1.26 Interplanetary Dust Particles J. P. Bradley Lawrence Livermore National Laboratory, Livermore, CA, USA 1.26.1 INTRODUCTION 1.26.2 PARTICLE SIZE, MORPHOLOGY, POROSITY, AND DENSITY 1.26.3 MINERALOGY 126.96.36.199 CP IDPs 188.8.131.52 Glass with Embedded Metal and Sulﬁdes 184.108.40.206 CS IDPs 1.26.4 OPTICAL PROPERTIES 1.26.5 COMPOSITIONS 220.127.116.11 Major Elements 18.104.22.168 Trace Elements 22.214.171.124 Isotopes 126.96.36.199 Noble Gases 1.26.6 CONCLUSIONS ACKNOWLEDGMENTS REFERENCES
How can you appreciate a castle if you don’t cherish all the building blocks? Stephen Jay Gould
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into the solar system, to what extent it has survived, and how it might be distinguished from solar system grains. To better understand the process of solar system formation, it is important to identify and analyze these solid grains. Since all of the alteration processes that modiﬁed solids in the solar nebula presumably had strong radial gradients, the logical place to ﬁnd presolar grains is in small primitive bodies like comets and asteroids that have undergone little, if any, parent-body alteration. Trace quantities of refractory presolar grains (e.g., SiC and Al2O3) survive in the matrices of the most primitive carbon-rich chondritic meteorites (Anders and Zinner, 1993; Bernatowicz and Zinner, 1996; Bernatowicz and Walker, 1997; Hoppe and Zinner, 2000; see Chapter 1.02). Chondritic meteorites are believed to be from the asteroid belt, a narrow region between 2.5 and 3.5 AU that marks the transition from the terrestrial planets to the giant gas-rich planets. The spectral properties of the asteroids suggest a gradation in properties with some inner and main belt C and S asteroids (the source
One of the fundamental goals of the study of meteorites is to understand how the solar system and planetary systems around other stars formed. It is known that the solar system formed from preexisting (presolar) interstellar dust grains and gas. The grains originally formed in the circumstellar outﬂows of other stars. They were modiﬁed to various degrees, ranging from negligible modiﬁcation to complete destruction and reformation during their B108 years lifetimes in the interstellar medium (ISM) (Seab, 1987; Mathis, 1993). Finally, they were incorporated into the solar system. Submicrometer-sized silicates and carbonaceous material are believed to be the most common grains in the ISM (Mathis, 1993; Sandford, 1996), but it is not known how much of this presolar particulate matter was incorporated 1
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region of most meteorites and polar micrometeorites) containing layer silicates indicative of parent-body aqueous alteration and the more distant anhydrous P and D asteroids exhibiting no evidence of (aqueous) alteration (Gradie and Tedesco, 1982). This gradation in spectral properties presumably extends several hundred astronomical units out to the Kuiper belt, the source region of most short-period comets, where the distinction between comets and outer asteroids may simply be one of the orbital parameters (Luu, 1993; Brownlee, 1994; Jessberger et al., 2001). The mineralogy and petrography of meteorites provide direct conﬁrmation of aqueous alteration, melting, fractionation, and thermal metamorphism among the inner asteroids (Zolensky and McSween, 1988; Farinella et al., 1993; Brearley and Jones, 1998). Because the most common grains in the ISM (silicates and carbonaceous matter) are not as refractory as those found in meteorites, it is unlikely that they have survived in signiﬁcant quantities in meteorites. To date only a few presolar silicates have been identiﬁed in meteorites.
