters that have been eroded by repeated small impactsalso suggest significant regolith development. Figure 8bshows a similar-sized section of the much smallerapproximately 500 m long asteroid 25143 Itokawa, withan abundance of fragments on its surface possiblyhinting at a deeper rubbly structure, and anapproximately 100 m scale “pond” of fine-grainedmaterial with few rocks larger than approximately 5 m.These “ponds” are found in gravitational potential“lows” on asteroids, suggesting migration of finematerial on the surface. Measurements from lunarcrater counts indicate that the impact flux at 3.8 Ga agowas approximately 102 times the present flux (Neukumet al. 2001), suggesting that substantial regolithproduction occurred relatively early in solar systemhistory. Thus, it is to be expected that a body like theUDB would accumulate a substantial volume ofregolith over the lifetime of the solar system.

We have argued above that clast compositions andsizes in 2008 TC3 were similar to those of otherpolymict ureilites, which are inferred to have beenderived from a surface or near-surface regolithenvironment (Goodrich et al. 2004; Downes et al. 2008).


However, from several lines of evidence it is clear that2008 TC3 was less well consolidated than other polymictureilites. Thus, we suggest that 2008 TC3 was afragment derived from outer layers of regolith formedon the UDB, while other polymict ureilites were derivedfrom deeper levels in the regolith (Figs. 1e and 1f).

Horstmann and Bischoff (2014) argued that 2008TC3 could not have been derived from regolith becauseno solar wind implanted gases were found in twoureilitic samples of AhS (Ott et al. 1993). However,solar gases are only implanted in the top millimeter ofan asteroidal surface (Wieler et al. 1986), and evenextensive gardening may not redistribute affected grainshomogeneously in subsurface regolith layers.Furthermore, the solar gases in other polymict ureilitesare known to be heterogeneously distributed (Ott et al.1993; Rai et al. 2003), implying that some bulk samplesof even strictly defined regolith breccias can be devoidof them. Thus, the absence of solar gases in materialthat represents a very small fraction of the mass of theasteroid does not rule out 2008 TC3 forming as regolith.

2008 TC3 itself must have been a surviving portionof a larger body, such as the proposed UDB, since a2008 TC3-sized body has a collisional lifetime much lessthan the age of the solar system (Farinella et al. 1998;Bottke et al. 2005). Given that the UDB could havedeveloped substantial regolith, as discussed above, couldthe diverse array of foreign material found in AhS bethe macroscopic remnants of some of those impactors(or perhaps even small intact impactors themselves)?Experiments have shown that at low speeds, a thingranular layer whose thickness is only a fraction the sizeof an impactor strongly inhibits the rebound velocity in

a low-velocity bouncing collision (Hartmann 1978;Colwell and Taylor 1999). Even at moderately highspeeds, a granular layer allows an impactor to “burrow”to deep layers and disperse energy in mechanicalfriction, thus reducing ejecta and facilitating capture ofthe impactor ejecta (Hartmann 1978, 1985; Melosh1989). “Ponds” of fine-grained material, as in Fig. 8b,may be extremely efficient at allowing incomingmeteoroids to be captured into the fine-grained matrix.To cite a conclusion from Hartmann (1978), “regolithbegets regolith.”

What is uncertain, however, is the maximumvelocity at which foreign impactors (such as ECs andOCs impacting ureilitic regolith on 2008 TC3) or theirfragments could be buried and preserved in the regolithpowders, rather than being pulverized or vaporized. Toour knowledge, little has been published on preservationof impactors as a function of velocity, becauseexperiments that might be relevant are usually focusedon cratering effects instead. The mean impact velocityin the current asteroid belt is roughly 5 km s�1 (and itlikely reached that value within a few tens of Ma of theformation of the solar system), although the distributionof velocities follows a broad curve such that impacts atsignificantly lower and higher velocities are also possible(e.g., Bottke et al. 1994). Schultz and Gault (1984)studied impacts into regolith-like powder targets andfound aluminum projectiles severely deformed at2.0 km s�1 and fragmented at 2.8 km s�1, while morebrittle pyrex projectiles completely comminuted at>1 km s�1. At least one experimental study reportspreservation of aluminum projectile fragments from acollision at 4.5 km s�1 (Flynn et al. 2013) into pumice

