planetary and pre-solar noble gases in meteorites

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Contents lists available at ScienceDirect

Chemie der Erde

j o ur na l ho mepage: www.elsev ier .de /chemer

nvited Review

lanetary and pre-solar noble gases in meteorites

lrich Otta,b,∗

University of West Hungary, Faculty of Natural Sciences, Savaria Campus, H-9700 Szombathely, HungaryMax Planck Institute for Chemistry, Hahn-Meitner-Weg 1, D-55128 Mainz, Germany

r t i c l e i n f o

rticle history:eceived 22 August 2013ccepted 29 January 2014

eywords:eteoritesoble gaseslanetary noble gaseshase Qresolar noble gases

a b s t r a c t

Noble gases are not rare in the Universe, but they are rare in rocks. As a consequence, it has been possibleto identify in detailed analyses a variety of components whose existence is barely visible in other ele-ments: radiogenic and cosmogenic gases produced in situ, as well as a variety of “trapped” components –both of solar (solar wind) origin and the “planetary” noble gases. The latter are most abundant in the mostprimitive chondritic meteorites and are distinct in elemental and isotopic abundance patterns from plan-etary noble gases sensu strictu, e.g., those in the atmospheres of Earth and Mars, having in common onlythe strong relative depletion of light relative to heavy elements when compared to the solar abundancepattern. In themselves, the “planetary” noble gases in meteorites constitute again a complex mixture ofcomponents including such hosted by pre-solar stardust grains.

The pre-solar components bear witness of the processes of nucleosynthesis in stars. In particular,krypton and xenon isotopes in pre-solar silicon carbide and graphite grains keep a record of physicalconditions of the slow-neutron capture process (s-process) in asymptotic giant branch (AGB) stars. Themore abundant Kr and Xe in the nanodiamonds, on the other hand, show a more enigmatic pattern,which, however, may be related to variants of the other two processes of heavy element nucleosynthesis,the rapid neutron capture process (r-process) and the p-process producing the proton-rich isotopes.

“Q-type” noble gases of probably “local” origin dominate the inventory of the heavy noble gases (Ar,Kr, Xe). They are hosted by “phase Q”, a still ill-characterized carbonaceous phase that is concentratedin the acid-insoluble residue left after digestion of the main meteorite minerals in HF and HCl acids.While negligible in planetary-gas-rich primitive meteorites, the fraction carried by “solubles” becomesmore important in chondrites of higher petrologic type. While apparently isotopically similar to Q gas,the elemental abundances are somewhat less fractionated relative to the solar pattern, and they deservefurther study. Similar “planetary” gases occur in high abundance in the ureilite achondrites, while smallamounts of Q-type noble gases may be present in some other achondrites. A “subsolar” component, pos-sibly a mixture of Q and solar noble gases, is found in enstatite chondrites. While no definite mechanismhas been identified for the introduction of the planetary noble gases into their meteoritic host phases,there are strong indications that ion implantation has played a major role.

The planetary noble gases are concentrated in the meteorite matrix. Ca-Al-rich inclusions (CAIs) arelargely planetary-gas-free, however, some trapped gases have been found in chondrules. Micrometeorites(MMs) and interplanetary dust particles (IDPs) often contain abundant solar wind He and Ne, but they
are challenging objects for the analysis of the heavier noble gases that are characteristic for the planetarycomponent. The few existing data for Xe point to a Q-like isotopic composition. Isotopically Q-Kr and Q-Xe show a mass dependent fractionation relative to solar wind, with small radiogenic/nuclear additions.They may be closer to “bulk solar” Kr and Xe than Kr and Xe in the solar wind, but for a firm conclusionit is necessary to gain a better understanding of mass fractionation during solar wind acceleration.
© 2014 Elsevier GmbH. All rights reserved.

Please cite this article in press as: Ott, U., Planetary and pre-solar noble gases in meteorites. Chemie Erde - Geochemistry (2014),http://dx.doi.org/10.1016/j.chemer.2014.01.003

∗ Correspondence to: Königsberger Str. 60a, D-55268 Nieder-Olm, Germany. Tel.: +49 1 160 5467230; fax: +49 6136 959875.E-mail addresses: [email protected], [email protected]

ttp://dx.doi.org/10.1016/j.chemer.2014.01.003009-2819/© 2014 Elsevier GmbH. All rights reserved.

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U. Ott / Chemie der Erde xxx (2014) xxx–xxx

ontents

1. Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002. Planetary noble gas patterns . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

2.1. The solar reference composition . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.2. Elemental abundance patterns . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.3. Isotopic patterns . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

2.3.1. Helium . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.3.2. Neon . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.3.3. Argon . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.3.4. Krypton . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 002.3.5. Xenon . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

3. Chondrites: origins, host phases and history. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 003.1. The Q component . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 003.2. Noble gases in pre-solar grains . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

3.2.1. Silicon carbide . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 003.2.2. Graphite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 003.2.3. Nanodiamonds (with glassy carbon). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

3.3. The solubles in carbonaceous and ordinary chondrites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 003.4. Enstatite chondrites: subsolar and sub-Q noble gases . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

4. Non-chondritic meteorites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 004.1. Ureilites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 004.2. Other achondrites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

5. Lesser components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 005.1. Xenon in chondritic metal . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 005.2. Xenon and krypton in silicate-graphite inclusions of the El Taco iron meteorite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 005.3. Xenon in acid residues of iron meteorites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

6. Interrelations, hosts, fractionation and trapping . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 006.1. Interrelations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 006.2. Hosts . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 006.3. Fractionation and trapping . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

7. Related materials . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 007.1. Asteroidal samples . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 007.2. Cometary material . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 007.3. Interplanetary dust particles (IDPs) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 007.4. Micrometeorites (MMs) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

8. Summary and general considerations: planetary versus planets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00Acknowledgments . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00References . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 00

. Introduction

Noble gases are not rare in the Universe, e.g., helium is the sec-nd most abundant element in the Solar System (Lodders et al.,009). They are rare, however, in solids, since – being noble – theyardly take any part in chemical interactions, are among the most

ncompatible elements and generally prefer the gas phase. As anxample, a typical abundance in even the most primitive carbona-eous chondrites for Xe of (in the units commonly used in noble gaseo- and cosmochemistry) 2 × 10−8 cm3 STP/g (Marti, 1967; Mazort al., 1970) corresponds to no more than about 0.1 ppb by weight1 cm3 STP/g = 1 cc/g = 2.668 × 1019 atoms per g).

In spite or (more likely) just because of that, fundamentalnsights in the field of cosmochemistry have come from the studyf the abundances and isotopic compositions of noble gases. Beingo low in abundance, their isotopic compositions are easily influ-nced by “foreign” or “exotic” additions. More often than not thesehow up even in the analysis of bulk materials, in particular intepwise heating analyses that tend to partially separate gases car-ied by different minerals according to their thermal stability. Aase in point is the discovery of presolar grains in primitive mete-rites: isotope abundance anomalies in noble gases were knownong before the discovery of presolar grains per se, such as nanodi-

abundant glassy carbon (see Section 3.2.3); for simplicity in the fol-lowing I will use the term “nanodiamonds” or just “diamonds” todescribe real samples containing both.

Usually noble gases in meteorites occur as a number of discrete“components”, where a component is defined by having a certainisotopic/elemental composition or a certain origin, but need notbe restricted to a certain mineral (“carrier phase”). This includescomponents that are of in situ origin, which means those producedfrom radioactive decay (e.g. 4He, 40Ar), and also those produced bynuclear spallation reactions induced by cosmic rays. In situ compo-nents are often dominant in meteorites of higher petrologic type,which tend to have low abundances of the “trapped” (also called“primordial”) noble gases. While cosmogenic contributions havetraditionally played a significant role in the study of noble gasesand can be used to determine cosmic ray exposure (CRE) ages, theyhave generally been negligible in the study of the isotopic systemsof “solid” elements. This is changing now due to the increased pre-cision with which these elements can be analyzed nowadays. Animportant example is the Hf-W decay system used to study core for-mation of early planetesimals in the early Solar System via extinct182Hf, where it is essential to take into account the cosmic rayinduced effects (e.g., Kruijer et al., 2013; and references therein).

Coming to the trapped noble gases, starting with the work of

Please cite this article in press as: Ott, U., Planetary and pre-solar nhttp://dx.doi.org/10.1016/j.chemer.2014.01.003

monds, silicon carbide and graphite, and it was the search for thearrier phases of these anomalies that ultimately led to the identifi-ation and isolation of these types of grains (e.g., Anders and Zinner,993; Ott, 1993). Note, in this context, that recent work (Stroudt al., 2011) has shown that nanodiamond separates contain also

John Reynolds at Berkeley, much of the work has centered on xenon,

oble gases in meteorites. Chemie Erde - Geochemistry (2014),

so much that an own term “xenology” was coined to describe thisfield of noble gas cosmochemistry (Reynolds, 1963). The first ofthe so-called “gas-rich meteorites” were discovered by Gerling andLevskii (1956), in the same year that Reynolds (1956) published


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dtotoowhtSasmd“cwt

esoondTitr

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Fig. 1. Basic “planetary” patterns as observed in various bulk meteorites and theatmospheres of Earth and Mars. Shown are abundance ratios relative to 132Xe of4He, 20Ne, 36Ar and 84Kr normalized to the solar ratios (Lodders et al., 2009). Datafor CM2 Maribo, CV3 Allende, the ureilite Novo Urei and the enstatite chondriteSouth Oman are from Haack et al. (2012), Ott et al. (1981; and unpublished), Göbelet al. (1978) and Crabb and Anders (1981), respectively (see notes to Table 2). He andNe abundances for these have been corrected for cosmogenic contributions, but maystill include some radiogenic 4He and implanted solar wind. Hence the He/Xe andNe/Xe ratios are upper limits to those for the truly planetary component, for whichthe fractionation relative to solar may be even stronger than plotted. Ar/Kr/Xe data

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etails of his innovative static noble gas mass spectrometer. Soonhereafter it was discovered that the elemental abundance patternsbserved in meteorites basically fell into one of two categories: (a)he gas-rich meteorites with a “solar” pattern – high abundancesf the lightest gases He and Ne compared to the heavy ones, therigin being implanted solar wind and (b) the “planetary” patternith strong elemental fractionation, i.e., strong enrichment of theeavy noble gases relative to the light ones, when compared tohe solar abundance pattern (Signer and Suess, 1963; Pepin andigner, 1965). Some examples – data for various meteorites as wells the compositions of the atmospheres of Earth and Mars – arehown in Fig. 1. Here I will deal with the type of noble gases ineteorites included in the “planetary component” as originally

efined, although the subject has become more complex since then.Planetary” noble gases are most abundant in the most primitivearbonaceous chondrites (Marti, 1967; Mazor et al., 1970), so thisill be dominated by the discussion of chondritic meteorites and

heir noble gases.While there is some basic similarity between the “planetary”

lemental abundance pattern in meteorites and those of the atmo-pheres of Earthand Mars (Fig. 1), there are differences in detail, notnly in the elemental, but also in the isotopic abundances. Whilene might argue, therefore, that the term “planetary” is a mis-omer and should better be avoided (Ott, 2002), this is probablyifficult to achieve since it has been and still is so commonly used.hus, it is important to keep in mind that planetary noble gasesn meteorites are quite unlike those found in the atmospheres ofhe terrestrial planets (a point addressed again at the end of thiseview).

As already mentioned and as shown in the individual chapterselow, planetary noble gases in meteorites are complex and consistf a number of different individual components characterized bynique isotopic (sometimes also elemental) compositions whichre carried by specific host phases. Some – termed isotopicallyapproximately) normal – are not too different in isotopic com-osition from solar noble gases (Pepin et al., 1995; Wieler, 2002;eber et al., 2009, 2012; Crowther and Gilmour, 2013; Meshik et al.,013, 2014) and may be (but need not be) of “local” origin. Othersre carried by presolar grains and show in their isotopic composi-


ions abundance “anomalies” that reflect excess contributions frompecific processes of nucleosynthesis in stars. A general overview isiven in Table 1, where the most important components are listed,ogether with some alternative names sometimes found in the

ig. 2. Elemental abundance ratios in major trapped components, 4He/36Ar vs. 20Ne/36Af Busemann et al. (2000), Huss and Lewis (1994a), Lewis et al. (1994), Göbel et al. (197ifferent samples, for example among different ureilites. In this case as an example data fre shown. For more detail concerning the selected values see footnotes to Table 2 and

lemental abundances was propagated.

for Mars are from Bogard et al. (2001), while the 20Ne/132Xe ratio is from Pepin(1991).

literature, often for neon or krypton/xenon only. Probably the mostabundant and widespread component is the “Q” component (Lewiset al., 1975; Wieler, 1994; Ozima et al., 1998), also named P1 (Husset al., 1996) or (for xenon) OC-Xe (Lavielle and Marti, 1992). Oth-ers, as noted above, are carried by presolar grains, while additionalones, such as ureilite and subsolar noble gases (Göbel et al., 1978;Rai et al., 2003a; Crabb and Anders, 1981; Patzer and Schultz, 2002),are primarily found in special types of meteorites only.


In parts, the present review builds on Ott (2002), with significantupdates and additions. Other useful recent reviews of the field arethose of Podosek (2003) and Wieler et al. (2006).

r (a) and 36Ar/132Xe vs. 84Kr/132Xe (b). Based on listed values or derived from data8) and Crabb and Anders (1981). Variations occur in compositions as observed inor the diamond-rich acid-resistant residue B-1 from Novo Urei (Göbel et al., 1978)

text. Where no errors are given in the original literature, an assumed 10% error in


ARTICLE IN PRESSG ModelCHEMER-25295; No. of Pages 26

4 U. Ott / Chemie der Erde xxx (2014) xxx–xxx

Table 1Important trapped “planetary” noble gas components. Also given: alternative names used in the literature (sometimes restricted to a single gas only).

Component Alternative Carrier phase Remarks

Q P1 Q (primarily) Dominates Ar, Kr, XeP3 Ne: Ne-A1 Presolar diamond Isotopes ∼ normalP6 – Presolar diamond ∼ normal (?)HL Ne: Ne-A2 Presolar diamond Nuclear component (Kr, Xe: r-, p-process?)N – Presolar graphite, SiC Isotopes ∼ normalG Ne: Ne-E(H); Kr, Xe: Kr-S, Xe-S Presolar graphite, SiC Nuclear component; (Kr, Xe: s-process)R Ne-E(L) Presolar graphite Radiogenic 22NeUreilite gases – Diamond, graphite Isotopes ∼ normal

(likely) enstatite Less fractionated

A e ureilites, as well as in sulfide, metal and iron meteorite samples.

2

2

pasciwop(Stsmrecs2aaS

FnoawsPa

Fig. 4. Kr isotopic compositions of various “approximately normal” trapped com-ponents: Q (P1), ureilite, subsolar, the P3 and the “normal” P6 version in presolardiamonds, plus N in presolar SiC (Table 4). The points are slightly shifted alongthe mass axis to aid in visibility. The more exotic components are shown in Fig. 9.Ratios are normalized to 84Kr, and shown are deviations in per mill of the iKr/84Krratios from the corresponding ratios measured for the solar wind (Table 4; Meshiket al., 2014). Not shown are 78Kr/84Kr ratios in the P3, P6 (because of their largeerrors) and N (no value reported) components. The gray shaded area is the possi-ble range estimated for Kr in the Sun’s outer convective zone (OCZ), as discussed inSection 2.1.

Subsolar –

dditional (?) minor (?) components: sub-Q, noble gases in achondrites other than th

. Planetary noble gas patterns

.1. The solar reference composition

The compositions of more or less pure nucleosynthetic com-onents (Sections 2.3 and 3.2) are informative in their own rightnd tell us about the conditions under which the elements wereynthesized. On the other hand, when it comes to less extremeompositions with a possible/probable Solar System origin, it isnstructive to see them in context, hence, it is necessary to compare

ith a reference, as in the following chapters and in Figs. 2–5. Thebvious choice for such a reference is the Solar System abundanceattern, for which the composition of the outer convective zoneOCZ) of the Sun will be used. However, determining exact Solarystem abundances is not an easy thing. Spectroscopic observa-ions with improved models for the solar atmosphere have recentlyignificantly revised downward the abundances of “metals” (ele-ents heavier than He), but are (somewhat) in conflict with the

esults of helioseismology (e.g., Asplund et al., 2009; but also Caffaut al., 2009, and Houdek and Gough, 2011). On the other hand, pre-ise data are available now from the Genesis mission for the bulkolar wind (e.g., Heber et al., 2009; Vogel et al., 2011; Pepin et al.,012; Crowther and Gilmour, 2013; Meshik et al., 2013, 2014), but

Please cite this article in press as: Ott, U., Planetary and pre-solar noble gases in meteorites. Chemie Erde - Geochemistry (2014),http://dx.doi.org/10.1016/j.chemer.2014.01.003

bundances in the solar wind are not exactly the same thing asbundances in the bulk Sun (which dominates abundances in theolar System as a whole).

ig. 3. Isotopic ratios 20Ne/22Ne and 36Ar/38Ar as determined for various trappedoble gas components (Table 3). Solar wind as well as the ratios inferred for the Sun’suter convective zone (Heber et al., 2012) are included for comparison. Not shownre Ne isotopic ratios for the nucleogenic G [Ne-E(H)] and R [Ne-E(L)] components,hich are almost or (probably) entirely pure 22Ne (see discussion in text). Also not

hown is the (theoretical) value for Ar-G. HL may contain small contributions from6. Obvious are the similarities between the Q and ureilite compositions as well as

similarity in 36Ar/38Ar between solar wind and the subsolar component.