Interplanetary dust particles (IDPs) are the smallest and most ﬁne-grained meteoritic objects available for laboratory investigation (Figure 1). In contrast to meteorites, IDPs are derived from a broad range of dust-producing bodies extending from the inner main belt of the asteroids to the Kuiper belt (Flynn, 1996, 1990; Dermott et al., 1994; Liou et al., 1996). After release from their asteroidal or cometary parent bodies the orbits of IDPs evolve by Poynting–Robertson (PR) drag (the combined inﬂuence of light pressure and radiation drag) (Dermott et al., 2001). Irrespective of the location of their parent bodies nearly all IDPs under the inﬂuence of PR drag can eventually reach Earth-crossing orbits. IDPs are collected in the stratosphere at 20–25 km altitude using NASA ER2 aircraft (Sandford, 1987; Warren and Zolensky, 1994). Laboratory measurements of implanted rare gases, solar ﬂare tracks (Figure 2), and isotope abundances have conﬁrmed that the collected particles are indeed extraterrestrial and that, prior to atmospheric entry, they spent 104–105 years as small particles orbiting the Sun (Rajan et al., 1977;
1 µm 20KV
20KV X9400 (c)
Figure 1 (a–c) Secondary electron images. (a) Anhydrous CP IDP. (b) Hydrated CS IDP (RB12A44). (c) Single-mineral forsterite grain with adhering chondritic material. (d) Optical micrograph (transmitted light) of giant cluster IDP (U220GCA) in silicone oil on ER2 collection ﬂag.
3 Solar flare tracks
Solar wind and solar flare effects
Figure 2 Dark-ﬁeld transmission electron micrographs. (a) Solar-wind sputtered rim on exterior surface plus implanted solar ﬂare tracks in chondritic IDP U220A19 (from Bradley and Brownlee, 1986). (b) Solar ﬂare tracks in a forsterite crystal in chondritic IDP U220B11 (from Bradley et al., 1984a). The track densities in both IDPs are B1010–1011 cm2 corresponding to an orbital exposure age of B104 years.
Hudson et al., 1981; Bradley et al., 1984a; McKeegan et al., 1985; Messenger, 2000). During atmospheric entry most IDPs are frictionally heated to within 100 1C of their peak heating temperature for B1 s and, to a ﬁrst-order approximation, the smallest particles are the least strongly heated. Although some IDPs may experience thermal pulses in excess of 1,000 1C for up to 10 s (depending on particle size, mass, entry angle, and speed) (Love and Brownlee, 1991, 1996), the presence of solar ﬂare tracks in an IDP establishes that it was not heated above B650 1C (Bradley et al., 1984a). Since IDPs decelerate from cosmic velocities at altitudes 490 km, where the maximum aerodynamic ram pressure is a factor of B103 less than that exerted on conventional meteorites, extremely fragile meteoritic materials that cannot survive as large objects can survive as small IDPs (Figure 1a) (Brownlee, 1994). Such fragile materials are suspected to be among the most primitive objects and potentially the most informative regarding early solar system and presolar processes. (Conventional meteorites penetrate deep into the atmosphere such that only relatively well-indurated rocks can survive.) Collected IDPs are brieﬂy exposed to the terrestrial environment but since their residence time in the stratosphere is short (B2 weeks), they are not subjected to longer-term
weathering that affects the surfaces of most meteorites (Flynn, 1994a). This chapter examines the compositions, mineralogy, sources, and geochemical signiﬁcance of IDPs. Additional reading can be found in reviews by Fraundorf (1981), Brownlee (1985), Sandford (1987), Bradley et al. (1988), Jessberger et al. (2001), Rietmeijer (1998), and the book edited by Zolensky et al. (1994). Despite their micrometer-scale dimensions and nanogram masses it is now possible, primarily as a result of advances in small particle handling techniques and analytical instrumentation, to examine IDPs at close to atomic-scale resolution. The most widely used instruments for IDP studies are presently the analytical electron microscope, synchrotron facilities, and the ion microprobe. These laboratory analytical techniques are providing fundamental insights about IDP origins, mechanisms of formation, and grain processing phenomena that were important in the early solar system and presolar environments. At the same time, laboratory data from IDPs are being compared with astronomical data from dust in comets, circumstellar disks, and the ISM. The direct comparison of grains in the laboratory with grains in astronomical environments deﬁnes the new discipline of ‘‘astromineralogy’’ (Jaeger et al., 1998; Bradley et al., 1999a, b;
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Molster et al., 2001; Keller et al., 2001; Flynn et al., 2002).