Fig. 8. a). A approximately 250 m portion of the surface of approximately 35 km long asteroid Eros, showing regolith structure.Faintly visible old craters that have been eroded by repeated small impacts also suggest regolith development (NEAR Mission,NASA/JHUAPL). b). A approximately 250 m wide portion of approximately 500 m long asteroid 25143 Itokawa, showingapparent rubble pile structure. The image includes a “pond” where a mass of fine-grained material is concentrated in a region oflow gravitational potential (Hayabusa image, JAXA).


targets. If this can occur, it suggests that the low-velocity tail of asteroid-belt collisions could potentiallyimplant a significant amount of foreign fragments into afine-grained regolith of ureilitic material.

Gayon-Markt et al. (2012) carried out aquantitative analysis to determine whether such aprocess could have led to formation of 2008 TC3. First,they compared spectra of asteroids to meteorite typesfound in AhS and identified the region around theNysa-Polana asteroid family as the most likely sourceregion of 2008 TC3 (as also suggested by Jenniskenset al. 2010), due to its particularly diverse range ofspectral types. Then, taking what they called an“arbitrary” upper limit impact speed of 0.5 km s�1 forcapture of foreign fragments and assuming the currentmass and taxonomic distribution of potential impactors,they calculated how much material could be capturedby a body in the inner asteroid belt. They found thatinsufficient foreign material would be captured, eitheronto 2008 TC3 itself (as a free-floating body), or ontothe surface of a larger body (such as the UDB) that wasthe progenitor of 2008 TC3.

There are several reasons why the analysis ofGayon-Markt et al. (2012) might underestimate thepossibility of forming 2008 TC3 by impact implantationof foreign clasts into the surface layers of the UDB.First, as we suggest above, the implantation of foreignmaterial into regolith could potentially occur at highervelocities than 0.5 km s�1 (Schultz and Gault 1984;Flynn et al. 2013), though further experimental studiesdesigned to test this hypothesis are badly needed. Thiswould mean that a larger amount of foreign materialcould be preserved. The larger impact rate in the past,especially prior to approximately 3.8 Ga ago when mostof the foreign clasts in AhS would have been acquired(see below), would also boost the implantation rate.Even more importantly, Gayon-Markt et al. (2012)implicitly assume that the part of 2008 TC3 that waslost during atmospheric entry had the same distributionand proportions of meteorite types (in particular,ureilitic versus foreign) as the fragments recovered onthe ground. However, if the lost material was primarilycomposed of pulverized ureilitic material, as we argueabove, then 2008 TC3 would have had a much lowerconcentration of foreign clasts than Gayon-Markt et al.(2012) assumed, i.e., ≤1% rather than 20–30%.Incorporating this small fraction of foreign impactorsmay be much less difficult than the results of Gayon-Markt et al. (2012) suggest.

In any case, it is clear that the 2008 TC3 parentbody (the UDB in our model), began incorporatingforeign fragments very early in solar system history.Meier et al. (2012) reported K-Ar retention ages of4.4 Ga and 4.7 Ga for two chondritic clasts (L4 and