Fig. 5. Xe isotopic compositions in various “approximately normal” planetarytrapped components: Q (P1), ureilite, subsolar, the P3 and the “normal” P6 versionin presolar diamonds, plus N in presolar SiC (Table 5). The points are slightly shiftedalong the mass axis to aid in visibility. The more exotic components are shown inFig. 10. Ratios are normalized to 132Xe, and shown are deviations in per mill of theiXe/132Xe ratios from the corresponding ratios measured for the solar wind (Table 5;Meshik et al., 2014). Not plotted is 129Xe in the subsolar component because of theradiogenic contribution in the measured bulk Soth Oman sample (Crabb and Anders,1981). The value for 136Xe in the N component is off scale at +334‰. As for the abun-dance of 126Xe, see discussion in text and Ott et al. (2005). The gray shaded areais the possible range estimated for Xe in the Sun’s outer convective zone (OCZ), asdiscussed in Section 2.1.



U. Ott / Chemie der Erde xxx (2014) xxx–xxx 5

Table 2Elemental abundance ratios (atom ratios) in various trapped planetary components. For comparison the solar abundance ratios according to Lodders et al. (2009) are listedin the last line.

component 4He/132Xe 20Ne/132Xe 36Ar/132Xe 84Kr/132Xe

Q (P1)a 374 ± 72 3.2 ± 0.5 76 ± 7 0.81 ± 0.05P3b (6.7 ± 2.0) × 104 80 ± 53 400 ± 25 2.7 ± 0.4HLb (3.00 ± 0.30) × 105 485 ± 26 50 ± 20 0.48 ± 0.04Gc (3.21 ± 0.45) × 105 108 ± 15 3.55 ± 0.50 0.465 ± 0.066Nc (6.61 ± 0.93) × 105 826 ± 117 79 ± 11 1.468 ± 0.208Ureilite (NU-B1)d 34.6 ± 15.1 0.81 ± .11 204 ± 9 1.15 ± 0.07Subsolare – 7.59 ± 1.07 2660 ± 376 5.86 ± 0.84Solarf 1.745 × 109 2.13 × 106 5.45 × 104 22.10

Variations occur within the individual elemental abundance patterns as observed in different samples (see discussion in text).Where no specific errors are given in the original literature, a 10% error in elemental abundances was assumed and propagated.

a Values listed here for Q (also named P1): the average deduced in the closed system etch analyses of HF/HCl residues by Busemann et al. (2000).b Values listed here for P3 and HL: “best estimate of primitive component” of Huss and Lewis (1994a).c Values listed here for G and N components: as given by Lewis et al. (1994) and/or calculated based on their sample KJ. The “typical” G and N patterns listed and shown are

based on their decomposition of Ne, Kr and Xe for the KJ sample, together with their reported (4He/22Ne)G = 193 and (4He/20Ne)N = 800. For the conversion of 22Ne abundancesinto 20Ne abundances their value for 20Ne/22Ne in the G component of 0.065 (Table 3; Heck et al., 2007, Heck et al. 2009a) and a value of 8.4 in the N component were used.Ar was calculated by partitioning, following their approach, assuming 38Ar/36Ar as 0.660 and 0.1705 in G and N, respectively (Table 3).

d Values listed here for ureilites: as an example data for the acid-resistant residue B-1 from Novo Urei (Göbel et al., 1978) are shown; note however substantial variationsfrom meteorite to meteorite and even among samples of the same ureilite.

ondriw Xe = 3

isbnm(2tcBaecrTsdt

Octit1n(lstNdtu

wa(aeIa

e Values listed here for the subsolar component: as measured in the enstatite chas measured by Okazaki et al. (2010) in a crushing step of EH4 Y 791790: 36Ar/132

f Lodders et al. (2009).

This is primarily because the elements and to some extent alsosotopes are fractionated during ionization and acceleration in theolar atmosphere. Some difference probably exists also betweenulk Sun and the Sun’s OCZ due to gravitational settling, which iseglected in the following considerations, since this is likely to beuch smaller than the differences between OCZ and solar wind

e.g., Turcotte and Wimmer-Schweingruber, 2002; Heber et al.,012). Elements in the solar wind appear fractionated accordingo first ionization potential (FIP) or first ionization time (FIT) as dis-ussed extensively in the literature (e.g., Meyer, 1985; Geiss andochsler, 1985; von Steiger and Geiss, 1989; Geiss, 1998; Gloecklernd Geiss, 2007). For our comparison of the planetary and solarlement patterns (Fig. 1) we follow Vogel et al. (2011), who for dis-ussing the difference between Sun and solar wind compositionely on the solar abundances compiled by Lodders et al. (2009).he approach may be considered circular (Vogel et al., 2011), sinceolar wind abundances were (also) used in defining the solar abun-ances, but Vogel et al. (2011) have also presented arguments thathis is not a serious problem.

Heber et al. (2012) have treated the difference between theCZ and bulk solar wind in the observed He, Ne and Ar isotopicompositions and give two sets of values for the composition ofhe Sun’s OCZ, one derived from a correlation approach (betweensotopic composition and the H/4He ratio), and another based onhe inefficient Coulomb Drag (ICD) model (Bodmer and Bochsler,998, 2000). The differences between the two approaches are sig-ificant only for He. For comparison with the planetary ratiosSections 2.3.1–2.3.3) I choose their solar values from the corre-ation approach: 3He/4He = 3.62 × 10−4 (220‰ lower than in theolar wind), 20Ne/22Ne = 13.36 (16‰ per mass unit lower than inhe solar wind) and 36Ar/38Ar = 5.37 (10‰ per mass unit lower).ote, though, that Pepin et al. (2012) report solar wind ratios thatiffer somewhat (in case of Ne outside reported analytical uncer-ainties) from the values of Heber et al. (2009, 2012), which aresed here.

Equivalent data from Genesis for Kr and Xe isotopes in the solarind have been more difficult to obtain because of the much lower

bundances. Definitive data sets have just now become availableCrowther and Gilmour, 2013; Meshik et al., 2013, 2014). Remark-


bly, these are almost indistinguishable from probably the bestxisting data set derived from lunar samples (Pepin et al., 1995).ntuitively, one expects mass fractionation effects to be consider-bly smaller in Kr and Xe than in the light noble gases (e.g., Wieler,

te South Oman by Crabb and Anders (1981). Note that a slightly higher 36Ar/132Xe100 ± 155, together with 84Kr/132Xe = 5.64 ± 0.28.

2002), but this is not a safe assumption. The actual physical mech-anism for isotope fractionation between Sun and solar wind is notclear at this time, but for an estimate of the difference in case ofKr and Xe it is instructive to use the predictions of the ICD modelas a working hypothesis, since for Ne and Ar it gives quite similarresults as the correlation approach (Heber et al., 2012). Applica-tion of the equations in Heber et al. (2012) requires knowledge ofthe He/H fractionation factor, for which, as Heber et al. (2012), Iuse the (SW ratio/photospheric ratio) = 0.0370/0.085 = 0.435. Alsoneeded are the mean charge states of solar wind Kr and Xe, whichare not known experimentally, but enter critically. Helium is essen-tially fully ionized and Ne almost so (average charge states are +2and +8), but the ionization degree seems to level off after that,with Ar mostly being in the +9 or +10 charge state (Heber et al.,2012). In line with this, Fe, the heaviest element for which mea-sured data appear to be available, has a maximum occurrence inthe +11 charge state (Gloeckler and Geiss, 2007). It seems reason-able, therefore, to assume only slightly higher charge states for Krand Xe, possibly in the range +11 to +15 for Kr and +12 to +16for Xe (see also Gloeckler and Geiss, 1989). For Kr, then the ICDmodel predicts mass fractionation in the range 5‰/amu (+15 state)to 11‰/amu (+11 state), while for Xe the corresponding results varymore strongly with charge state: 5‰/amu (+16 state) to 19‰/amu(+12 state). In each case the solar wind is enriched in the light iso-topes relative to the true solar OCZ composition. To allow for acomparison in Sections 2.3.4 and 2.3.5, “solar” compositions areused that are (somewhat arbitrarily) fractionated by 10 (±5)‰/amuin case of Kr and 12 (±6)‰/amu in case of Xe. The nominal val-ues correspond to mean charge states of +12 and +13, but clearlythere are significant uncertainties. In any case, as discussed below,the fractionation between Sun and solar wind may have importantimplications, when trying to establish a relation between “almostnormal” planetary components and the solar or solar wind pattern.

2.2. Elemental abundance patterns

Elemental abundance ratios for several of the components listedin Table 1 are given in Table 2, and are shown in Fig. 2. Unlike(in general) the isotopic compositions discussed in the following


chapter, most of the elemental compositions investigated showvariations between different analyzed samples. In some cases atleast this is due to metamorphism (cf. Huss and Lewis, 1994a;Busemann et al., 2000). So the elemental ratios listed in Table 2,




Table 3Isotopic compositions of He, Ne, and Ar in various trapped components. (Bulk) solar wind and derived solar ratios (Heber et al., 2012) are included for comparison. Errors inthe last digits are given in parentheses.

Component 3He/4He [10−4] 20Ne/22Ne 21Ne/22Ne 36Ar/38Ar References

Q (P1) 1.23 (2) 10.67 (2); 10.11 (2) 0.0294 (10) 5.34 (2) [1,2,3][1]

P3 ≤1.35 (10) 8.910 (57) 0.029 (1) 5.26 (3) [4]HL ≤1.70 (10) 8.500 (57) 0.036 (1) 4.41 (6) [4]G – < 0.1 < 0.0015 – [5]G (theory) ∼0 0.065 0.00059 1.52 (20) [5–8]N ≤2.6 – – 5.87 (7) [5]R – <0.01 <0.0001 [9]Ureilite – 10.70 (25) – 5.26 (6) [10]

10.4 (3) [11]Subsolar – – – 5.45 (4) [12]

∼2.1 ∼11.65 – ∼5.36 [13]Solar wind* 4.645 (8) 13.777 (10) 0.03289 (7) 5.47 (3) [14]

– 13.972 (25) – 5.501 (5) [15]– 13.74(2) 0.03361 (18) – [16]– 14.00 (4) – 5.501 (14) [17]

– 5.5005 (40) [18]Solar (OCZ) 3.62 13.36 0.03236 5.37 [14]

References: [1] Busemann et al. (2000); [2] Wieler et al. (1992); [3] Busemann et al. (2001a); [4] Huss and Lewis (1994b); [5] Lewis et al. (1994); [6] Gallino et al. (1990), [7]Heck et al. (2009a,b); [8] Heck et al. (2009a); [9] Amari et al. (1995); [10] Göbel et al. (1978); [11] Ott et al. (1985a); [12] Crabb and Anders (1981); [13] Busemann et al.(2001b); [14] Heber et al. (2012); [15] Meshik et al. (2007); [16] Mabry et al. (2009); [17] Pepin et al. (2012); [18] Meshik et al. (2014).Notes: (a) He isotopic ratios can be generally regarded as upper limits only because of omnipresent spallogenic contributions. Radiogenic 4He should make at most a minorcontribution. (b) For the G components experimental upper limits for the Ne isotopic ratios are listed plus (in italics) the theoretical values used in the decomposition of gasesmeasured in samples of presolar silicon carbide. (c) For 3He/4He in P3 a value of (0.45 ± 0.04)x10−4 has been suggested by Busemann et al. (2001a). (d) “solar” (OCZ) values aret /22Ne

v n and

aunmfettOe1ta

tiQosr2lraAmebilt(–wafta

hose derived by Heber et al. (2012) using the correlation approach; except for 21Nealue (Heber et al., 2009, 2012) and an assumed fractionation between Ne in the Su

nd shown in Fig. 2, are kinds of “typical” or “average” values (Q,reilites) or best estimates of “primitive” compositions (see foot-ote to Table 2 for details). For ureilites, where all ratios vary byore than an order of magnitude and often within samples taken

rom a single meteorite in a manner as yet not understood, it isspecially difficult to define a representative composition. Onlyhe subsolar composition (for lack of better knowledge) is simplyhat found by Crabb and Anders (1981) in the E-chondrite Southman. Where no explicit errors are reported in the original lit-rature, errors have been assigned for this compilation assuming0% uncertainties in elemental abundance determinations. Uncer-ainties due to model-dependent assumptions were not taken intoccount.

All patterns are fractionated, relative to the solar abundance pat-ern (Figs. 1 and 2), to various degrees, however. Extremely strongs the fractionation of He and Ne relative to the heavy gases for

and ureilites. The similarity of the elemental patterns and alsof isotopic compositions (Ott et al., 1985a; Busemann et al., 2000;ee Section 3) has led to speculation that these components areelated, if not identical (Ott et al., 1984, 1985b; Busemann et al.,000). Gases contained in diamonds and in SiC generally show a

ess fractionated pattern; but given that their origin is not directlyelated to “solar system components”, a comparison between thesend the solar elemental abundance ratios is of limited usefulness.lso less fractionated is the subsolar component found in enstatiteeteorites, which has been the reason for giving it its name. How-

ver, there appears to be a general break in fractionation patternsetween light and heavy noble gases. The subsolar component, e.g.,

s less fractionated in Ar/Xe and Kr/Xe than the other componentsisted in Table 2, but Ne/Xe is among the most fractionated, whilehe G component in presolar SiC shows just the opposite behaviorFig. 2). A similar effect – a break between light and heavy gases

has also been seen by Busemann et al. (2000) in the variationsithin the Q component alone. Possibly more than one fraction-


tion mechanism has been at work, and the one(s) dominatingractionation among the heavy noble gases are not the same ashose primarily responsible for the fractionation between heavynd light ones.

(not reported by these authors), which has been calculated based on the solar wind solar wind of 16 ‰/amu (Heber et al., 2012). For more detailed discussion see text.

2.3. Isotopic patterns

Given that in many cases the elemental compositions of the vari-ous components are variable, isotopic compositions (Tables 3–5) aremore characteristic for the presence of a specific component. ForHe, Ne and Ar comparison is made with solar wind compositions asdetermined on Genesis samples as well as the “true” solar compo-sitions derived from these (Heber et al., 2009, 2012). In case of Krand Xe,the “best current estimate” of Meshik et al. (2014) is used,together with an estimate of the fractionation between Sun andsolar wind as discussed above (Section 2.1) There are basically twoways in which the compositions of the individual trapped compo-nents have been derived: (a) consistent determinations in a varietyof samples and/or release steps where they were assumed to occurin essentially “pure” form (e.g. Q, ureilites); (b) in cases where pre-dominantly mixtures of various components were measured, fromobserved mixing lines and extrapolation to an assumed value forone endmember isotopic ratio (the P3, HL, G and N components).As for the latter approach, the value chosen for extrapolation isarbitrary in principle, although there are of course arguments forthe specific choices. This limitation should be kept in mind whenassessing the meaning of the different compositions and of the dif-ferences between the various near-normal components. Only inthe case of the G component hosted by presolar silicon carbideand (for Kr and Xe) corresponding to material synthesized in the s-process, the assumption involved in the decomposition (essentiallyzero 136Xe/130Xe in s-process-Xe, Section 3.2.1) is based on a solidand well-founded physical understanding. Generally also not takeninto account are possible fractionation effects during trapping (butsee Section 3.2.3 on the P3 and HL components in nanodiamond).

2.3.1. HeliumCompositions are listed in Table 3. Given the small grain size

of the most important carrier phases (Q and diamond), these must


have acquired basically the surrounding matrix concentrationsof spallogenic 3He and radiogenic 4He, due to recoil out of othermatrix material. With 3He/4He generally low (∼10−4) in trappedcomponents, and 3He abundantly produced by cosmic ray




Table 4Isotopic composition of Kr in various trapped planetary components. Also given for reference is the composition of the solar wind (Meshik et al., 2014).