PARTICLE SIZE, MORPHOLOGY, POROSITY, AND DENSITY
Individual IDPs span the diameter range 1– 50 mm, although most are between 5 and 15 mm (Figures 1a–1c). Larger 50–500 mm diameter particles (10–20% of collected IDPs) that fragment into many pieces when they impact the ﬂags are known as giant ‘‘cluster’’ particles (Figure 1d). Two principal morphological groups of IDPs are recognized, porous and smooth (Figures 1a and 1b). Since their bulk compositions are similar to chondritic meteorites (of types CI and CM), they are referred to as chondritic porous (CP) and chondritic smooth (CS) IDPs. CP and CS IDPs are also mineralogically distinct classes of materials (see Section 1.26.3). The morphologies of CP particles resemble a bunch of grapes (Figure 1a). Porosities as high as 70% and densities ranging between 0.3 and 6.0 g cm2 have been measured, although IDPs with densities above 3.5 g cm2 typically contain a large FeNi sulﬁde grain (Fraundorf et al., 1982a; Love et al., 1994). Such low densities, high porosities, and fragile microstructures are consistent with the particulate matter in cometary meteors (Bradley and Brownlee, 1986). Most cluster particles belong to the CP class, presumably because their fragile microstructures predispose them to fragmentation during impact onto the collection substrates. Less common low-porosity CP IDPs are referred to as chondritic ﬁlled (CF) (Schramm et al., 1989). The CS IDPs are mostly solid objects with platy and/or ﬁbrous surface textures (Figure 1b). It is important to note that although there are many particles that fall neatly into this category, some particles have characteristics of more than one particle type or they have unique characteristics. Some IDPs are composed mostly of refractory calcium-, aluminum-, silicon-rich minerals (Zolensky, 1987), and others are single mineral grains like olivines, pyroxenes, and iron-rich sulfides, some of which have adhering ﬁnegrained chondritic material. Their morphologies are deﬁned, to a large extent, by the shape of the mineral grain (Figure 1c).
It has been established using infrared (IR) and electron microscopic studies that there are three principal mineralogical classes of
chondritic IDPs (Figure 3). They are referred to as ‘‘pyroxene,’’ ‘‘olivine,’’ and ‘‘layer silicate’’ after their most abundant silicate minerals. In a study of 26 IDPs, Sandford and Walker (1985) classiﬁed B25% as ‘‘pyroxenes,’’ B25% as ‘‘olivines,’’ and B50% as ‘‘layer lattice silicate’’ IDPs. The pyroxene and olivine classes are usually porous CP IDPs and contain only anhydrous minerals, while the layer silicate classes are usually smooth CS IDPs and contain hydrous silicates (clays). Most IDPs fall into this mineralogical classiﬁcation scheme, although IDPs with intermediate morphology and mineralogy are relatively common. Examples include anhydrous IDPs with similar amounts of pyroxene and olivine (Sandford and Walker, 1985; Bradley et al., 1989, 1992), porous IDPs containing minor amounts of hydrated layer silicates (Thomas et al., 1995), and smooth layer silicate particles containing large anhydrous silicate grains (Germani et al., 1990).