H5, respectively) from AhS. Turrin et al. (2013) andBeard et al. (2013) reported 40Ar-39Ar measurements forsamples of the same two clasts. The former obtainedeight Ar-Ar fusion ages and plateau ages all fallingbetween 4.57 and 4.49 Ga. The latter obtained poorlydefined plateaus, with best summed ages of >4.2 Ga and4.17 Ga (for the L4 and H5, respectively). Meier et al.(2012) also obtained 4He retention ages (by assumingplausible chondritic U and Th abundances) of 3.8 Gafor both clasts. All three studies agree that thesechondritic clasts have experienced no significantreheating events since very early in solar system history.Beard et al. (2013) concluded that either thesefragments were incorporated into the ureilite parentbody >4200 Ma ago, or the collisional events thatimplanted them occurred at low enough velocity toprevent thermal metamorphism. However, based onevidence for an early pre-irradiation phase in the clasts,Meier et al. (2012) inferred that the incorporation of thechondrite fragments into the ureilite host must havetaken place ≥3.8 Ga ago.

Meier et al. (2012) also suggested that the 3.8 Ga4He age might record the time of a major breakup of aureilitic daughter body to form a third generation body.This age is associated with what appears to be a periodof intense impacts throughout at least the inner solarsystem during the 4.1–3.5 Ga interval (the Late HeavyBombardment or LHB), and has been seen in ureilitesbefore. Some ureilites have been found to have Ar-Arages consistent with the time of the LHB (Bogard andGarrison 1994). Also, some show a Sm-Nd isochron ageof approximately 3.8 Ga (Goodrich et al. 1991;Goodrich and Lugmair 1995), which cannot representigneous crystallization (Goodrich et al. 2010b). Ittherefore seems likely that the time of the LHB was oneof significant regolith development on the UDB.However, we disagree with the suggestion that a thirdgeneration offspring body formed at this time. The3.8 Ga Sm-Nd age does not reflect major resetting ofthe silicate minerals in ureilites (Goodrich and Lugmair1995). It probably records a subcatastrophic impactevent that caused local reheating on the UDB.

Gayon-Markt et al. (2012) argued that if intactimpactor implantation could occur at the current meanasteroidal collision velocity of 5 km s�1, AhS-typemeteorite falls would be much more common thanobserved. The same could possibly be said even ifimplantation could only occur at lower, but still notuncommon, velocities, as we argue here. This impliesthat other factors must also be responsible for the rarityof AhS-like objects in the meteorite record, or even therarity of AhS-like meteorites among ureilites (assumingthey all originated on the same daughter body, as wepropose). An important factor is likely that only the


outermost layers of regolith will have a porous, weaklycemented structure of fine regolith with some largerpieces of foreign and ureilitic material, whereas thedeeper layers (which will volumetrically dominate thematerial ejected in a cratering event or catastrophicdisruption of the ureilite daughter body) will consist ofthe much more common polymict (more cohesive,fragmental breccias, with less distinction between matrixand clasts) and main group ureilites. It is also possiblethat only ideal local conditions in small areas of thesurface favor the formation of a foreign clast-richregolith—perhaps it requires formation in a pond-likedeposit of deep, fine regolith like that seen in Fig. 8b.Another significant factor is that porous objects like2008 TC3 are more likely to break up entirely uponatmospheric entry, leading to no large single meteorite.Finally, as also noted by Gayon-Markt et al. (2012),Almahata Sitta is the first meteorite fall for which wehad advance warning and concerted tracking andrecovery efforts were made, and this certainlyconstitutes an additional selection effect.

RECENT BREAKUP OF THE UDB AND DELIVERY

OF UREILITES TO EARTH

We have suggested that all ureilitic material thathas ever been found on Earth resided in a UDB in theinner asteroid belt since at least 5 Ma after CAI. Theliberation of main group ureilites from the interior ofthis body (Fig. 1f) requires either a large cratering eventor complete disruption of the UDB. In this section wediscuss evidence for when this might have happened,and mechanisms by which UDB fragments could havebeen delivered to Earth.