Component 78Kr/84Kr 80Kr/84Kr 82Kr/84Kr 83Kr/84Kr 86Kr/84Kr References

Q (P1) 0.00603 ± 0.00003 0.03937 ± 0.00007 0.2018 ± 0.0002 0.2018 ± 0.0002 0.3095 ± 0.0005 [1]P3* 0.0064 ± 0.0010 0.0395 ± 0.0004 ≡0.2023 0.2032 ± 0.0005 0.3128 ± 0.0006 [2]HL* 0.0042 ± 0.0010 0.0305 ± 0.0010 ≡0.1590 0.1989 ± 0.0010 0.3623 ± 0.0018 [2]P6* 0.0059 0.0381 ± 0.0008 ≡0.2023 02013 ± 0.0010 0.3147 ± 0.0030 [2]P6(exotic)* 0.0052 0.0354 ± 0.0004 ≡0.1839 0.1999 ± 0.0005 0.3348 ± 0.0012 [2]G*× – 0.0133–0.0183 ≡0.4167 0.1192 ± 0.0054 0.454–1.176 [3]N* – 0.03962 ± 0.00040 0.2028 ± 0.0021 0.2018 ± 0.0032 0.2842 ± 0.0036 [3]Ureilite 0.00601 ± 0.00008 0.0399 ± 0.0005 0.2037 ± 0.0018 0.2026 ± 0.0015 0.3091 ± 0.0012 [4]Subsolar 0.00634 ± 0.00016 0.04050 ± 0.00029 0.2045 ± 0.0009 0.2027 ± 0.0007 0.3073 ± 0.0017 [5]Solar wind 0.00642 ± 0.00005 0.0412 ± 0.0002 0.2054 ± 0.0002 0.2034 ± 0.0002 0.3012 ± 0.0004 [6]

References: [1] Busemann et al. (2000); [2] Huss and Lewis (1994b), renormalized following Busemann et al. (2000); [3] Lewis et al. (1994), [4] Göbel et al. (1978), avg. of 7ureilites, except for 78K/84Kr (ureilite Kenna; Wilkening and Marti, 1976). [5] Crabb and Anders (1981), South Oman. [6] Meshik et al. (2013).Notes: For a discussion of differences between bulk Sun and solar wind see Section 2.1. * Compositions for G and N components in SiC derived from mixing lines, involvinga 86Kr/S

sb3

cHlbuisrfmhcQrriteiSiodpQ

TI

URuNrolt

ssumptions about 84Kr/82Kr ratios in endmembers (see text). x Ratios 80Kr/82Kr andiC. For graphite, 86Kr/84Kr extends to even higher values.

pallation, 3He generally is more strongly affected than 4Hey extraneous contributions. Precise compositions of trappedHe/4He have been obtained only in a few cases, i.e., for the Qomponent (Busemann et al., 2000, 2001a) as well as the P3 andL gases hosted by presolar diamond (Huss and Lewis, 1994b), all

ying in the range 1.2–1.7 × 10−4, with P3 possibly characterizedy a lower ratio of ∼0.45 × 10−4 (Busemann et al., 2001a). A similarpper limit of 2.6 × 10−4 has been inferred for the N component

n presolar silicon carbide (Lewis et al., 1994). However, strictlypeaking, even in the best cases obtained 3He/4He ratios must beegarded as upper limits. The exact size of the cosmogenic inter-erence depends both on the cosmic ray exposure age of the host

eteorite and on the concentration of trapped He in the respectiveost phase and, hence, is minor in phases like diamond whichontains trapped He in high concentration. The 3He/4He ratio in the-component is based on closed system etch analyses of HF/HCl-

esistant residues (Busemann et al., 2000, 2001a), which selectivelyelease Q gases as operationally defined (see Section 3.1). The valuen Table 3 of 1.23 × 10−4 (Busemann et al., 2001a) was obtained inhe analysis of a residue from the Isna CO3 chondrite, which has anxtremely short cosmic ray exposure. The exact value for 3He/4Hen Q has wider implications, e.g., in cosmology (Yang et al., 1984;chramm and Turner, 1998; Iocco et al., 2009). If the Q components indeed representative of the primordial noble gas component


f the solar system (Wieler, 1994; Ozima et al., 1998; see alsoiscussion of Xe in Q below), (3He/4He)Q should correspond to therotosolar value. In turn, from the difference between 3He/4He in

and 3He/4He in the solar wind, it should be possible to determine

able 5sotopic composition of Xe in various trapped components. Also listed for comparison is t

Component 124Xe/132Xe 126Xe/132Xe 128Xe/132Xe 129Xe/132Xe

Q (P1) 0.00455 (2) 0.00406 (2) 0.0822 (2) 1.042 (2)

P3* 0.00446 (6) 0.00400 (4) 0.0806 (2) 1.042 (4)

HL* 0.00839 (9) 0.00564 (8) 0.0905 (6) 1.056 (2)

P6* 0.00433 (25) 0.00440 (28) 0.0890 (20) 1.114 (8)

P6(exotic)* 0.00679 (8) 0.00516 (8) 0.0899 (5) 1.078 (2)

G× ≡0 ≡0 0.2159 (23) 0.1182 (112)

N× 0.00470 (13) 0.00357 (18) 0.0785 (11) 1.000 (14)

Ureilite 0.00463 (6) 0.00416 (4) 0.0827 (5) 1.035 (5)

Subsolar 0.00490 (16) 0.00432 (14) 0.0843 (8) –

Solar wind 0.00491 (7) 0.00416 (9) 0.0842 (3) 1.0405 (10)

ncertainties in the last digits are given in parentheses.eferences: [1] Busemann et al. (2000); [2] Huss and Lewis (1994b), renormalized followireilites; [5] Crabb and Anders (1981), South Oman. [6]Meshik et al. (2014), current best

otes: For a discussion of differences between bulk Sun and solar wind see Section 2.1. *

atios in end members (see text). Two possible compositions for P6 are listed based on diff s-process only 130Xe. Such a composition would be obtained by further extrapolatingines, involving assumptions about 136Xe/130Xe ratios in endmembers (see text). 124Xe anhese isotopes in Xe-G the theoretically expected value (≡0) is listed. As for Xe-N, see tex

82Kr in the G component are variable. Ranges are for grain size fractions of presolar

the protosolar D/H ratio (e.g., Wieler, 2002; and referencestherein).

2.3.2. NeonCompositions are also listed in Table 3, together with the mea-

sured composition of bulk solar wind and the inferred compositionin the Sun’s OCZ (Heber et al., 2009, 2012). The OCZ value is thatobtained by Heber et al. (2012) using the correlation approach,while 21Ne/22Ne has been calculated from their solar wind valueusing their fractionation value of 16‰/amu. The 20Ne/22Ne ratios,except for the most exotic ones, are shown in Fig. 3, togetherwith 36Ar/38Ar ratios. Neon isotopic compositions, together withthose in xenon, have turned out to be the most diagnostic ones foridentification of different components and determination of theirrelative abundances. Strongly standing out is the composition ofNe in the graphite (R) and in the SiC-G component, both almostpure 22Ne (Table 3; not shown in Fig. 3), also termed Ne-E (Black,1972). In early work (e.g., Alaerts et al., 1980; Swindle, 1988), theNe-E subcomponents had been given the names Ne-E(L) and Ne-E(H), because the early separation attempts had shown them to belocated in two different carrier phases: one of low density and lowrelease temperature (L; now known to be graphite) and another ofhigh density and high release temperature (H; now known to besilicon carbide).


Strictly speaking, only very low upper limits on the abundanceof 20Ne and 21Ne in the R and G components have been obtained(Table 3). But it is widely accepted that Ne-R [=Ne-E(L)] is (almost)pure 22Ne, dominantly, but maybe not exclusively (Amari et al.,

he composition of solar wind Xe based on Genesis data (Meshik et al., 2014).

130Xe/132Xe 131Xe/132Xe 134Xe/132Xe 136Xe/132Xe References

0.1619 (3) 0.8185 (9) 0.3780 (11) 0.3164 (8) [1]0.1589 (2) 0.8247 (19) 0.3767 (10) ≡0.3096 [2]0.1542 (3) 0.8457 (13) 0.6356 (13) ≡0.6991 [2]0.1658 (11) 0.8229 (47 0.3288 (50) ≡0.3096 [2]0.1587 (3) 0.8370 (13) 0.5176 (13) ≡0.5493 [2]0.4826 (42) 0.1858 (117) 0.0222 (53) ≡0.00343 [3]0.1603 (13) 0.8109 (130) 0.4183 (59) 0.4006 (32) [3]0.1627 (5) 0.8195 (13) 0.3776 (12) 0.3152 (19) [4]0.1649 (10) 0.8301 (34) 0.3765 (25) 0.3095 (20) [5]0.1650 (4) 0.8256 (12) 0.3691 (7) 0.3001 (6) [6]

ng Busemann et al. (2000); [3] Lewis et al. (1994); [4] Göbel et al. (1978), avg. of 7estimate.Compositions derived from mixing lines, involving assumptions about 136Xe/132Xeferent assumptions. “Pure” nucleosynthetic Xe-HL is commonly assumed to be free

to 130Xe≡0 (see discussion in Section 3.2.3). x Compositions derived from mixingd 126Xe derived in this approach may be influenced by spallation contributions. Fort and reassessment of the 124Xe/126Xe ratio in Ott et al. (2005).




Fig. 6. Distribution of 132Xe and 22Ne among the various components in the Orgueil CI carbonaceous chondrite (Ott, 2002). Because in the most primitive meteorites virtuallyall (non-solar) trapped gases are contained in acid-resistant phases (e.g. Ott et al., 1981), where a better separation of components is achieved during stepwise heating, then nd xec bute m

1t“ct1fiTu

ateTog1tbetutp

imNtPatatc1

2

iapia

umbers shown are based on such residue data (Huss and Lewis, 1995): for neon aomponents on the etched residue data. Plotted are only components which contri

995; Amari, 2009) the decay product of 22Na (T½ = 2.6 a), which inurn was trapped when the grains condensed. It is this apparentlyradiogenic origin”, which led Amari et al. (1995) to re-name thisomponent Ne-R. The G component is thought to be material fromhe He burning shell of Red Giant AGB stars (e.g., Gallino et al.,990; see Section 3.2.1) and theoretical values for the small, butnite amounts of 20Ne and 21Ne accompanying 22Ne are listed inable 3 (in italics) together with the experimentally determinedpper limits.

Neon in the other planetary components in Table 3/Fig. 3 is char-cterized by higher 20Ne/22Ne ratios, although in all cases lowerhan in solar wind neon, occupying a range ∼8 to ∼11, whichncompasses the ratio observed in the terrestrial atmosphere (9.8).wo values are listed for Ne in the Q component because one groupf meteorites gives a 20Ne/22Ne ratio of ∼10.7, while another groupives a lower ratio of ∼10.1 (Busemann et al., 2000; Schelhaas et al.,990). Also for ureilite Ne, two values are listed: 10.70, which ishe average ratio for the ureilite carbon/diamond samples analyzedy Göbel et al. (1978), and a lower ratio of 10.4 reported by Ottt al. (1985a) for the Hajmah ureilite with very low spallation con-ributions. The similarity in isotopic composition between Q andreilite neon strengthens the case for a close relationship betweenhe two components as already evident in the elemental abundanceattern (Fig. 2; Section 2.2).

Neon in presolar diamonds (P3, HL, P6) dominates the neonnventory of primitive meteorites (Fig. 6; see Section 3), and for-

erly identified components Ne-A (Pepin, 1967), or Ne-A1 ande-A2 (e.g., Alaerts et al., 1980; Swindle, 1988) can be related to

hese newly identified primary components: A1 corresponds to3, A2 to HL+P6, and A is probably a mixture of all three plus, inddition, contributions from Ne-G and/or Ne-R. Since the P6 con-ribution is difficult to resolve, the value listed for HL in Table 3nd shown in Fig. 3 probably contains some P6 contribution. Inhis sense Ne-HL more closely corresponds to the old Ne-A2, butontributions from P6 in all likelihood are minor (Huss and Lewis,994b).

.3.3. ArgonCompositions for the various planetary components are listed

n Table 3, together with He and Ne. Again, measured solar wind


nd the inferred composition of the Sun’s OCZ are listed for com-arison. Given are only 36Ar/38Ar ratios. Measured 40Ar/36Ar ratios

n trapped-gas-rich samples are generally low (often <1) and actu-lly reported values critically depend on blank corrections and/or

non in the Q (P1) component on the HF/HCl residue data, for the Xe-P3 and Xe-HLore than 1% to the total.

purity of the analyzed samples so as to minimize radiogeniccontributions to 40Ar from K-bearing phases. Upper limits for the Q-component are (40Ar/36Ar)Q < 0.12 (Wieler et al., 1992; Busemannet al., 2000), for the P3 and HL components (40Ar/36Ar)P3 < 0.03 and(40Ar/36Ar)HL < 0.08 (Huss and Lewis, 1994b), respectively. The low-est value measured for 40Ar/36Ar is (2.9 ± 1.7) × 10−4, i.e., 6 ordersof magnitude lower than in the terrestrial atmosphere, and wasobtained on a diamond sample of the ureilite Dyalpur (Göbel et al.,1978).

Variations in 36Ar/38Ar are much less pronounced than in20Ne/22Ne (Fig. 3). The only component that is clearly distinct inFig. 3 is the HL component with 36Ar/38Ar = 4.41 ± 0.06 (Huss andLewis, 1994b). Noteworthy are also the similarities of 36Ar/38Ar inthe subsolar component and the solar wind, (again) in Q and ure-ilite noble gases, and in particular, the close agreement of thesewith both terrestrial atmosphere and the inferred OCZ value.

2.3.4. KryptonTraditionally, krypton is known to show the smallest differences

in isotopic composition among different solar system objects. Iso-topic compositions are listed in Table 4, along with the solar windcomposition reported by Meshik et al. (2014) from Genesis. Datafor the “more normal” of the planetary components (Q, P3, N, ure-ilite and subsolar) are shown in Fig. 4 as deviations (in per mill)from solar wind. Also included in Fig. 4 is the “normal variant” ofP6 (see below). The gray area shows the possible range estimatedfor the “true solar” composition of the Sun’s OCZ, as described inSection 2.1. Not shown in Fig. 4 are the 78Kr/84Kr ratios for the P3and P6 components (because of their large uncertainies) as well asfor the N component (no value reported). Again, as in the case of Ar,the subsolar component shows a marked similarity to solar wind.The others show very similar mass fractionation patterns relativeto solar wind, favoring the heavy isotopes by something like 0.8%per mass unit. Remarkably, they all more or less fall into the rangeestimated as possible for the true solar composition. The only clearexception is at mass 86, where relative to “solar” the N componentis depleted by about 7% and where P3 may contain a small excess.

Large anomalies are observed in the nucleosynthetic compo-nents carried by the presolar diamond, SiC and graphite grains(Table 4). Kr-G (s-process krypton) and Kr-N are carried by SiC


grains, where they occur in roughly equal abundance, and havebeen studied extensively (Sections 3.2.1 and 3.2.2). The composi-tions listed in Table 4 are those derived by Lewis et al. (1994) for SiCby partitioning the measured compositions with the assumption


ING ModelC

Erde x

oofsnt1Ttg

taoripmbafteo

2

twbhhLXotg

aefatacfmaolrapca

gmte1

K1

ic


U. Ott / Chemie der

f a 84Kr/82Kr ratio for the G component of 2.40, as based on the-retical grounds. The thus derived ratios 80Kr/84Kr and 86Kr/84Kror the G component are variable and constitute a sensitive mea-ure of physical conditions during their nucleosynthesis by the sloweutron capture process (s-process; Section 3.2.1). Similar kryp-on of s-process origin is present in presolar graphite (Amari et al.,995), where even more extreme variations have been observed.his may be due to the presence in varying abundance ratios ofwo discrete components, Kr-S(H) and Kr-S(L) carried by graphiterains of different density.

In contrast to the approach for G and N in silicon carbide, wherehe reported compositions are extrapolations, the values for Kr-P3nd Kr-HL in presolar diamond (Huss and Lewis, 1994b) are basedn the most extreme measured compositions, and Kr-HL of theeported composition is unlikely to correspond to Kr as producedn a nucleosynthetic event. Defining a composition for the P6 com-onent – which is released from diamonds at high temperature,ixed with HL gases – has turned out to be difficult. Two possi-

le compositions have been suggested by Huss and Lewis (1994b)nd are listed in Table 4: one assumes a “normal” 84Kr/82Kr ratioor P6 corresponding to the one assumed for the Kr-P3 component,he other (“exotic” P6) assumes a 84Kr/82Kr just beyond the mostxtreme measured data point in the Kr-HL/Kr-P6 mixtures. Detailsf the anomalies and their implications are discussed in Section 3.2.

.3.5. XenonBesides neon, xenon is the noble gas most diagnostic in its iso-

opic composition. An important role has been played by Xenon-HL,hich was the first of the nucleosynthetic isotope anomalies to

e discovered (Reynolds and Turner, 1964). The HL componentas received its name for the simultaneous overabundance of theeavy xenon isotopes (≡Xe-H) and the light xenon isotopes (≡Xe-). Because the H part originally was more reliably determined,enon-HL was first believed to be associated with fission, possiblyf a superheavy element (e.g., Anders et al., 1975), but in the endhe search for its host phase led to the discovery of the existence ofrains of presolar origin in primitive meteorites (Lewis et al., 1987).

Compositions of the planetary components are listed in Table 5,gain together with solar wind xenon (Meshik et al., 2014) as a ref-rence. This solar wind composition is not only based on Xe datarom the St. Louis group obtained by the authors themselves, butlso takes into account the results from the Manchester labora-ory (Crowther and Gilmour, 2013). Like in the case of Kr, Xe-HLs listed is based on mixing lines, assuming 136Xe/132Xe in the HLomponent to have a value of 0.70 (0.6991 after re-normalizationollowing Busemann et al., 2000; Ott, 2002), which is close to the

ost extreme measured value and the intersection of the P3-HLnd P6-HL mixing lines (Huss and Lewis, 1994b). To derive thether end composition of the mixing line, which is observed in theow-temperature release and corresponds to Xe-P3, a 136Xe/132Xeatio of ∼0.310, just below the lowest measured value has beenssumed (Table 5; Huss and Lewis, 1994b). Again, for P6, the com-onent with which Xe-HL mixes at higher release temperature, twoompositions are listed, based on two different (normal vs. exotic)ssumptions for 136Xe/132Xe.