This class of IDPs has been most intensively examined primarily because their high porosities and ﬂuffy microstructures are unique among other known classes of extraterrestrial materials. IR spectra indicate that silicates are the most abundant minerals in CP IDPs and that the absence of hydrous silicates (clays) is a fundamental distinguishing property (Sandford and Walker, 1985; Bradley et al., 1992). Transmission electron microscopy studies of electron transparent thin sections conﬁrm the absence of hydrous minerals and reveal that CP IDPs are heterogeneous aggregates of predominantly submicrometer-sized crystalline mineral grains (olivine, pyroxenes, iron-rich sulfides, and FeNi metal), polycrystalline aggregates (e.g., glass with embedded metal and sulfides (GEMS), see Section 188.8.131.52), silicate glasses, and carbonaceous material (Bradley, 1994a) (Figures 4a and 4c). Enstatite and forsterite with o5 mol.% Fe are the most common crystalline silicates, although crystals with up to 30 mol.% Fe are also observed (Christoffersen and Buseck, 1986; Bradley et al., 1989, 1999b). Enstatite-rich particles are typically more ﬁne-grained than olivinerich particles. Solar ﬂare tracks have been observed in both enstatite- and forsteriterich IDPs, conﬁrming that they are indeed extraterrestrial and that they were not strongly heated during atmospheric entry. However, while tracks are observed in most enstatite-rich IDPs they are conspicuously absent in most but not all forsterite-rich IDPs. Some forsterite-rich IDPs exhibit melt textures,
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ol 8.0 9.0 10.0 11.0 12.0 Wavelength (µm)
(e) U230A43 SI
Carbonate Layer silicate 7.0 8.0 9.0 10.0 11.0 12.0 Wavelength (µm)
Figure 3 (a–c) IR spectra from thin sections of pyroxene-rich CP IDP W7027H14, olivine-rich CP IDP U220A14, and layer silicate-rich CS IDP U230A43. (d–f) Corresponding X-ray point count analyses obtained from the thin sections on a two-dimensional grid using a 200 keV electron probe with o50 nm spatial resolution at each point. Solid boxed area in (f) shows Mg–Fe–Si composition of the layer silicate and dotted boxed area shows carbonate Mg–Fe composition. Source: Bradley et al. (1992).
equilibrated silicate mineralogy, and surfaces that are decorated with magnetite that almost certainly formed from frictional heating during atmospheric entry (Germani et al., 1990). Therefore, it is likely that the olivine-rich subset of CP IDPs includes and is perhaps dominated by particles that were thermally modiﬁed during atmospheric entry. The morphologies, crystal structures, and compositions of some enstatite and forsterite
crystals in (track-rich) CP IDPs suggest that they formed by vapor-phase growth. (Gas-tosolid condensation is the fundamental mechanism by which grains are formed from nebular gases throughout the galaxy.) Enstatite (MgSiO3) crystals in IDPs have distinctive ultrathin platelet morphologies while others are whiskers with crystallographic screw dislocations characteristic of vapor-phase growth (Figure 5) (Bradley et al., 1983). Forsterite
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Chondritic porous (CP)
Chondritic smooth (CS)
(b) Chondritic smooth (CS)
Chondritic porous (CP) FeS Carbonaceous material
Saponite FeNi metal (kamcite) 0.1 µm
Saponite FeS 0.1 µm
Figure 4 Bright-ﬁeld electron micrographs of ultramicrotomed thin sections: (a) CP IDP U222B42, (b) CS IDP U230A43, (c) CP IDP U222B42 with carbonaceous material, GEMS, crystalline silicates (enstatite), and sulfides (FeS), (d) CS IDP U222C29 with ﬁbrous layer lattice silicate (saponite), pyrrhotite (FeS), and the pyroxene diopside (Di).