The 21Ne cosmic ray exposure ages for main groupureilites range from <1 to 33 Ma (Fig. 9). For polymictureilites other than AhS (only four samples have beenanalyzed) they range from approximately 9 to 46 Ma.Six ureilitic clasts (Welten et al. 2010) and two OCclasts (Meier et al. 2012) from AhS have 21Ne ages ofapproximately 20 Ma. This large spread in CRE ages,with no apparent groupings or systematic differencesbetween polymict and main group samples, suggestsseveral possible breakup scenarios. One possibility isthat the UDB experienced a complete breakup (at46 Ma or older) in which various “large” fragmentswere created. Subsequently, these fragments experiencedone or more additional collisions that broke them intosmaller fragments at various times. Alternatively, the46 Ma and 33 Ma CRE ages (Fig. 9) may representisolated cratering events, with complete breakup of theUDB then occurring at approximately 20 Ma. Asteroid2008 TC3 would have been created as an object a fewmeters in size during this breakup. As discussed by

Meier et al. (2012), this body was small enough that ifit was released from a larger fragment approximately20 Ma ago, the CRE age clock would have started forall incorporated clasts at that time, leading to identicalages for the ureilitic and OC samples of AhS (Fig. 9).

Fragments of the UDB could have been transferredfrom the inner main-belt to Earth-crossing orbits byseveral strong dynamical resonances, the 3:1 mean-motion resonance with Jupiter at a = 2.5 AU, and theυ6 secular resonance with Saturn, which defines theinner edge of the Main Belt at roughly 2.1 AU for zero-inclination bodies and somewhat further out asinclination increases (Gladman et al. 1997). If theureilite daughter body was located in the Nysa-Polanaregion, that would place it within 0.1 AU of the 3:1resonance. Both resonances can increase the eccentricityof a body to approximately 1 on Ma timescales (e.g.,Farinella et al. 1994), and if material from a disruptedasteroid is directly injected into a resonance, this willlead to much faster transfer times to Earth thansuggested by cosmic-ray exposure ages (e.g., Morbidelliand Gladman 1998). However, the Yarkovsky Effect, anongravitational force resulting from the absorptionand reemission of solar radiation, can allow bodies onthe order of meters to tens of meters to driftsignificantly in semimajor axis (e.g., Farinella et al.1998), such that direct injection of material intoresonances is not required. Simulations incorporating

Fig. 9. 21Ne exposure ages of 18 main group ureilites, 6ureilitic clasts in AhS, 4 other polymict ureilites, and 2 OCclasts in AhS. Data for main group ureilites and polymictureilites other than AhS from Aylmer et al. (1990), Rai et al.(2003) and compilation in Goodrich (1992). Multiple analysesfor the same meteorite have been averaged. Data for ureiliticclasts in AhS from Welten et al. (2010). Data for OC clasts inAhS from Meier et al. (2012).


Yarkovsky drift into dynamical models of meteoritedelivery yield results consistent with the CRE agedistributions of major meteorite groups (e.g., Bottkeet al. 2000).

COMPARISION WITH OTHER MODELS FOR

FORMATION OF 2008 TC3 AND ALMAHATA

SITTA

Herrin et al. (2010) and Horstmann and Bischoff(2014) both discussed models for the formation of 2008TC3 which, like ours, depend on early catastrophicdisruption of the UPB and posit that 2008 TC3

originated from one of the daughter bodies formed inthis event. Herrin et al. (2010) proposed that after thebreakup of the UPB its daughter bodies experiencedfurther collisional events over some span of time, inwhich encounters with a variety of asteroid types led toacquisition of the chondritic materials now seen inAlmahata Sitta. This model is generally similar to ours,although it is not specific as to where and when theUPB was disrupted or the subsequent collisionsoccurred.

In contrast, Horstmann and Bischoff (2014) andBischoff et al. (2010b) propose that the nonureiliticmaterials in 2008 TC3 were acquired at the time of UPBbreakup—specifically, they were “reaccreted” along withfragmented ureilitic material to form mixed daughterbodies with high proportions of chondritic lithologies,as in AhS. Horstmann and Bischoff (2014) cite supportfrom Gayon-Markt et al. (2012) in suggesting that alltypes of chondritic fragments may have been present as“pebbles” in a debris disk around the Sun at the time ofUPB breakup, and could have been mixed intoreaccreting ureilitic bodies at low collisional velocities.We argue that this model is unlikely, for severalreasons.