The compositions of Xe-G and Xe-N present in presolar SiC andraphite grains – in analogy to Kr – have been derived based onixing lines and theoretical guidelines, i.e., assumption of essen-

ially no 136Xe (136Xe/130Xe = 0.0071) for the G component (Lewist al., 1994). For the N component, these authors chose the ratio36Xe/130Xe to be equal to 1.305 times that measured in the ureilite


enna. This leads (by definition) to excess 136Xe (and also to excess34Xe), but a Kenna-like (ureilite-like) composition for the lightersotopes in Xe-N (Fig. 5). A shortcoming of this treatment – when itomes to 124Xe and 126Xe – is that cosmogenic contributions were

PRESSxx (2014) xxx–xxx 9

ignored. But there must have been such contributions producedduring exposure to the cosmic radiation while the SiC grains residedin the interstellar medium For this reason, for 124Xe and 126Xe in theG component, instead of the Lewis et al. (1994) values, the theoret-ically expected values of zero are given in Table 5. Still, cosmic rayexposure ages for SiC appear to be short based on the amount of cos-mogenic Ne in “Jumbo grains” (Heck et al., 2009a; see also Section3.2.1), and also for Xe spallation contributions appear to be small asargued by Ott et al. (2005). From a re-assessment of the Lewis et al.(1994) data, these authors also argue for the 124Xe/126Xe ratio inXe-N to be ∼1.07 rather than the value ∼1.32 of Lewis et al. (1994)listed in Table 5 and plotted in Fig. 5. Using the revised 124Xe/126Xeratio value together with the nominal 124Xe abundance will removethe strong excursion of 126Xe into the negative range in Fig. 5.

As for Kr, the nucleosynthetic components in Xe will be dis-cussed in more detail below (Section 3.2). The more “normal”planetary Xe compositions (Q, P3, N, ureilite, subsolar and “nor-mal P6”) are shown in Fig. 5. Again, as in Kr, a general characteristicis an overall relative depletion of the light isotopes relative to theheavy ones, when compared to solar wind. At the lighter isotopesthere is some scatter, relative uncertainties are large there, how-ever. More interesting is that there seem to be differences not onlyin 136Xe/132Xe but also in 130Xe/132Xe (see Sections 3.2 and 6). Note-worthy again is the strong similarity between Q-Xe and that foundin ureilites, and the fact that Q, P3 and ureilite Xe – except for 129Xe– are in the range that may correspond to Xe in the “bulk Sun”derived as discussed in Section 2.1 and shown as the gray shadedarea. Another noteworthy feature is that except for 134Xe and 136Xe,the subsolar pattern is within (unfortunately rather large) errorsidentical with solar wind.

3. Chondrites: origins, host phases and history

The exact origin and the history for most of the trapped compo-nents discussed here still present a puzzle. This in particular holdsfor the Q component, which dominates the inventory of the heavynoble gases in the most primitive meteorites (Fig. 6). The Q com-ponent is likely to be representative for the most important noblegas reservoir in the Solar System outside of the Sun, and it is oftenthought to be derived from a gas of originally solar composition(e.g., Wieler, 1994; Ozima et al., 1998; Gilmour, 2010; Crowtherand Gilmour, 2013; Meshik et al., 2013, 2014). Somewhat clearerare our ideas regarding the components hosted by grains of presolarorigin, as a consequence of the unique nucleosynthetic informationrecorded in their isotopic signatures, which point to the stellar sitesfrom which they originate. “Approximately normal” componentssuch as Q may, to first order, be related to solar wind compositionby a mass-dependent fractionation process as suggested in Figs. 3–5(see also, e.g., Ozima et al., 1998), but detailed modeling of Kr andXe indicates that small additional variations may be present aris-ing from variable nuclear contributions (Sections 3.1, 3.2.3 and 6;Gilmour, 2010; Crowther and Gilmour, 2013; Meshik et al., 2013,2014). It should be kept in mind, though, that all the compositionaldata have limited accuracy and apparent differences should notbe over-interpreted, in particular when comparing datasets fromdifferent laboratories. A case in point is, e.g., the difference in Ne iso-topic ratios (outside reported analytical uncertainties) determinedfor the solar wind by the analysis of Genesis samples in differentlaboratories (Table 3; Heber et al., 2009, 2012; Pepin et al., 2012).

3.1. The Q component


For the heavy noble gases Ar, Kr and Xe, the Q component dom-inates the inventory of primitive meteorites. For example, in thecase of the chemically highly primitive Orgueil CI carbonaceous


IN PRESSG ModelC

1 Erde xxx (2014) xxx–xxx

c(“brpaRtrryvsi(c(mgQgi

cmmwp1bigaittabt2

∼

8

TlBrtmtfigaatsr

c(12pd

Fig. 7. Variation of 36Ar/132Xe and 84Kr/132Xe in samples from HF/HCl-resistantresidues from Allende and ALH 81032 (L3.4.) separated according to density/size.The dark red line shows the trend observed among Q in a variety of meteoritesas shown in Fig. 8 of Busemann et al. (2000). The variable ratios may possibly beexplained as due to different density/size distributions of the Q phase. Shifts due tothe presence of the HL-component (roughly at the left end of the line; Table 2) arenegligible (cf. Fig. 6). Data sources: (a) “Allende AI centr”: differential centrifuga-tion of Allende AI residues (Ott et al., 1981); (b) “81032 density”: density separatesfrom ALH 81032 residue (Haubold et al., 2012 and unpublished); (c) “Allende AII

by Marrocchi et al. (2005b) seem to point into a similar direction.


0 U. Ott / Chemie der

hondrite, about 95% of all 132Xe belongs to the Q(P1) componentFig. 6; Huss and Lewis, 1995; Huss et al., 1996). The definition ofQ gas” as a gas component is operational (Lewis et al., 1975): Theulk of planetary noble gases survives treatment with hydrochlo-ic and hydrofluoric acid (which leaves typically on the order of aercent of a meteorite), and during following treatment with nitriccid, essentially all of the “Q-gases” are lost, with little loss of mass.ecent work has shown that there are alternatives to the HF/HClreatment and that samples closely corresponding to the HF/HClesidues can also be obtained – relying on the small size of the car-ier grains – by purely physical means, however with much lowerield (Matsuda et al., 1999; Amari et al., 2003). After initial contro-ersies following its original discovery (Lewis et al., 1975), it noweems well established that Q is a carbonaceous phase. Convinc-ng evidence for this comes from the combined results of Ott et al.1981), who showed that noble gas abundances correlated with Content in acid residues of Allende, and the study of Frick and Pepin1981) who showed that the noble gas carrier is combustible. Pri-

arily because of their rather normal isotopic composition, it isenerally assumed (but in no way certain) that Q-gases (and hence) are of “local” or related, probably Solar System origin, unlike theases hosted by the presolar circumstellar phases being discussedn Section 3.2.

Within the Q component as operationally defined, variations inomposition have been observed, partly at least connected withetamorphic history (Huss et al., 1996; Busemann et al., 2000). Theost clear-cut trend is the decrease of the abundance of Q gasesith increasing metamorphic grade of the host meteorites, cou-led with an increase in median release temperature (Huss et al.,996). As far as isotopic compositions are concerned, there maye variations caused by re-trapping of other components (primar-

ly HL gases) by the Q phase in meteorites of higher metamorphicrade (Huss et al., 1996). However, this interpretation, which isn attempt to explain observations from stepwise heating exper-ments, is probably not unique. Less than “perfect behavior” ofhe various host phases under these conditions (i.e. the assump-ion of clear separation in stepwise pyrolysis or, in other cases, thessumption that only Q is attacked by the oxidative treatment) maye alternatives. The only clear-cut variation seen during closed sys-em etching (i.e., definitely in Q as defined) is in neon: ratios for0Ne/22Ne in Q from different meteorites group around two values,10.1 and ∼10.7 (Table 3; Busemann et al., 2000).

Also the elemental ratios vary: for example, 36Ar/132Xe and4Kr/132Xe ratios vary by about ±50% around the average listed inable 2, and even larger variations occur in the abundance ratio ofight to heavy elements (Huss et al., 1996; Busemann et al., 2000).usemann et al. (2000) suggest the existence of two types of subcar-iers “Q1” and “Q2” with slightly different chemical properties. Inheir interpretation variations in Ar/Xe and Kr/Xe reflect both ther-

al metamorphism and aqueous alteration of the carriers, whereashe relatively high He/Ar and Ne/Ar ratios in meteorites that suf-ered strong aqueous alteration indicate that the subcarriers differn how susceptible they are to aqueous alteration. From an investi-ation of density separates from the L3.4 chondrite ALHA 81032 and

survey of previous data on Allende obtained by Ott et al. (1981)nd Amari et al. (2003), Haubold et al. (2012) have suggested thathe subcarriers may differ in density and/or grain size (Fig. 7). Aimilar trend may be present in colloidal separates from an acidesidue of L4 Saratov studied by Amari et al. (2013).

In some way the most puzzling property of Q is the apparent highoncentration of noble gases. There have been a variety of attemptse.g., Yang and Anders, 1982; Wacker et al., 1985; Zadnik et al.,985; Nichols et al., 1992; Sandford et al., 1998; Marrocchi et al.,


005a, 2011) to simulate gas trapping by Q. But in no case, with theossible exception of Xe in Marrocchi et al. (2011), have inferredistribution coefficients approached those inferred from the gas

density”: density separates from Allende residue AII (Ott et al., 1981); (d) “AllendeAmari density”: density separates of Amari et al. (2003). (e) average Q and trend:Busemann et al. (2000).

abundance in Q and pressures expected in the solar nebula duringtrapping. And even if gas concentrations were extremely high in thestudy of Marrocchi et al. (2011), the trapped gases were in no wayas retentively sited as Q-gases. A widely acknowledged model inorder to explain the Q phenomenon is that of Wacker (1989), whosuggests that Q gases are physically adsorbed on interior surfacesformed by a pore labyrinth within amorphous carbon. Adsorp-tion/desorption times in this model are controlled by choke pointsthat restrict the movement of noble gases within the labyrinth.Another interesting mechanism (Hohenberg et al., 2002) involvesimpingement of low energy noble gases onto a growing surfaceaccompanied by the formation of chemical bonds, under condi-tions, where chemically more active elements are present in lowabundance only. In order for the process to work, the residencetime of the weakly bound noble gas atoms at the surface must belong enough so that they can be buried by the growing surface andthus retained (“active capture and anomalous adsorption”).

While the gas concentration must be high in the Q phase, it isquite uncertain how high it really is; hence in most cases concen-trations of Q gases are reported for HF/HCl residues only and not forthe Q phase proper (e.g., Huss et al., 1996; Busemann et al., 2000).Weight loss during chemical treatment with nitric acid (the proce-dure releasing Q gases and presumably dissolving Q) is small andgenerally difficult to determine, which surely causes some of thevariability in weight loss observed in different experiments. Mostimportantly, we lack detailed knowledge of the chemical processesgoing on during the etching. In particular, (a) it is highly likely thatother phases not hosting Q gases are dissolved at the same time(which will lead to an underestimate of the true concentration in Q)and (b) instead of being due to dissolution of the carrier, gas releasemay also be caused by structural rearrangement of the phase host-ing the noble gases, not being accompanied by actual physical loss(leading to an overestimate of the concentration).

Pointing toward the importance of the second possibility areresults from a recent Raman spectroscopic study by Matsuda et al.(2010a), who found that “the oxidization not only dissolved andremoved oxidized carbon, but also changed the carbon structureitself to a more amorphous (disordered) state”. Results of a study


These authors studied the effect of solvation by pyridine, a sub-stance known to cause swelling of organic matter, on a HF/HClresistant residue from Orgueil, and found that the heavy rare gases


ARTICLE ING ModelCHEMER-25295; No. of Pages 26

U. Ott / Chemie der Erde x

Fig. 8. Xe isotopic compositions in Q as well as in the P3 component in presolardiamonds (Table 5), enlarged from Fig. 5. Ratios are normalized to 132Xe, and shownare deviations in per mill of the iXe/132Xe ratios from the corresponding ratios mea-sured for the solar wind (Table 5; Meshik et al., 2014). The points are slightly shiftedalong the mass axis to aid in visibility. The gray shaded area is the possible rangeestimated for Xe in the Sun’s outer convective zone (OCZ), as discussed in Section 2.1.The lines show the effect of assumed mass-dependent fractionations of 10.8‰/amu(Gilmour, 2010), 13‰/amu (Marrocchi et al., 2011) and 8‰/amu (most likely accord-ing to the most recent evaluations; Meshik et al., 2013, 2014; Crowther and Gilmour,2013), respectively. Obviously Xe-Q contains an excess of radiogenic 129Xe from 129Idecay. The small enhancements at the lightest and the heaviest isotopes relative tothe Gilmour line can be explained by a modest addition of Xe-HL (Gilmour, 2010;Meshik et al., 2013, 2014; Crowther and Gilmour, 2013). The negative excursion atmG

wfiloeCiotptpoXnpn

lnebafifafwttai121s1t

(>90%) of SiC grains probably come from this stellar source (e.g.,

ass 128 may indicate a deficit of s-process Xe in P3 (Gilmour, 2010; Crowther andilmour, 2013).

ere largely lost. However, their finding was only partially con-rmed by Matsuda et al. (2010b), who found significantly smaller

osses in the case of Orgueil, and essentially no change in the casef an Allende residue. To complicate matters further, no significantffects were found in a study by Busemann et al. (2008) of variousM and CR meteorites, nor by Spring et al. (2011), who studied var-

ous primitive meteorites, notably also including Orgueil. On thether hand, recent results by Amari et al. (2013) are interpreted byhe authors as more indicative of the first possibility, i.e. the hosthase being a minor carbon residue fraction. In their study of a par-icularly gas-rich sample from Saratov (L4) these authors identifiedorous carbon “consisting of domains with short-range graphemerder” as the main constituent. Based on the relative abundances ofe and C, they then go on to conclude that most likely “individualoble gas atoms are associated with only a minor component of theorous carbon, possibly one or more specific arrangements of theanoparticulate graphene”.

Ionization energies may have played an important role in estab-ishing the fractionated elemental abundance pattern in Q-typeoble gases (Figs. 1 and 2; Ott, 2002), like in ureilites (e.g., Webert al., 1971, 1976; Göbel et al., 1978). This points toward the possi-ility that both originated in a similar process involving ionizationnd ion implantation (e.g., Rai et al., 2003a), a process also favoredor Q by Matsuda et al. (2010a) based on their Raman and TEM stud-es. Such reasoning, of course, assumes that the source compositionor the Q gases was solar. A clue may come from the isotopes. For Krnd Xe, to first order, Q type noble gases show a mass-dependentractionation relationship relative to solar wind (Figs. 4 and 5),hich for Xe is shown in expanded form in Fig. 8. The question

hen is how and when this fractionation took place and which ofhe two compositions is the original one. The common assumptionppears to be that the Q noble gases were mass fractionated start-ng from a gas of solar wind composition (e.g., Lavielle and Marti,992; Wieler, 1994; Ozima et al., 1998; Crowther and Gilmour,013; Meshik et al., 2013, 2014). A mass fractionation of about%/amu, as observed for Xe (Figs. 5 and 8), has, in fact, been found in


everal ion implantation simulation experiments (e.g., Frick et al.,979; Bernatowicz and Fahey, 1986; Koscheev et al., 2001), andhus may be consistent with the “plasma model” favored for Q


noble gases by, e.g., Matsuda et al. (2010a). However, the degreeof mass fractionation (if there is) probably depends on the setup ofthe experiments (Bernatowicz and Fahey, 1986), which for the sim-ulation experiments may well differ from the conditions prevalentduring trapping of Q gases in nature.

A different situation is suggested by the fact that the solar wind36Ar/38Ar ratio after correction for mass-dependent fractionationduring acceleration, i.e., the derived “solar” 36Ar/38Ar (from theOCZ; Section 2.1; Heber et al., 2012) matches within uncertaintiesthat of Ar in Q (Table 3 and Fig. 3). Along the same lines – again ina first approximation – the composition of Kr and Xe in Q falls wellinto the range of solar Kr and Xe covered by reasonable choices forthe fractionation between the Sun and solar wind. Thus it appearspossible, and maybe even likely, that Q-Ar, -Kr- and -Xe simply aresolar Ar, Kr and Xe (or at least very close), a possibility raised for thecase of Kr and Xe by Gilmour (2010) already. Some small additionsseem to be required, though (see Fig. 8). Following the model-ing of Gilmour (2010), Crowther and Gilmour (2013) and Meshiket al. (2013, 2014), in case of Xe this would be addition to Xe-Q ofradiogenic 129Xe from 129I decay and some Xe-HL (which may bedue to re-trapping during metamorphism or incomplete separationfrom Xe-HL carrying nanodiamonds; Table 5; Section 3.2.3). Theamount of Xe-HL required in these models is about 1.5% (at 132Xe)and the mass fractionation between Q and solar wind Xe is about0.8%/amu, well within the range considered reasonable for fraction-ation between bulk solar and solar wind Xe from the application ofthe ICD model (Section 2.1; and Figs. 5 and 8). In addition there maybe a weak hint for a tiny (∼0.003%?) contribution from s-process Xe(Meshik et al., 2013, 2014; but see Crowther and Gilmour, 2013).