and enstatite grains in some IDPs contain up to 5 wt.% MnO, in contrast to pyroxenes and olivines in meteorites that typically contain o0.5% MnO. Since iron and manganese are coupled during crystallization from a liquid melt and decoupled during vapor-phase condensation, Klo¨ck et al. (1989) propose that the (low-iron) manganese-enriched forsterite and enstatite grains in IDPs are vapor-phase condensates. CP IDPs contain more silicate glass than any other class of chondritic materials. Nonstoichiometric magnesium-rich silicate glass is most common although other glass compositions with highly variable amounts of oxygen, magnesium, aluminum, silicon, calcium, titanium, and iron are observed. The glasses occur as discrete grains, rims on silicate crystals, and within GEMS (Section 184.108.40.206) (Bradley, 1994a, 1994b; Brownlee et al., 1999). Sulfides are the second most abundant class of minerals in CP IDPs (Figure 6). The most
common sulﬁde is pyrrhotite with up to B20 at.% Ni and a hexagonal unit cell (Fraundorf, 1981; Zolensky and Thomas, 1995; Dai and Bradley, 2001). Some crystals exhibit superlattice reﬂections. Grain sizes span a huge range from B10 nm to B5 mm. Rare troilite, pentlandite, sphalerite, and NiS crystals have also been observed or reported (Christoffersen and Buseck, 1986; Zolensky and Thomas, 1995). A cubic ‘‘spinel-like’’ sulﬁde with a composition similar to the hexagonal pyrrhotite has also been identiﬁed in CP IDPs (Dai and Bradley, 2001). Both polymorphs are sometimes coherently intergrown on a unit cell scale. The cubic polymorph is metastable and it transforms into hexagonal pyrrhotite when it is heated in the electron beam (Figures 6c and 6d). Since most IDPs are pulse-heated above 200 1C during atmospheric entry (Sandford and Bradley, 1989; Love and Brownlee, 1991), it is possible that much of the hexagonal pyrrhotite in IDPs is a secondary thermal alteration product of
a a 400
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Figure 5 (a) Secondary electron image of a CP IDP with embedded enstatite whiskers (and a platelet, upper left). (b) (Left) Bright-ﬁeld transmission electron micrograph of a segment of a clinoenstatite rod viewed down the b crystallographic axis. The two features parallel to the axis of the rod are screw dislocations and the features cutting across the rod are (100) stacking faults. (Right) Selected area electron diffraction pattern showing the (h00) reciprocal lattice row from the rod. The side bands (splitting) of the diffracted beams, especially visible for the 600 and 800, are related to the helical lattice distortions arising from the dislocations (see Bradley et al., 1983). (c) Bright-ﬁeld electron micrograph of ﬁve ultrathin enstatite ribbons and platelets. (d) Dark-ﬁeld electron micrograph of an enstatite ribbon. Striations in the crystal result from extreme stacking disorder associated with unit-cell scale intergrowths of orthorhombic, orthoenstatite, and monoclinic clinoenstatite. All of enstatite crystals are likely condensates from a nebular gas. Reproduced by permission of Nature Publishing Group from Bradley et al. (1983).
the cubic sulﬁde. Although cubic spinel-like sulfides with pyrrhotite compositions have not been reported in nature, a crystallographically similar nickel-free pentlandite was synthesized in the laboratory by low-temperature (o200 1C), low-pressure vapor-phase growth (Nakazawa et al., 1973). Therefore, it is possible that some or all of the sulfides in IDPs, such as the pyroxene whiskers and platelets (Figure 5), formed by direct condensation from a nebular gas. Alternatively, the sulfides may have formed by gaseous sulﬁdization of preexisting FeNi metal grains (Lauretta and Fegley, 1994, 1995; Lauretta et al., 1996; Zolensky and Thomas, 1995). Since the cubic sulﬁde has not been found in chondritic meteorites, it is likely that some sulfides in IDPs formed under
conditions signiﬁcantly different from those in conventional meteorites. Flynn (2000) observed selenium levels in IDPs 60% higher than those in meteoritic sulfides and concluded that IDP sulfides may have formed in a different environment than the sulfides in meteorites. Carbonaceous material is widespread throughout CP IDPs both as discrete inclusions of noncrystalline material and as a semicontinuous matrix with embedded mineral grains. Often it has a vesiculated appearance consistent with an organic component (Figure 7a). The bulk abundance of carbon in CP IDPs varies from B4% to B45% with an average of B13% (Keller et al., 1994). In contrast to the ﬁne-grained matrices of carbonaceous chondrites, ordered (graphitic) carbon exhibiting B3.4 A˚ lattice
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0006 0226 0220
200 nm  (a)
111 0110 200
Figure 6 (a) Bright-ﬁeld electron micrograph of a pyrrhotite crystal in a thin section of CP IDP U222B28. (b) Selected area electron diffraction pattern from the pyrrhotite crystal in (a) exhibiting prominent superlattice reﬂections consistent with the a ¼ 2A, c ¼ 6C superstructure. Inset lower right is a magniﬁed view of the central region of the pattern showing the 6C reciprocal lattice periodicity. (c–d) Electron micro-diffraction patterns from a pyrrhotite crystal in CP IDP W2070-8D before and after extended electron irradiation. The initial pattern (c) shows two reciprocal lattice nets, the stronger indexed as hexagonal pyrrhotite viewed along [0 1 0] (thick-line box) and the weaker as cubic spinel-like sulﬁde viewed along [0 1 1] (thin-line box). After the grain was illuminated in electron beam for several tens of seconds (c) only the strong reﬂections remain, indicating that the cubic sulﬁde has transformed into hexagonal pyrrhotite. Source: Dai and Bradley (2001).