First, it is unlikely that there was a large mixture ofmeteorite types anywhere in the asteroid belt, either aslarge bodies or small debris, at 5 Ma after CAI (theprobable time of UPB breakup). The large-scalevariation in asteroid types with heliocentric distance(E?S?C) was probably established within the first fewMa of solar system history, controlled by thetemperature gradient with heliocentric distance. Thismeans that even if a significant amount of pebble-sizeddebris was present, most of it would have been locallyderived. Thus, the outer belt, where the UPB accreted,would not have contained significant EC and OCmaterials. While the Nysa-Polana region noted byGayon-Markt et al. (2012) as a possible source for 2008TC3 does contain diverse materials, it is highly unlikelythat the family or families in that region formed in thefirst few tens of Ma of solar system history. If they had,

the scattering and depletion occurring in the asteroidbelt at that time would have rendered themundetectable today. Recent work suggests that at leastsome of the breakup events in the Nysa-Polana regionoccurred less than 2 Ga ago (Walsh et al. 2013). Hence,these families would not have been present at the timeof UPB breakup. In addition, if mixed types of debrishad been widespread in the early solar system, thenmixed meteoritic breccias should be common when infact they are extremely rare (Bischoff et al. 2006).

Second, even if mixed debris was present in theregion of the UPB breakup, it is unlikely that much ofit would have reaccreted along with the reassemblingureilite fragments. The fragments of the UPB wereoriginally gravitationally bound to each other, andreassembled on a free-fall time scale of days to weeks,allowing little chance for capture of exogenous material.Gayon-Markt et al. (2012) suggested that pebble-sizedparticles of mixed debris were concentrated and mixedwith ureilitic debris by turbulent eddies or streaminginstability in the nebular gas. However, by 5 Ma afterCAI, most of the original mass of solids in the solarnebula was already in the form of large planetesimalsand planetary embryos; pebble-sized fragments wouldbe more likely to accrete to existing large bodies than toform new ones. Admittedly, nebular gas could havedamped velocities of small particles (Weidenschillinget al. 2001) and this damping would aid their accretionto the reassembling UDB, but the amount of gaspresent would have been greatly diminished at 5 Maafter CAI. Thus, the amount of foreign material thatwould have “re-accreted” to the UDB (especially giventhe short time scale for daughter body reassembly)would have been quite small, much less than the highproportions required in the Horstmann and Bischoff(2014) model.

In summary, the salient features of our model thatdistinguish it from other models that have beenproposed for the origin of 2008 TC3 are (1) that 2008TC3 contained only a small component of nonureiliticmaterials, similar to other polymict ureilites; (2) that allureilitic material that has ever been sampled on Earthcomes from the immediate parent of 2008 TC3, whichoriginated as an offspring body of the UPB and is likelya biased sample of the UPB; and (3) that a stochasticprocess such as the long-range migration in heliocentricdistance that we propose for the UPB is required toexplain the unique diversity of foreign materials inpolymict ureilites.

Acknowledgments—We thank Addi Bischoff, PeterJenniskens, Noriko Kita, Susanne Schwenzer, and MikeZolensky for interesting and helpful discussions. Thefirst author is indebted to Michael J. Drake for many


years of support, encouragement, and interest inureilites. We are grateful to reviewers Edward Scott,Allan Treiman, and the associate editor John Jones fortheir helpful comments and suggestions, which led tosignificant improvement of this manuscript. This workwas supported by NASA OSS grant #NNX12AI84Gand NASA Cosmochemistry grant #NNX12AH74G.P.M. acknowledges support from the French spaceagency CNES.

Editorial Handling—Dr. John Jones

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