3.2. Noble gases in pre-solar grains

Unlike the case of Ar, Kr, Xe, where Q is the dominant noblegas component, gases in presolar phases make a significant contri-bution to the planetary He and Ne budget in primitive meteorites(Fig. 6). In particular, presolar nanodiamonds are the dominant hostphase for planetary Ne, while 22Ne-rich Ne-G (Table 3) makes a verynoticeable contribution to the 22Ne budget (Fig. 6). As for the heavynoble gases, as discussed below, they show distinct isotopic com-positions reflecting the action of the nucleosynthesis processes bywhich they were produced: the slow neutron capture process (s-process) in the case of silicon carbide and graphite, some variantsof the rapid neutron capture process (r-process) and the p-process(by which the proton-rich isotopes are produced), respectively, inthe case of the nanodiamonds. Kr and Xe isotopic compositions forthe exotic components carried by silicon carbide (G, N) and dia-mond (HL) are plotted in Figs. 9 and 10 in an analogous fashion asthe more normal components in Figs. 4 and 5.

3.2.1. Silicon carbideThe G and N components of noble gases are carried by preso-

lar grains of silicon carbide. Diagnostic as far as their origins areconcerned, are the isotopic patterns observed in the G componentfor Ne, Kr and Xe: Ne-G [=Ne-E(H)] is dominated by 22Ne, withonly small amounts of 20,21Ne; while Kr-G and Xe-G, like manyother heavy trace elements in SiC (e.g., Yin et al., 2006), show largeoverabundances of the isotopes made in the slow neutron captureprocess (s-process) of nucleosynthesis (Kr-S, Xe-S; Tables 4 and 5;Figs. 9 and 10). This clearly points to an origin from material madein the He-burning shell of carbon-rich Red Giant stars in their AGBphase (Gallino et al., 1990; Lewis et al., 1990, 1994). It is also con-sistent with observations in other elements, and the large majority


Hoppe and Ott, 1997; Zinner, 1998; Hoppe, 2008).Krypton has proven to be a diagnostic for the conditions dur-

ing s-process nucleosynthesis (e.g., Ott et al., 1988; Gallino et al.,




Fig. 9. Kr isotopic compositions of “exotic” trapped components: HL and the “exotic”P6 version in presolar diamonds plus G in presolar SiC (Table 4). The points areslightly shifted along the mass axis to aid in visibility. Ratios are normalized to 84Kr,and shown are deviations in per mill of the iKr/84Kr ratios from the correspondingratios measured for the solar wind (Table 4; Meshik et al., 2014). Not shown arethe highly uncertain 78Kr/84Kr ratios in the “exotic” P6 and G components. Also notshown are the 80Kr/84Kr and 86Kr/84Kr ratios for the G component, which are variable(aa

1dapioH8

ba8

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s

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Fig. 11. Three-isotope plot of 86Kr/82Kr vs. 83Kr/82Kr in “bulk SiC” and grain sizeseparates from Murchison; data from Lewis et al. (1994). The lines connect the Kr-Ncomposition (right) with an end point on the left having 83Kr/82Kr as expected for

Ott et al., 1988; Lewis et al., 1990, 1994; see also text and Fig. 11). The gray shadedrea is the possible range estimated for Kr in the Sun’s outer convective zone (OCZ),s discussed in Section 2.1.

990; Pignatari et al., 2006). This is because of “branchings”, i.e.,ifferent possible pathways, during the s-process. For example,t low neutron fluxes 85Kr (with a half-life of 10.8 years) com-letely decays before it can capture another neutron and thus 86Kr

s bypassed; while at higher neutron densities a significant numberf neutrons is captured by 85Kr leading to production of stable 86Kr.ence the effective neutron density is recorded in the observed

6Kr/84Kr ratio. Similarly there is a branching at unstable 79Se, andecause the half-life of this nuclide is sensitive to stellar temper-ture, both neutron density and temperature are recorded in the0Kr/84Kr ratio. The sensitive Kr isotopic ratios vary as a function ofize of the SiC grains (Fig. 11; Lewis et al., 1994). These variationsccur coupled together with variations in the elemental composi-ion, in particular the abundance ratio Ne/Xe in the G componentNe-E(H)/Xe-S; Lewis et al., 1994; cf. also Russell et al., 1997; Ott


nd Merchel, 2000).Unlike for Ne (see below) and some more refractory elements,

uch as Zr, Mo, Ba (Nicolussi et al., 1997, 1998; Savina et al., 2003;

ig. 10. Xe isotopic compositions in “exotic” components: HL and the “exotic” P6ersion in presolar diamonds plus G in presolar SiC (Table 5). Ratios are normalizedo 132Xe, and shown are deviations in per mill of the iXe/132Xe ratios from the cor-esponding ratios measured for the solar wind (Table 5; Meshik et al., 2014). Notehat “pure” nucleosynthetic Xe-HL is commonly assumed to be free of s-process only30Xe. Such a composition would be obtained by further extrapolating to 130Xe≡0see discussion in Section 3.2.3). The points are slightly shifted along the mass axiso aid in visibility, while the gray shaded area is the possible range estimated for Xen the Sun’s outer convective zone (OCZ), as discussed in Section 2.1.

the G component (Table 4). The data points fall on separate mixing lines, indicatingvariable 86Kr/82Kr in the G component (Ott et al., 1988; Lewis et al., 1990, 1994).

Marhas et al., 2007), no telling single grain analyses have beenpossible so far for Kr and Xe in SiC. The only reported attempt isthat of Crowther et al. (2006, 2008), who analyzed 43 grains andfound nothing above their detection limit of 950 atoms for 132Xe.The trends with grain size, however, are opposite in the noble gasisotopic patterns to what would be expected from the trends inneighboring elements like Ba (Prombo et al., 1993; Gallino et al.,1993; Lewis et al., 1994). In addition, the noble gas G componentappears little/not elementally fractionated relative to its stellarsource, while the N component is fractionated [especially withinthe heavy noble gases and in the ratio of the heavy noble gasesrelative to He and Ne (cf. Fig. 2; Lewis et al., 1990, 1994)]. Takentogether, these observations have led Lewis et al. (1990) to suggesta scenario, in which the SiC grains condensed in the expandingenvelopes of AGB stars and were then impregnated with noblegas ions from a stellar wind. Ion implantation would thus domi-nate the noble gas pattern, while chemically active elements suchas Ba would have also been taken up during condensation andtheir inventory would be dominated by this source. The scenariois consistent with the elemental abundance ratios that point to a>104 × higher Ba/Xe ratio in the grains as compared to the source.To also accommodate the elemental trends within the noble gases,two wind components might be required: a minor component froma fully ionized region, contributing most of He, Ne and Ar, and acomponent from a cooler, only partially ionized region, relativelyenriched in Kr and Xe because of their lower ionization potentials.For a quantitative discussion and some alternative ion implantationscenario see Verchovsky et al. (2003, 2004), who suggest implanta-tion of a low-energy component during the main stage of AGB starevolution, with a higher-energy component implanted later duringthe stage of planetary nebula formation. However, while some frac-tion of the more refractory elements will also have been implanted(cf. Marhas et al., 2007), it is clear that co-condensation must haveplayed the major role for, e.g., the Rare Earth Elements (REE), whichin their elemental abundance pattern show a volatility trend, withYb strongly and Sm and Eu moderately underabundant relative tothe other REE (Yin et al., 2006).

There are overall trends with grain size, e.g., in purity of Ne-G (Lewis et al., 1994), the isotopic composition of Kr-G (Lewiset al., 1994; similar: Ott et al., 1988) and the Ne/Xe ratio of theG component (Lewis et al., 1994; Russell et al., 1997). However,the distribution among and concentration in individual grains isuneven, at least for He and Ne, the elements for which singlegrain measurements were analytically possible. Early work indi-


cated that Ne-G was concentrated in a small fraction of SiC grains(Nichols et al., 1993), but subsequent higher-sensitivity work byHeck et al. (2007) found about 40% of SiC grains with Ne-G contents



U. Ott / Chemie der Erde x

Fig. 12. Presolar cosmic ray exposure ages for “Jumbo” SiC grains from cosmogenic2

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1Ne. Data from Ott et al. (2009). Most analyzed grains were mainstream grainsfrom AGB stars), with a few AB or possible AB grains (characterized by very low2C/13C < 20).

bove detection limit. Concentrations ranged up to 1.5 × 10−2 cm3

TP/g for grains from Murchison and up to 2.7 × 10−2 cm3 STP/gor a grain from Murray. These authors could also confirm that inhe grains richest in 22Ne, the Ne-G component was always accom-anied by 4He, as predicted by astrophysical models (Heck et al.,007).

During residence in the interstellar medium, presolar grainsere exposed to the cosmic radiation. Cosmogenic noble gases pro-uced during this time are also “presolar” and therefore qualify as

subject of this review, besides those gases that were “trapped”.s with the cosmogenic noble gases produced during a meteoroid’s

ourney from parent body to Earth, this allows to calculate a cos-ic ray exposure (CRE) age – in this case a presolar CRE age. Of the

re-solar minerals which are subject of this review (SiC, graphite,anodiamond) only SiC can be reasonably used for this purpose,ecause it contains a useful target element (Si) on which cosmo-enic Ne can be produced (only easily lost cosmogenic helium wille made on graphite and diamond, since these consist of carbonnly).

A complication arises from the fact that, from grains of ∼�m-sizer smaller, a significant fraction of cosmogenic produced nuclidesill be immediately lost due to the recoil momentum received dur-

ng the spallation reaction. This essentially (also) rules out the use ofosmogenic He and also the smaller size fractions in case of cosmo-enic neon (Ott and Begemann, 2000; Ott et al., 2009; Trappitschnd Leya, 2013). The problem has been overcome by Heck et al.2009a) and Ott et al. (2009) by the use of very large “Jumbo” grains,ith sizes of tens of �m, where Ne recoil loss is small. Surprisingly

hort 21Ne CRE ages – in comparison to expected lifetimes of inter-tellar grains – were found by these authors, mostly less than 200a and many below detection limit (Fig. 12). Apart from the fact

hat we know little about flux and spectra of cosmic ray particlesn the interstellar medium >4.6 Ga ago, resulting in uncertain cos-

ic ray production rates, there are two additional caveats to thisnding. For one, the very large “Jumbo” grains that were analyzeday not be typical for the majority of SiC grains, e.g., abundances

f Ne-G are rather low. And second, CRE ages determined via cos-ogenic (?) 6Li on a set of similarly large grains were not as low as

he 21Ne ages, with the reason for the difference not understood.

.2.2. GraphiteIn contrast to silicon carbide, which is dominated by grains orig-


nating from a single source, namely AGB stars, graphite grains haveoticeable contributions from various sources, where relative con-ributions correlate with density (e.g., Zinner, 1998; Hoppe, 2008):ow-density graphites appear to be primarily of supernova origin,


while the denser types’ origins are more mixed (Jadhav et al., 2013).The AGB contributions are noticeable in the noble gases (Amariet al., 1995) by way of the signature of the s-process in Kr and Xeisotopes. Kr isotopic ratios that are sensitive to physical conditionsduring the process display an even wider range of values as in caseof SiC. For example, (86Kr/82Kr)G from graphite ranges from <0.5up to ∼5 (Amari et al., 1995), as compared to the maximum of∼3 observed in SiC (Fig. 11; Lewis et al., 1994) and the distribu-tion appears bimodal rather than continuous. In addition, the mostconspicuous feature in graphite noble gases is the “R component”,previously Ne-E(L), which appears to be monoisotopic 22Ne fromthe decay of 22Na (Amari et al., 1995; Amari, 2006), with a half-lifeof 2.6 a.

Results from single grain analyses by Nichols et al. (1994) indi-cated that, like Ne-G in SiC, Ne-R in graphite is concentrated in asmall fraction of the graphite grains. Unlike the case of SiC, however,where the fraction of grains with detectable Ne increased remark-ably when analyzed with improved detection limits (Heck et al.,2007), this has not been the case for graphite, where Heck et al.(2009b) found ∼22% (11 out of 51) graphite grains with detectableNe-R. In line with the multiple origins of graphite – apparentlyreflected also in density/morphology – the abundance of gas-richgrains also seems to depend on grain density (Nichols et al., 1994;Amari, 2006). What the analyses with improved precision have alsoshown, however, is that not all 22Ne-rich grains have necessarilymono-isotopic radiogenic 22Ne-R. Meier et al. (2012) analyzed 91grains from Murchison KFC1, a high-density separate with multi-ple stellar sources, and found several grains with finite amountsof 20Ne as well as one grain with detectable 21Ne, plus in additiontwo grains with 4He, but no detectable 22Ne. Clearly the situationis more complex than indicated by the initial analyses of “bulk”graphite samples or separates according to size and density (Amariet al., 1995; Amari, 2009).

As for the trapping of graphite gases, given the evidence for ionimplantation as an important trapping mechanism in the case ofpresolar diamond (see below) and silicon carbide grains, it appearslikely that this applies also to graphite (save probably the radiogenicNe-R). Elemental abundance patterns for Ar, Kr, Xe in graphite aresimilar to those in SiC, but Ne is low by one to two orders of mag-nitude, for which Amari et al. (1995) consider diffusive loss as themost likely explanation. This is in line with results for nitrogen iso-topes in individual presolar graphite grains, which appear to havebeen affected by exchange with isotopically “normal” nitrogen (e.g.,Hoppe et al., 1995; Jadhav et al., 2013).

3.2.3. Nanodiamonds (with glassy carbon)Nanodiamonds in primitive meteorites occur in high abundance

exceeding 1000 ppm (e.g., Huss and Lewis, 1995). Their average sizedetermined from TEM observations (Daulton et al., 1996) is about2.6 nm (roughly 10,000 carbon atoms), which is also consistent withmass spectrometric “bulk size measurements” by MALDI (Lyon,2005) and LA/ionization-TOF-MS (Maul et al., 2005). As a conse-quence, except for maybe carbon (Heck et al., 2012; Lewis et al.,2012), there is no chance for useful single grain isotopic analysisfrom which to draw conclusions regarding the source.

Nevertheless, there are hints for a multiplicity of sources. Thenanodiamonds were first identified by Lewis et al. (1987) as animportant host phase of isotopically exotic Xe, which seemed totie them to a supernova. In addition, a small amount concentratedin the heaviest density fractions appears to come from AGB stars(Verchovsky et al., 2006), as inferred from a heavy carbon isotopecomposition and the presence of s-process Xe. This is similar to SiC


mainstream grains discussed in Section 3.2.1 and thus will not bediscussed in further detail here. In any case, simple arithmetic tellsthat only a small fraction of the diamonds can be the actual gas car-riers (e.g., about 1 out a million diamonds contains a Xe atom, as has


IN PRESSG ModelC

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Fig. 13. Comparison of thermal release of 132Xe from Orgueil nanodiamonds (Hussand Lewis, 1994b) with those from artificial nanodiamonds implanted with noblegas ions (UDD1-1) by Koscheev et al. (2001). Both show two release peaks. Shown isthe percent release per ◦C temperature interval in stepwise degassing (left axis).



een frequently pointed out) and that the majority is gas-free. Ofourse, the situation is far less extreme when regarding the ordersf magnitude more abundant helium (Table 2) instead of xenon.evertheless, together with the normal carbon isotopic composi-

ion (in bulk), the situation is suggestive of a Solar System origin forhe majority of nanodiamonds (e.g., Zinner, 1998; Verchovsky et al.,006; see Dai et al., 2002, for additional arguments). Interestingly,ecent work by Stroud et al. (2011) has revealed that nanodiamondeparates actually contain substantial amounts of glassy carbon, aact not recognized previously. Currently it is not clear what the rel-tive importance of glassy carbon is as a gas carrier – in addition tor maybe even instead of the diamonds proper. As a consequencend as noted in the introduction, the term “nanodiamonds” wille used here as representative for “real samples”, i.e., for separates

ncluding both nanodiamonds and glassy carbon.There are also a variety of pathways for the formation of nan-

diamonds, and so not all nanodiamonds may have formed theame way. Commonly associated with high-pressure processes,hey may as well be formed during chemical vapor deposition,

process favored by TEM structural analysis (e.g., Daulton et al.,996). On the other hand, based on their recent discovery of glassyarbon along with the nanodiamonds, Stroud et al. (2011) favor aupernova-shock induced conversion of both, from organic carbonn the interstellar medium. Along similar lines, Marks et al. (2012),ased on atomistic simulation evidence, suggest nonequilibrumormation in space involving carbon onions as initial condensate,ollowed by conversion to diamond due to an “energetic impact”.aman and cathodoluminescence spectroscopy have been triedo distinguish between high- and low-pressure origin, but resultsave not been conclusive (e.g., Gucsik et al., 2008, 2012).

The P3, HL and P6 noble gas components are the most prominentn the nanodiamonds, and among the noble gases carried by preso-ar phases in primitive meteorites, P3 and HL are by far the mostbundant. Remarkably, for the most primitive meteorites they evenake a significant contribution to the whole noble gas inventory:

or the heavy noble gases like Xe, whose abundance is dominatedy the Q component, diamonds contribute typically on the order of0%, while for He and Ne the gases carried by the diamonds evenominate the inventory (Fig. 6; Huss and Lewis, 1995). As a conse-uence, the original “planetary” Xe (AVCC = average carbonaceoushondrite Xe; e.g., Eugster et al., 1967) in reality is a mixture of Q-e and diamond Xe, while the composition of the “classical” Ne-Aomponent (20Ne/22Ne = 8.2 ± 0.4; Black, 1972), which is the Ne ofhe “planetary” component in primitive meteorites, is dominatedy Ne in the diamonds’ Ne-HL (=Ne-A2; Section 2.3.2).

Apart from the minor fraction of probable AGB star originVerchovsky et al., 2006), the noble gas inventory of the presolariamonds comprises a total of three isotopically distinct compo-ents: P3, with an approximately normal isotopic composition ofhe heavy noble gases Kr and Xe, the isotopically anomalous HLomponent, and the less abundant and less well-characterized P6omponent (Figs. 4, 5, 8–10). P3 is thermally significantly moreabile than HL and P6 (peak release of P3 in pyrolysis of Orgueil nan-diamonds is at ∼490 ◦C; Huss and Lewis, 1994a; cf. Fig. 13), and itsbundance relative to that of HL – as well as, on the higher temper-ture side, the ratio P6/HL – may be used to estimate metamorphicemperatures (Huss and Lewis, 1994a, 1995).

The most complete characterization of the isotopic compo-itions has been given by Huss and Lewis (1994b), and theompositions listed in Tables 4 and 5 are based on their data. For P3nd HL these are based on mixing lines defined by gases releasedt low temperature (<1235 ◦C), while compositions of gas released


t higher temperatures are interpreted as a mixture of HL and P6.he decomposition of the components involves assumptions aboutndmember compositions for one of the isotopic ratios in order toerive the full set, e.g., assumed 136Xe/132Xe = 0.3096 and 0.6991

Also shown is the mass fractionation of the 136Xe/132Xe relative to the startingcomposition (air) in the UDD implantation experiment (right axis).

in P3 and HL, respectively (renormalized from 0.310 and 0.700 fol-lowing Busemann et al., 2000). The value for HL corresponds to theintersection of the P3-HL and P6-HL mixing lines, but no definiteconstraints exist on the other end of those lines and therefore alsonot for the P6 composition. Huss and Lewis (1994b) give two setsof ratios, both listed in Table 5: one assuming P6 to be a “normalcomponent” (136Xe/132Xe = 0.3096 as in P3), and another assumingP6 to be “exotic”, with the ratio just slightly lower than in the tem-perature steps with the highest P6/HL ratio (136Xe/132Xe = 0.5493;Table 5).

From an astrophysical point of view the most interesting com-ponent is the HL component of likely supernova origin, in whichthe Xe isotopes produced solely by the p-process (124Xe, 126Xe; Xe-L) and those produced only in the r-process (134Xe, 136Xe; Xe-H),are strongly enriched relative to those of intermediate mass thathave contributions from the s-process (Fig. 10). Prominent amongthe puzzles surrounding the HL component is the close associa-tion of the excesses in the light and heavy Xe isotopes, and so farall reports of observed separations between Xe-H and Xe-L havenot been confirmed by later experiments trying to reproduce theresults (cf. Huss and Lewis, 1994a; Meshik et al., 2001).

Another puzzle is the connection between the P3 and HL compo-nents. The isotopic composition of P3-Xe is isotopically close to thatof solar Xe and Q-Xe (Figs. 5 and 8) and can, to first approximation,described as fractionated solar wind, with the degree of mass frac-tionation – remarkably – the same as that for Xe-Q. There are twoadditional distinctive features in Xe-P3 (Fig. 8). One is the enhanced129Xe abundance relative to solar, similar as for Q-Xe. The other isthat, relative to mass fractionated solar wind, P3 shows a deficitin the s-process only isotope 130Xe (Fig. 8). A possible explanation(Gilmour, 2010; Crowther and Gilmour, 2013) is that P3 is in fact a“presolar” component isolated from a reservoir that later evolvedto the solar composition, with P3 missing a small late addition ofs-process material (∼1% at 132Xe; Crowther and Gilmour, 2013).From the isotopic signature in Kr, Gilmour (2010) argues that thiss-process component had the signature not of the “main”, but the“weak” s-process. The Xe-P3 signature may also suggest that 129I,the parent of the excess 129Xe, may have been alive in the P3 com-ponent in the early solar system and that the isolation betweenP3 and bulk solar Xe took place less than 100 Ma before solar sys-tem formation (Gilmour, 2010; for a more extensive discussion of


interrelations see also Gilmour and Turner, 2007).There are indications that the noble gases were introduced into

the nanodiamonds by ion implantation. One argument can be basedon the fact that they do not contain excessive abundances of 129Xe


IN PRESSG ModelC


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Fig. 14. Results from various attempts of reproducing the “pure” Xe-H pattern (i.e.extrapolated to zero 130Xe). Besides the Xe-H pattern, predictions from the followingmodels are shown: (a) the neutron burst of Howard et al. (1992); (b) the neutronburst of Meyer et al. (2000); (c) r-process network calculations in the high entropywind scenario as described in Farouqi et al. (2010) and references therein. For thecalculations an electron density Ye = 0.45 was assumed, and an expansion velocity of7500 km/s. The EFTSI mass model (Pearson et al., 1996) was used and the resultingcomposition is that obtained for the entropy range 204–260 kB/baryon. As reportedin Ott et al. (2012). (d) The rapid separation model (Ott, 1996), assuming separation


U. Ott / Chemie der

s would have been expected, if incorporation was by trappinguring grain formation in an early stage after the explosion of aupernova. This is because at this time 129I (made by r-processucleosynthesis only) would have been alive, and because chem-

cally active elements should have been more effectively trappedhan chemically inert noble gases (Lewis and Anders, 1981). Subse-uent decay would have significantly increased the abundance of29Xe (Lewis and Anders, 1981). Other evidence for an ion implan-ation process comes from pyrolysis/combustion experiments onifferent grain size fractions of diamond (Verchovsky et al., 1998;ilmour et al., 2005). The process was studied in the laboratory byerchovsky et al. (2000) using Ar ions and by Koscheev et al. (1998,001), who implanted ∼1 keV ions (He, Ar, Kr, Xe) into syntheticanodiamonds of similar size as the natural ones. Both Verchovskyt al. (2000) and Koscheev et al. (2001) found a bimodal release sim-lar to that in the meteoritic diamonds (Fig. 13). This result, though,n particular the observed abundance ratio of low to high temper-ture component, depends sensitively on irradiation conditions, inarticular with respect to radiation dose (Verchovsky et al., 2000;or further discussion see also Fisenko et al., 2010).

Koscheev et al. (2001) infer from their results that the P3 com-onent, besides the evident low temperature part, should have also

high temperature part and that s-process-only isotopes such as30Xe observed in the (high-temperature) release and commonlyssigned to the HL component (e.g., Huss and Lewis, 1994b) mayctually belong to the high temperature part of P3. As a conse-uence, in order to obtain “pure HL”, the composition given inable 5 and shown in Fig. 10 would need to be modified by extrap-lating the mixing lines to a 130Xe-free composition (Koscheevt al., 2001; Huss et al., 2008). In addition, the isotopic fraction-tion observed in the high-temperature portion of ion-implantedases (Koscheev et al., 2001) needs to be taken into account, whichas been done by Huss et al. (2000, 2008). By doing so, because the

ntrinsic effects are very large, the inferred isotopic compositione-HL changes only slightly, but changes in case of the other nobleases are severe. The Kr-H composition becomes less extreme, ar-L composition becomes apparent (i.e., an excess rather than aeficit relative to solar at 78Kr and 80Kr), the isotope anomaly at Aroes largely away, and Ne and He show signs for a pre-irradiationy cosmic rays. While this makes for a seemingly coherent picture,

n detail it leaves some puzzles, like, e.g. the different distributionf P3 and HL among grain size separates (Gilmour et al., 2005) andhe existence of the (admittedly small) P6 component, i.e. trimodalather than bimodal release from the meteoritic diamonds.

An (almost) 130Xe-free composition for the HL component islso more in line with most astrophysical scenarios. Since relativexcesses in 124Xe and 126Xe are not identical and nor are the rel-tive excesses of the two r-only isotopes (134Xe and 136Xe), theucleosynthetic sources cannot be the p- and r-processes proper.ittle effort has been expanded on the origin of the light (L) partf HL-Xe since the early work of Heymann and Dziczkaniec (1979,980), and it would be desirable to resume work on the topic with

mproved nuclear physics and astrophysics input. The heavy (H)art has received more attention. A favored model is that of the neu-ron burst (i.e., neutron captures at neutron densities intermediateetween those of s- and r-process) during the explosion of a super-ova, acting on a He shell seed which experienced a weak s-processefore (Clayton, 1989; Howard et al., 1992). The model quite suc-essfully explains (using the right parameters) the composition ofo in supernova SiC (Type X) grains (Meyer et al., 2000), but it

oes considerably less well for the case of Xe-H (for which it wasnitially constructed). For this reason an alternative has been sug-


ested, which is a regular r-process (producing r-process nuclidesn solar system abundance ratios) augmented by separation of sta-le Xe isotopes from radioactive precursors on a short time scalehours) after the nucleosynthesis process (Ott, 1996), which itself

after 7750 s and an addition of 4% of unseparated Xe. Note also the sensitivity of,in particular of the 132Xe abundance in Xe-H, to mass fractionation during trapping(see detailed discussion in Ott et al., 2012).

takes place on a timescale of about a second. This does not only givea better match to Xe-H, but is also much better than the neutronburst in reproducing the corresponding pattern for Te (Richter et al.,1998), the only element besides Xe for which useful data exist. Thechallenge here is to find a scenario where such a separation processcould take place. In a third approach, modifications of the r-processin the high entropy wind (HEW) scenario for r-process nucleosyn-thesis are explored (Ott et al., 2012; Kratz et al., 2012) by varyingparameters like entropy, neutron richness and expansion veloc-ity. The current status is shown in Fig. 14, where ratios observedand predicted are shown, normalized to 136Xe. Obviously, all mod-els with the exception of the rapid separation have problems withreproducing the exceptionally low 132Xe/136Xe ratio, which, on theother hand, is rather sensitive to assumed isotopic fractionationduring trapping (see discussion in Ott et al., 2012).

3.3. The solubles in carbonaceous and ordinary chondrites

HF/HCl-resistant residues retain a large fraction of the planetarygases in the bulk meteorites including both the Q gases and thosehosted by presolar materials. This is especially so for Xe, where, e.g.,the residues investigated by Huss et al. (1996) “typically” accountfor >75% of Q-Xe in the bulk meteorites (Huss et al., 1996). In detail,however, the situation is more complex, and unfortunately the “sol-ubles”, i.e., those gases lost during the acid treatment, have receivedlittle attention after the early exploratory work following the Lewiset al. (1975) discovery of the Q phase (e.g., Alaerts et al., 1979a,b;Matsuda et al., 1980).

A closer look at the data reveals that in many cases, while thepercentage of Xe is as high as reported in Huss et al. (1996), itis much lower for Ar and Kr. The difference is most convincinglydemonstrated by a comparison of 36Ar/132Xe ratios in bulk andHF/HCl-residue samples (Fig. 15), since this does not depend onthe frequently uncertain weight yields of the chemical treatment.The ratios are often considerably higher in the bulk samples than inthe residues. In fact, from mass balance calculations, Alaerts et al.(1979b) derive 36Ar/132Xe ratios in solubles from C3O meteorites


that range up to ∼850 for Kainsaz, where about 90% of trapped36Ar is in the soluble fraction. These authors, as well as Alaertset al. (1979a) from work on LL-chondrites, also tentatively iden-tify Xe in the soluble fraction as having a solar isotopic pattern, but



16 U. Ott / Chemie der Erde x

Fig. 15. 36Ar/132Xe ratios in bulk meteorites vs. the same ratio in HF/HCl-resistantresidues. The dashed line corresponds to ratios being equal. Data from Alaerts et al.(b3

nasbt2

1vttdctft(t

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npi1edm(dchSitUbgpc2m1m1l

1979a,b), Matsuda et al. (1980) and Schelhaas et al. (1990). Slight corrections haveeen applied for cosmogenic Ar. Obviously there is a gas component with high6Ar/132Xe that is lost during acid treatment.

aturally the uncertainties in these calculations are large. Lookingt Ar instead of Xe might seem more promising, but even for Kain-az with the highest inferred 36Ar/132Xe ratio this is complicatedy the presence of cosmogenic Ar. Still taken at face value, a correc-ion to Kainsaz Ar based on the measured abundance of cosmogenic1Ne and a ratio (38Ar/21Ne)cos of ∼0.15 for CO chondrites (Eugster,988), increases the 36Ar/38Ar ratio by ∼2.4% to ∼5.42, close to thealues for solar wind and the subsolar component (Table 3). Overall,he observations point to a possible link between the solubles andhe subsolar component that is more clearly seen in enstatite chon-rites (see next section), which may thus be present in at least somearbonaceous and ordinary chondrites as well. This is an impor-ant and long neglected topic which deserves more attention inuture work. Studies of bulk meteorites using the in vacuo etchingechnique so far have been performed on enstatite chondrites onlyBusemann et al., 2003a,b), and their application to other meteoriteypes would be a step in that direction.

.4. Enstatite chondrites: subsolar and sub-Q noble gases

The noble gas components known from carbonaceous and ordi-ary chondrites, such as the Q component and the gases carried byresolar diamond and SiC (Sections 3.1 and 3.2), are also present

n the enstatite chondrites (e.g., Huss and Lewis, 1995; Huss et al.,996; Russell et al., 1997; Patzer and Schultz, 2002; Verchovskyt al., 2002), with indications for a lower matrix-normalized abun-ance of Q gases (similar to CV3 meteorites) as compared toeteorites like Orgeuil and unequilibrated ordinary chondrites

Huss et al., 1996). A distinctive feature of the enstatite chon-rites, however, is the occurrence of the “subsolar” noble gasomponent (Crabb and Anders, 1981), in particular in those ofigher petrologic type (Crabb and Anders, 1981, 1982; Patzer andchultz, 2002; Okazaki et al., 2010). Subsolar gases are character-zed by high Ar/Xe and Kr/Xe ratios when compared to the “typical”rapped noble gas pattern (excepting solar wind) as shown in Fig. 2.nlike Q, the major host phase of Ar/Kr/Xe in ordinary and car-onaceous chondrites (Fig. 6), the phase that hosts the subsolarases is largely soluble in HF/HCl (although apparently not com-letely so; Busemann et al., 2001b), and a non-negligible fractionan also be released by crushing the meteorites (Okazaki et al.,010). A favored candidate as the major host phase is the majorineral in enstatite chondrites, i.e., enstatite (Crabb and Anders,


981, 1982; Okazaki et al., 2010). Alternative or additional hostsay be phases associated with enstatite (Crabb and Anders, 1981,

982) or “friable phases” (Okazaki et al., 2010). Notably, subso-ar type gases have also been found in chondrules from enstatite

PRESSxx (2014) xxx–xxx

meteorites (Okazaki et al., 2001). Based on this observation, Okazakiet al. (2010) suggest a scenario in which solar noble gases (laterconverted to the subsolar pattern through metamorphism) wereimplanted into chondrule precursors, which then accreted andformed enstatite chondrites showing the subsolar signature. Note,however, although the traditional belief has been that chondrulesshould have lost all noble gases at the time they were melted(Okazaki et al., 2001; and references therein), the occurrence oftrapped noble gases in chondrules is not a unique feature ofenstatite chondrites, but has also been seen elsewhere (e.g., Vogelet al., 2004a; Beyersdorf-Kuis et al., 2013).

As discussed above (Section 3.3), the subsolar (or a similar)component appears to be present also in at least some carbona-ceous and ordinary chondrites, although this fact has achieved lessattention. Its occurrence is shown by the fact that Ar/Xe often is sub-stantially higher in bulk meteorites as compared to Q-dominatedacid-resistant residues (Fig. 15; Alaerts et al., 1979a,b; Schelhaaset al., 1990). Isotopically, subsolar gases appear similar to solarwind noble gases (Crabb and Anders, 1981; Busemann et al., 2001b),most clearly visible in the crushing steps of Okazaki et al. (2010).The “subsolar” may not be an independent component, though, aspointed out by Busemann et al. (2003a,b) based on a detailed etchstudy on St. Mark’s (EH5). According to these authors it may contain,besides solar-like noble gases also contributions from Q and a ter-restrial component. The mixing ratios appear to be variable amongE chondrites, however, since in the case of St. Mark’s (bulk) the20Ne/36Ar ratio of ∼0.04 (Busemann et al., 2003a) is about an orderof magnitude higher than that for South Oman (∼0.003; Table 2,Fig. 2; Crabb and Anders, 1981).

While subsolar gases show up mostly in enstatite chondritesof higher types, Patzer and Schultz (2002) report evidence in themore primitive E3 chondrites for still another component thatthey name sub-Q, with very low 36Ar/132Xe (23 ± 4) and 84Kr/132Xe(0.68 ± 0.34) abundance ratios. While these authors consider “sub-Q” a separate component, it is worthwhile to remember that thereis quite significant variation within the Q component proper (seeFig. 7; Busemann et al., 2000), so the so-called sub-Q gases mayjust be an extreme example of that. Another possible explanation, atleast in some cases, is indicated from the additional consideration ofisotopic compositions. Doing so Okazaki et al. (2010) infer that sub-Q present in desert meteorites is mostly elementally fractionatedair.

4. Non-chondritic meteorites

4.1. Ureilites

The presence of trapped noble gases in ureilites belongs to theproperties that make this class of achondrites “enigmatic”, as theyoften have been called because they combine properties of primi-tive and not-so-primitive meteorites. The high concentrations ofnoble gases that approach or in some cases even surpass thosefound in the most primitive chondrites (Mazor et al., 1970; Göbelet al., 1978; Rai et al., 2003a) is among the “primitive” ones. Thehigh abundances as well as the similarity to the Q component in iso-topic composition and elemental abundance pattern (Figs. 1–5) hasrepeatedly been cited as evidence for a close relationship betweenthese two components (Ott et al., 1984, 1985b; Busemann et al.,2000; Rai et al., 2003a). Also, like in the case of Q, the highly frac-tionated elemental pattern, which correlates with the respectiveionization energies (Weber et al., 1971, 1976) is suggestive for an


introduction of the noble gases by ion implantation.In most ureilites the noble gases are predominantly hosted by

diamonds (Göbel et al., 1978; Rai et al., 2003a), with evidencethat some may reside in a different carbonaceous phase (Göbel


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t al., 1978). The ureilite diamonds, unlike the nanodiamonds inrimitive chondrites discussed earlier, are definitely of solar sys-em, not of presolar, origin. A popular assumption is that theyere produced from graphitic carbon due to shock transformation

Lipschutz, 1964; Vdovykin, 1970). Graphite has been suggesteds the noble gas carrier in less shocked ureilites such as ALH78019Wacker, 1986), but some kind of “amorphous” carbon, maybe simi-ar to the Q phase seems more likely, given that the noble gases arelmost completely lost upon oxidizing chemical treatment usingitric or perchloric acid (Rai et al., 2002, 2003a). Whatever thease, this may point into the direction of a “shock scenario”, inhich trapped noble gases in all ureilites were originally carried

y graphite/amorphous carbon and were retained during shockransformation into diamond. Trapping of noble gases and diamondormation would thus be decoupled. A problem arises, however, forhis chain of events from the large difference in nitrogen isotopicomposition between diamond (�15N < −100‰; Yamamoto et al.,998; Rai et al., 2003b) and co-existing other carbon phases, thus

nebular origin for the diamonds may be indicated instead (Rait al., 2003b). Matsuda et al. (1991, 1995) studied both processesnd their conclusion is that the “vapor-growth model” for diamonds a better match to the observations than the shock productioncenario.

There are two more interesting observations concerning ureiliteoble gases: One is that a moderate percentage of the trapped ure-

lite gas can be released by crushing the meteorite sample (Okazakit al., 2003). What is remarkable here is that in case of ALH 78019,

diamond-free ureilite, 25% of total gas was released, with ele-ental ratios from crushing indistinguishable from the values for

he total; while in case of Kenna, a diamond-bearing ureilite, ther/Xe and Kr/Xe ratios for the crushing were substantially lower

han for the totals (Ar release during crushing: ∼5%, Xe: ∼19%).he other interesting case is that of the Antartic ureilite ALH 82130ith 36Ar/132Xe and 84Kr/132Xe ratios that are not only extremely

ow (∼20 and ∼0.1, respectively), but also deviate from the gen-ral (extrapolated) trend (Ott et al., 1986; Rai et al., 2003a). This isresently not understood, but must be telling us something abouthe origin and history of the ureilites.

.2. Other achondrites

Among the other achondrites, the HED meteorites, acapulcoites,rachinites and angrites are interesting cases. The acapulcoites and

odranites belong to the primitive achondrites, and so it is not toourprising that they were able to retain some planetary gases, asound, in the case of the acapulcoites, by e.g., Palme et al. (1981)nd Schultz et al. (1982). The case is different for the ureilites andhe brachinites (Bischoff, 2001; Krot et al., 2014; Mittlefehldt et al.,996, 2003), but nevertheless these also contain trapped nobleases of the planetary type. Noble gases in the ureilites, in partic-lar, occur in substantial concentrations rivaling those of the mostrimitive carbonaceous chondrites (for references and discussionee Section 4.1) While abundances are lower in the acapulcoites,hey are still surprisingly high, corresponding to those found inype 4 ordinary chondrites (Schultz et al., 1982). Abundances in therachinites are lower (Ott et al., 1985c, 1993; Weigel et al., 1997),ut elemental abundance ratios as well as Xe isotopes again showimilarities to the ureilites and to Q.

The HED meteorites (howardites, eucrites, diogenites) are theesult of magmatic differentiation processes on their parent body,ost likely Vesta (e.g., Drake, 2001; McSween et al., 2011). Anal-

ses of the diogenite Tatahouine (certain temperature steps) by


ichel and Eugster (1994) created much interest in the noble gasommunity. This is, because, like similar results for the more prim-tive achondrite Lodran (certain mineral separates) by Weigel andugster (1994), they appeared to provide direct observations of


a particularly primitive Xe component named U-Xe (e.g., Pepinet al., 1995; see Section 6) – hitherto hypothetical, except fora tentative hint in a sample from CM2 Murray (Niemeyer andZaikowski, 1980). Careful re-analysis by Busemann and Eugster(2002) did, however, not confirm the earlier reports. Instead, theBusemann and Eugster (2002) data for Lodran and several otherlodranites, the diogenite Tatahouine, the eucrite Pasamonte, fiveaubrites and two angrites are all fully compatible with mixtures ofQ-gas and adsorbed air, plus cosmogenic and fissiogenic contrib-utions (Busemann and Eugster, 2002). In particular, the isotopicdata strongly point to the presence of Q-Xe also in these achon-drites. The elemental patterns for Lodran and the aubrites, on theother hand, point toward the presence of subsolar gas in thesemeteorites. A further interesting observation is the presence ofsolar gas in glass from the angrite D’Orbigny, but not in bulk sam-ples of D’Orbigny and another angrite, Sahara 99555, which did notcontain any detectable trapped noble gas (Busemann et al., 2006).

5. Lesser components

Besides the components discussed in detail above, a number oflesser and/or less-well characterized components have been iden-tified, primarily for xenon. These cannot be exhaustively discussedhere; moreover, in several cases a confirmation of their existenceas a distinct component would be desirable. One of these, foundin sulfides of the Allende meteorite, is characterized by overabun-dances of 124Xe in gas fractions released in the 1400–1500 ◦C range,without corresponding overabundances in the other light isotopesas expected for Xe-L or spallation xenon (Lewis et al., 1979). Oth-ers are connected with iron meteorites, some of which containinclusions rich in trapped noble gases with compositions similar tothose of the “normal” components discussed above (Bogard et al.,1971; Niemeyer, 1979; Mathew and Begemann, 1995). Notewor-thy in addition is also the occurrence of solar-type noble gases insome iron and stony-iron meteorites (Becker and Pepin, 1984a;Mathew and Begemann, 1997). Only few specific examples arebriefly described below.

5.1. Xenon in chondritic metal

Marti et al. (1989) have identified a distinct xenon component(FVM-Xe) in a metal separate of the Forest Vale (H4) chondrite. It ischaracterized by relative abundances of the heaviest isotopes withunusually high 134Xe. A possible explanation is recoil of fission frag-ments into the metal grains, possibly from 244Pu, 248Cm or neutroninduced fission of 235U (Marti et al., 1989). This suggestion is con-sistent with the observed grain size dependence of FVM-Xe whichfavors location near the surface of the metal grains.

5.2. Xenon and krypton in silicate-graphite inclusions of the ElTaco iron meteorite

In their analysis of silicate-graphite inclusions of the El Taco IABiron meteorite, Mathew and Begemann (1995) found not only atrapped component of Ar, Kr, Xe in the silicates with Xe isotopi-cally close to ureilite xenon, but in addition a different componentin the high temperature release from graphite and schreibersite[(Fe,Ni)3P]. The reported Xe composition is unlike any of the well-established types of Xe and can be generated as a mixture of massfractionated U-Xe (the putative primary component of “planetary”Xe; Pepin et al., 1995; Pepin, 2000; Pepin and Porcelli, 2002; Sec-


tion 6) and 244Pu fission Xe. The fission component would havebeen added prior to incorporation of Xe into the El Taco parentbody, while the mass fractionation would have occurred after themixing of U-Xe and fission Xe.


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.3. Xenon in acid residues of iron meteorites

Murty et al. (1983) analyzed noble gases in acid-resistantesidues from the Canyon Diablo and Campo del Cielo (El Taco)eteorites. Xe in the nonmagnetic fractions was found to be air-

ike except for 124Xe and 126Xe, with the relative abundance ofhese isotopes even lower than in air (in which their abundancelready is the lowest among well-characterized solar system reser-oirs), by up to 15%. Murty et al. (1983) speculate that this may beresolar nebular Xe which was trapped in micro-inclusions of ironeteorites.

. Interrelations, hosts, fractionation and trapping

Having described the characteristic abundance patterns androperties of the primary planetary noble gas components, it

s useful at this stage to briefly sum up possible interrelations,hat is known about the carrier phases and ideas about trappingechanisms. For more details see the individual chapters and the

espective references.

.1. Interrelations

The easy cases are the subsolar and nucleosynthetic compo-ents (G, HL in “pure” form). It is all but certain that the subsolarases are related to solar wind (possibly with addition of Q anderrestrial air; Busemann et al., 2003a) as may be the componentith high Ar/Xe ratio contained in the acid-soluble fraction of

rdinary and carbonaceous meteorites. The nucleosynthetic com-onents, on the other hand, reflect nuclear processes in their parenttars, so an obvious relation to the “local” components cannote expected. Still, presolar silicon carbide and diamonds containlso more normal components, N and P3. For the N componentncertainties may be too large for a proper statement, but P3 iso-opically clearly resembles the solar composition (Figs. 3–5, 8).his is remarkable since it indicates that the relative contrib-tions from the various nucleosynthetic processes to P3 and theolar component must have been rather similar. As obvious inigs. 5 and 8, relative to mass fractionated solar wind, P3 shows

deficiency in the s-only isotope 130Xe. A possible explanationGilmour, 2010; Crowther and Gilmour, 2013) is that P3 is a preso-ar component that evolved in the same reservoir from whicholar Xe was derived, but both separated from each other “late”,ith P3 lacking a late contribution of weak s-process material.ased on the slight excess of 129Xe, Gilmour (2010) suggests theeparation even may have taken place ∼100 Ma before Solar Sys-em formation. A late introduction into the diamonds of P3 is, inact, in agreement with the suggestion of Koscheev et al. (2001)hat the “pure” HL component was introduced into the diamondsarlier than P3 and subsequently lost its less retentively sitedart, showing up only at higher release temperatures during step-ise heating analysis. The almost normal P3 component, whichas introduced later, would not have seen such a loss event,

herefore having also a low-temperature release in addition toigh temperature release where it dilutes the HL signature. For

further thorough discussion of structures in the presolar Xeomponents based on three-dimensional fits aided by theoreticalonsiderations, see Gilmour and Turner (2007). In particular andurprisingly, according to these authors the results are stronglyuggestive of a nucleosynthesis process that produces “r-process”sotopes not only without producing s-process isotopes (128Xe,


30Xe), but also “without producing the conventional r-process iso-ope 136Xe”.

Another relation suggested by the patterns in Figs. 4, 5 and 8 ishat of a mass fractionation relationship also between the isotopes


in Q and solar wind, with the addition to Q gas of some radiogenic129Xe, possibly a small amount of Xe-HL (∼1.5%; Gilmour, 2010;Crowther and Gilmour, 2013; Meshik et al., 2013, 2014) and maybea tiny amount of s-process Xe (∼0.003%; Meshik et al., 2013, 2014).Notably, the mass dependent effects appear quite similar in case ofP3 and Q. In this context it is useful to keep in mind that mass frac-tionation is also to be expected between solar wind and the bulksolar composition (see Section 2.1), and from the sign and the sizeof the differences the composition of Q gas (with the exception ofthe radiogenic 129Xe) may thus well be close to, if not identical to,the composition in the Sun. In this case, models to make Q gasesout of solar (wind) gases (e.g., Ozima et al., 1998; Pepin, 1991, 2003)may be obsolete. But, of course, this clearly needs a thorough under-standing of mass fractionation during acceleration of the solar wind(Section 2.1).

A Xe isotopic composition more primitive than solar or Q, namedU-Xe, has been derived from multi-dimensional fits to carbona-ceous chondrite stepwise heating data (Pepin and Phinney, 1978;see also Pepin, 1991; Pepin et al., 1995). U-Xe is identical to solarwind Xe for the light isotopes, but has lower abundances of the twoheaviest ones, 134Xe and 136Xe. The observed compositions duringstepwise heating of primitive chondrites are then thought to bederived from U-Xe by addition of Xe-H (Xe-G and Xe-N are muchlower in abundance and are ignored in this approach); while, inthis picture, compositions observed in achondritic meteorites ofthe HED type, following an earlier suggestion by Takaoka (1972),are thought to be derived by the addition to U-Xe of 244Pu-fissionXe.

If true, there are astonishing consequences. For one, solar Xe(as well as Q and P3, which appear closely related as discussedabove) would contain a surprisingly large extra fraction (∼ 8% at136Xe) of Xe-H. Furthermore, also puzzling appears that in primi-tive meteorites Xe-H and Xe-L, where they are carried by presolarnanodiamonds, both always appear to occur in constant abundanceratio, but that according to the evaluation by Pepin and Phinney(1978), Xe-H and Xe-L would need to be decoupled (Pepin, 2000;Pepin and Porcelli, 2002): extra Xe-H would be present in the solarwind, unaccompanied by Xe-L.

U-Xe may have been directly observed in a low temperaturerelease step from the non-colloidal part of an acid-resistant residuefrom the Murray meteorite (Niemeyer and Zaikowski, 1980). It hasalso been reported as being present in two achondritic meteorites(Tatahouine and Lodran; Michel and Eugster, 1994; Weigel andEugster, 1994), but these findings have been discredited by morerecent careful re-analysis (Busemann and Eugster, 2002), so thereality of U-Xe appears questionable at this time.

6.2. Hosts

Again, the presolar components are most straightforward. Inparticular, single grain analyses of silicon carbide and graphite(Heck et al., 2007, 2009b; Meier et al., 2012), have demonstratedthat a significant fraction of these grains are carriers of the respec-tive Ne components, and there is no apparent reason why thisshould be different for the heavy noble gases. Q is also well estab-lished as a carbonaceous phase (Ott et al., 1981), shows, however,some complexity (e.g., Busemann et al., 2000; Marrocchi et al.,2005b; Spring et al., 2011). These phases seem to be concentratedin the meteorite matrices (e.g., Huss and Lewis, 1995; Huss et al.,1996). Laser microprobe analyses of CM meteorites indicate that alarge fraction may reside in accretionary rims around chondrules


(Nakamura et al., 1999). Some gas is also present in the chondrulesproper (e.g., Vogel et al., 2004a; Beyersdorf-Kuis et al., 2013), butCa-Al-rich inclusions appear to be free of trapped gas (Vogel et al.,2004b).




Fig. 16. Elemental ratios of various components normalized to 132Xe and the solar abundance ratios (Table 2). Shown are data for average Q (Busemann et al., 2000), P3 andHL (Huss and Lewis, 1994a) as well as a ureilite pattern. To minimize interference of radiogenic 4He, for the ureilites a “representative” acid-resistant residue is shown rathert energy1 00 K

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han data for bulk (Novo Urei B-1; Göbel et al., 1978). (a) Left: plotted vs. ionization

952; see also Göbel et al., 1978), with plasma temperatures of 7000, 8000, and 90lmost straight line for HL (cf. also Ott et al., 1981).

The bulk of the subsolar component is most likely carried byhe major silicate phase of the enstatite meteorites, enstatite. If theases in the “solubles” in ordinary and carbonaceous meteoritesave a similar origin as the subsolar gases, it is very well possi-le that they are also predominantly hosted by the major silicateaterials, but this is less certain.Overall though, carbon phases seem to play a much more sig-

ificant role as host phases as one may expect simply from theirbundance. Chemical properties may have played a role there.or example, silicon carbide and graphite are major condensatesrom their parent stars at the time and location where noble gasesith the observed nucleosynthetic properties were synthesized

nd ejected. Chemical properties may also have played a role inorm of resistance to alteration. Diamonds, graphite, SiC and the Qhase are chemically highly resistant and simply may have (par-ially) survived, while others did not, a fact which is also reflectedn the generally surprisingly high abundance of presolar SiC rela-ive to presolar silicates (e.g., Floss and Stadermann, 2009; Leitnert al., 2012). On the other hand, refractory oxides such as corundumnd spinel are similarly chemically resistant as SiC, and the preso-ar oxide/presolar silicate ratio is being used to infer the degreef alteration of primitive meteorites (e.g., Bose et al., 2012; Leitnert al., 2012). So, based on chemical inertness alone one might expectlso these oxides to contain trapped noble gases. However, a care-ul study by Lewis and Srinivasan (1994) was not able to confirmhis expectation.

.3. Fractionation and trapping

One key to an understanding of the trapping mechanism is theigh (for a noble gas) abundance in the important host phases, inarticular the Q phase if it is indeed not the bulk of, but a minorhase within, the insoluble organic matter, a question not definitelyettled yet; another is the fractionation pattern which also may beelated to the trapping process. Once more, the case is easy for theubsolar and, if related, the high Ar/Xe component in the solubles

most likely originally implanted solar wind that was elementally


ractionated during metamorphic loss.As for the more typical planetary gases, while purely mass

ependent fractionation processes have been extensively consid-red (e.g. Rayleigh type; Ozima et al., 1998), they are not doing

. The dotted lines give expected values for the ionized fraction in a plasma (Elwert,(from left to right). (b) Right: plotted vs. difference in mass relative to 132Xe. Note

very well in reproducing both the observed elemental and isotopicfractionation patterns. Two major trapping processes have seri-ously been considered: adsorption (e.g., Fanale and Cannon, 1972;Bernatowicz and Podosek, 1986; Zadnik et al., 1985; Wacker, 1989;Huss et al., 1996) and ion implantation (e.g., Bernatowicz and Fahey,1986; Matsuda et al., 1991; Verchovsky et al., 2000, 2003, 2004;Koscheev et al., 2001). An interesting variant, which clearly needsfurther active study, the “active capture and anomalous adsorptionscenario” of Hohenberg et al. (2002), combines features of both.

What is a strong argument for the involvement of ions is theelemental fractionation pattern relative to the solar composition inthe case of the Q and ureilite components (e.g., Rai et al., 2003a),which follows a clear trend with ionization energy (Weber et al.,1971, 1976). This is demonstrated in Fig. 16a, which compares theelemental trends for several of our components with the degree ofionization as a function of plasma temperature. Surprisingly, whilethere is a good correlation in case of Q and ureilites, in case ofthe nanodiamonds, where the case for ion implantation is espe-cially strong from the laboratory studies (Verchovsky et al., 2000;Koscheev et al., 2001), this is not so, and the diamond element pat-tern is much better described by a mass dependent relationship(Ott et al., 1981; Fig. 16b). But then, of course, the solar abundancepattern may not be the appropriate reference for the HL component.

Interestingly, while equilibrium adsorption has no noticeableeffect on isotopic compositions (Bernatowicz and Podosek, 1986;Marrocchi and Marty, 2013), laboratory experiments studyingnoble gas trapping via ions has regularly resulted in isotopic frac-tionation, on the order of 1%/amu in the case of Xe (e.g., Bernatowiczand Fahey, 1986; Koscheev et al., 2001). This has also to be takeninto account when comparing the Q pattern with that of the solarwind and the Sun proper (in addition to the fractionation of thesolar wind; Section 2.1).

7. Related materials

Most meteorites come from the asteroidal belt (e.g., Bischoff,2001; Krot et al., 2014), but maybe some of the most primitive


ones have a connection to comets. It is thus instructive to com-pare observations on meteorites with such of bona fide asteroidaland cometary matter, which has become available for laboratorystudy through recent space missions Hayabusa and Stardust. Other


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xtraterrestrial materials arriving on Earth are the small interplan-tary dust particles (Brownlee, 1985) and micrometeorites (Genget al., 2008). Unfortunately in almost all of these cases available datare few and have been mostly restricted to the light noble gases,ith no information on the more diagnostic Kr and in particular Xe

sotopes.

.1. Asteroidal samples

A number of small grains have been brought back by theayabusa spacecraft from asteroid Itokawa, whose chemical and

sotopic composition links it to metamorphosed LL chondritesNakamura et al., 2011). Noble gas results for three grains haveeen published by Nagao et al. (2011) and for two more by Nagaot al. (2013), who found that the inventory of He, Ne and Ar is dom-nated by implanted solar wind. Given the small size of the grainsvailable for study and assuming concentrations typical for highlyetamorphosed LL chondrites, any planetary Kr and Xe may be

eyond the reach of current instrumentation. Nevertheless, therere plans to analyze the grains previously analyzed for He, Ne, Ary Nagao et al. (2013), using the ultra-sensitive RELAX and RPMS

nstruments at the University of Manchester (Nagao et al., 2013),nd it will be interesting to see those results.

.2. Cometary material

The Stardust mission has brought back cometary grains cap-ured during high-velocity impact in aerogel (Brownlee et al., 2006).etailed results for five particles taken from the cavity walls of theapture tracks have been reported by Marty et al. (2008), howeveror He and Ne only. High concentrations were found. Ne isotopicompositions in the two grains with rather precise data are similaro Ne in the meteoritic Q component, while 3He/4He is interme-iate between Q and solar wind He. Measured 4He/20Ne ratios are25 (down to ∼1 in one particle), much lower than in Q (Table 2).ore preliminary data were reported by Palma et al. (2009, 2010,

012), including puzzling results for aerogel samples without vis-ble tracks or observed particle fragments. One of these yielded ae isotopic composition intermediate between Q and solar wind,hile another released very large amounts of He and Ne, with a low

He/4He of <2.2 × 10−4 and an extraordinary high 20Ne/22Ne ratiof >18, far higher than in any known components.

Mohapatra et al. (2011) have argued that the Stardust aerogeluring its passage through the comet’s coma should have trappedot only cometary grains but also cometary gas, which mightxplain that Palma et al. (2010, 2012) found He and Ne also inerogel without tracks. Similarly puzzling results were found byohapatra et al. (2013) in their study of Ar, Kr and Xe in track-free

erogel. In one of their samples, with Ar and Kr at background levelsike in non-flight aerogel, these authors found a very large amountf Xe with an isotopic composition corresponding (to first order) tohat of mass-fractionated air Xe.

.3. Interplanetary dust particles (IDPs)

A major fraction of the interplanetary dust collected in thetratosphere has a cometary origin, as clearly shown by the collec-ion of IDPs associated with comet Grigg-Skjellerup (e.g., Busemannt al., 2010; Pepin et al., 2011). Busemann et al. (2010) found veryigh Xe concentrations in individual Grigg-Skjellerup IDPs of up to

× 10−7 cm3 STP/g, an order of magnitude higher than upper limitsn a previous study of IDPs from the NASA stratospheric dust col-ection by Kehm et al. (2009), though Hudson et al. (1981) found a


imilarly high abundance (∼1 × 10−7 cm3 STP/g) in a set of 13 par-icles. These are higher than in bulk meteorites and approach thoseor the acid-resistant, carbon-rich, residues. While it is likely, andhile both Hudson et al. (1981) and Busemann et al. (2010) relate


the Xe abundances they find to the presence of planetary or Q-gases,no Xe isotopic data with sufficient precision have been obtained sofar that would prove this without doubt.

Otherwise, most studies of IDPs have concentrated on He andNe and occasionally Ar (e.g., Kehm et al., 2002), elucidating cos-mic ray exposure and heating during atmospheric entry (e.g., Nierand Schlutter, 1993), with a puzzling overabundance of 3He in asubset of samples (Pepin et al., 2000). Unique among these is thestudy by Pepin et al. (2011) of IDPs probably derived from comets26P/Grigg-Skjellerup and 55P/Tempel-Tuttle, with extraordinarilylow 4He/20Ne and high 20Ne/22Ne ratios, which the authors inter-pret as evidence for presolar grains from novae. The results are,however, very sensitive to the accuracy of the corrections for blank(∼54–77% at 22Ne; Pepin et al., 2011) and interferences, so in theeyes of this writer, they are interesting, but have to be taken witha grain of salt.

7.4. Micrometeorites (MMs)

Micrometeorites have generally experienced stronger heatingduring atmospheric entry than the IDPs, nevertheless even (fullymelted) cosmic spherules sometimes contain remnant extraterres-trial gases (Osawa et al., 2003), and partially molten (scoriaceous)MMs have often retained measurable amounts of trapped noblegases (Osawa et al., 2003; Baecker et al., 2012a). Compared to IDPs,their measurement is aided by the larger size (mostly in the range50–1000 �m vs. typically <30 �m for IDPs; Genge, 2008). Never-theless, as for IDPs, most analyses have been restricted to He (e.g.,Stuart et al., 1999) and Ne, with clear evidence for the presence ofimplanted solar wind (see summary in Wieler, 2002). Only rarelyhas Ar been included in the study of MMs (e.g., Marty et al., 2002;Osawa and Nagao, 2002; Osawa et al., 2003).

As for Kr and Xe, early data for one very large micromete-orite (0.23 mg; Sarda et al., 1991) revealed the presence of Xewith a planetary-like isotopic composition, with 36Ar/132Xe and84Kr/132Xe ratios (53 and 0.85, respectively) also falling into therange of typical planetary components (Table 2). Osawa and Nagao(2002) reported Kr and Xe isotope data for a large number ofmicrometeorites, however, these do not include the rare isotopeslighter than 82Kr and 129Xe, resp., and also otherwise are of limitedvalue due to the large analytical errors. Only recently have morecomplete (He through Xe) and precise data sets been obtainedby two groups. Baecker et al. (2012b,c) investigated micromete-orites from the Concordia collection (Duprat et al., 2007) and froma “trap” in the Transantarctic Mountains (Rochette et al., 2008),while Olinger et al. (2013) analyzed “Giant” unmelted micromete-orite samples from the Cap Prudhomme sampling site, with similarresults. Two trends become apparent in these analyses: microme-teorites with Xe that in isotopic composition resembles Q-Xe; andsuch with Xe having a composition suggestive of mass-fractionated(terrestrial) air, most likely acquired during atmospheric entry.

8. Summary and general considerations: planetary versusplanets

Meteorites contain, besides noble gases produced in situ, thosethat are trapped (often also called primordial). Besides implantedsolar wind, the latter comprises what has been traditionally called“planetary” gas, which was trapped early in Solar System historyor even before that. Originally distinguished by its elemental abun-dance pattern with strong depletion of the light noble gases relativeto the heavy ones, planetary gases also have distinct isotopic pat-


terns, which allow to further identify subcomponents that are partof the complex mixture for which the term “planetary” stands.

As discussed in the preceding individual chapters, the mostimportant of the planetary components is the Q component, which


IN PRESSG ModelC


dtthfraiatcTs“dftsotamawtp

namhna(n(tb2beccwtagt(

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“

Fig. 17. Xenon isotopic compositions in planetary gases in meteorites vs. plane-tary atmospheres. Shown are differences relative to solar wind (Table 5; Meshiket al., 2014) in per mill. Shown are Q-Xe (Busemann et al., 2000) and U-Xe (Pepin,1991) in comparison with xenon in the atmospheres of Jupiter (Wieler, 2002), Earth

not quite that simple. Interesting in this context is the fraction-ation between Ar, Kr and Xe determined for noble gas trappingin water ice at 22–27 K at low ice deposition rates (Notesco et al.,

Fig. 18. 36Ar/132Xe vs. 84Kr/132Xe for a representative selection of chondrites, ure-ilites and shergottites, in a log-log representation. Compositions for the Sun, Earthand Mars atmosphere are shown for comparison. The shergottites fall on a differenttrend line from the chondrites, passing roughly through Earth and Mars atmosphere.Adapted from Ott and Begemann (1985) and Ott et al. (1985c). C1 and C2 chondrites:Mazor et al. (1970). C3V chondrites: Matsuda et al. (1980). C3O chondrites: Alaertset al. (1979b). LL chondrites: Alaerts et al. (1979a) and Schelhaas et al. (1990). Hand L chondrites: Schelhaas et al. (1990). Enstatite chondrites: Crabb and Anders


U. Ott / Chemie der

ominates the heavy (Ar, Kr, Xe) noble gas budget of the primi-ive meteorites. Its isotopic Ar, Kr and Xe compositions differ fromhose in the solar wind, but may correspond to those of the bulk Sun,owever, for a firm conclusion reliable knowledge about isotopic

ractionation of the solar wind is required. Exotic components car-ied by nanodiamond/glassy carbon, silicon carbide and graphitere of “stellar origin” and record the processes of nucleosynthesisn their parent stars. Notably, due to the strong elemental fraction-tion in the Q component, it is He and Ne from the diamonds ratherhan from Q, which dominates planetary He and Ne (Fig. 6). Carbonompounds play an extraordinary role as host phases of these gases.he subsolar component most likely started out as early implantedolar wind, which was subsequently elementally fractionated. Asoluble” component present in carbonaceous and ordinary chon-rites may have a similar origin, but deserves more attention inuture work. Other extraterrestrial materials available for labora-ory studies (IDPs, micrometeorites) appear to have sampled theame components as the “normal” meteorites. However, analysesf these are not easy, given how tiny the samples are and the facthat the noble gases are the trace elements par excellence. Considers an example a ∼40 �m diameter sample (∼0.1 �g) of LL5-likeaterial returned from Itokawa by the Hayabusa spacecraft. With

LL5-like 132Xe abundance of ∼2 × 10−10 cm3 STP/g, this sampleould contain just about 500 atoms of 132Xe – obviously, even if

he analysis is perfect, statistical fluctuations in nature come intolay at this stage.

A point only fleetingly addressed above is that the differentoble gas components show different thermal stabilities. Thus, totals well as individual noble gas abundances show a record of ther-al and – to a more limited extent – aqueous alteration of their

ost meteorites. This has been noted in the early days of planetaryoble gas research already, in e.g., the correlation of trapped Ar, Krnd Xe abundances with petrologic type of chondritic meteoritesMarti, 1967), and even the finer subdivision of unequilibrated ordi-ary chondrites (UOC) is reflected in the trapped 36Ar abundanceAnders and Zadnik, 1985). These trends, of course, largely reflecthe abundance of the Q (and the solubles) component. The subject,oth with regard to Q (Huss et al., 1996; see also Busemann et al.,000), and, in particular with regard to the presolar phases haseen treated in great detail (Huss and Lewis, 1994a, 1995; Husst al., 2003, 2013). For example, among the diamond noble gasomponents P3 is most easily lost (with little influence from thehemical environment), while HL and P6 are more tightly bound,hich can be used to estimate temperatures experienced by the

hermally least metamorphosed meteorites (Section 3.2.3; Hussnd Lewis, 1994a). Complete loss of the carriers of exotic nobleases, in the case of unequilibrated ordinary chondrites, occurs inhe order graphite, SiC, diamond with none surviving in types >3.8Huss and Lewis, 1995).

As pointed out in the Introduction, the planetary noble gasomponent in meteorites is not the same as noble gas in plane-ary atmospheres. As for the latter, there is obviously a differenceetween the Gas Giants (Jupiter, Saturn) and the small, rocky innerlanets like Earth and Mars. Jupiter contains Ar, Kr and Xe in roughlyolar proportions, with all somewhat (∼3×) enriched relative toydrogen, and also enriched relative to He and Ne (Mahaffy et al.,000; Wieler, 2002), in contrast to the strong fractionation of theerrestrial and Martian atmospheres (Fig. 1). There are also differ-nces in the isotopic compositions between planetary gas and thelanets, in particular Ne and Xe – and of course He, which is notravitationally bound on Earth and Mars. Unfortunately the errorsre somewhat large in case of Jupiter-Xe, so, as illustrated in Fig. 17,


t is difficult to tell whether Jupiter-Xe is more like solar wind or-Xe or even the hypothetical Xe-U.

Coming back to the elemental patterns, the difference betweenplanetary” and the terrestrial planets Earth and Mars becomes

(contemporary as well as non-radiogenic; Pepin, 1991) and Mars (Swindle, 2002;composition #1). The data points are slightly shifted along the mass axis to aid invisibility.

even more obvious than in Fig. 1, if shown in a three isotope plotof 36Ar/132Xe vs. 84Kr/132Xe (Fig. 18), where the “primitive” and“evolved” bodies fall roughly on two separate “correlation” lines.The difference can be explained (to first order) by either a deficit ofXe in the “evolved” bodies or an excess of Kr. A popular assumptionhas been that of the “missing Xe” (e.g., Ozima and Podosek, 2002;and references therein). However, 36Ar/84Kr in the atmospheres ofEarth and Mars as well as in the shergottites is distinctly differ-ent from what is seen in primitive meteorites, so the situation is


(1981). Ureilites: Göbel et al. (1978) and references therein. Shergottite samples:Becker and Pepin (1984b) and Swindle et al. (1986) for EETA 79001 Lith. C, othersfrom Ott (1988) and Schwenzer (2004). Also shown is the composition of solar gaswhen fractionated as observed by Notesco et al. (2003) for trapping of noble gasesin water ice deposited at low rates at 22–27 K (labeled “ice”).


ING ModelC

2 Erde x

2mgbieado1faP

A

PWpirs

R

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A

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003). This is the only possibly relevant case of trapping known toe where Kr is favored over both Ar and Xe. Starting with a solar

as mixture, the resulting composition plots close to the evolvedodies line in Fig. 18, from which one may infer a role for comets

n the origin of the terrestrial and Martian atmospheres (Notescot al., 2003; Dauphas, 2003). The process, however, at least if oper-ting at higher temperature, does not account for the strong massependent difference in Xe isotopes between Earth and Mars, onne hand, and planetary or solar Xe on the other (Notesco et al.,999), which obviously needs to be taken into account in modelsor the origin of the terrestrial planet atmospheres. This is a sep-rate topic beyond the scope of this review – see, e.g. the work ofepin (1991, 1994, 2000, 2006).

cknowledgments

Thanks go to Johannes-Gutenberg University (Institute forhysics), Mainz, for hospitality during writing of this review. Rainerieler contributed a thorough and constructive review, including

ointing out recent relevant results on solar noble gases that weren press/in revision at the time of writing. Another constructiveeview was provided by Jamie Gilmour. This Invited Review wasolicited and handled by Associate Editor Klaus Keil.

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planetary and pre-solar noble gases in meteorites

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