fringes is rare in chondritic IDPs where it is rarely observed as rims on the surfaces of FeNi metal and FeNi carbide grains (e.g., Bradley et al., 1984b). The carbonaceous material is present in clumps or as a semicontinuous matrix throughout which submicrometer-sized grains (e.g., GEMS and sulfides) are distributed. While elemental carbon is almost certainly a major component, IR spectra showing prominent C–H stretching resonances establish the presence of an aliphatic organic component in chondritic IDPs (Figure 7b) (Flynn et al., 2000). Using two-step laser desorption mass spectroscopy (mL2MS), Clemett et al. (1993) have shown that both porous and smooth IDPs contain polyaromatic hydrocarbons (PAHs). Nitrogen may be associated with the PAHs because the mass spectra of the PAHs are dominated by odd mass species in the intermediate molecular weight range from 200 to 300 amu. This is in contrast to the results obtained previously from PAHs in primitive meteorites using the same analytical technique. The odd mass peaks could be due to the
substitution of functional groups containing odd numbers of nitrogen atoms such as cyano (–CN) and amino (–NH2) groups. Spatially correlated 13C depletion and 15N enrichments were recently observed in an IDP (Floss and Stadermann, 2003). Nanodiamonds have been identiﬁed in large cluster CP IDPs but not so far in smaller noncluster CP IDPs (Figures 7c and 7d) (Dai et al., 2002). The nanodiamonds are similar in size distribution and abundance to those found within carbonaceous chondrites although in one IDP (U220GCA, Figure 1d) their abundance appears to be B10 higher. Defect structures in meteoritic nanodiamonds suggest that they formed by a vapor deposition process as opposed to shock metamorphism (Daulton et al., 1996). The carrier of nanodiamonds in IDPs is a disordered (amorphous) carbonaceous material that may contain organic components (e.g., PAHs) (Dai et al., 2003). It is unclear why nanodiamonds are found in cluster IDPs and depleted or absent in smaller
Orqueil acid residue
3.4 3.5 3.6 Wavelength (µm)
0.206 nm (d)
Figure 7 Bright-ﬁeld transmission electron micrograph of vesiculated amorphous carbonaceous material (in CP IDP W7027A1) that likely contains organic and inorganic carbon. (b) Comparative IR spectra from Orgueil (CI) meteorite acid residue and CP IDP W7207A3 showing prominent C–H stretch features at 3.4–3.5 mm indicative of aliphatic hydrocarbons (Brownlee et al., 2000). (c) Nanodiamond embedded in a carbonaceous mantle on the surface of a sulﬁde crystal at the edge of a GEMS (cluster IDP U220GCA). (d) A single nanodiamond within an acid-etched residue of cluster IDP W7110A-2E-D. Reproduced by permission of Lunar and Planetary Institute from Bradley et al. (1996a).
noncluster IDPs. Moreover, the relationship and distinction between large cluster and smaller noncluster IDPs is also unclear. It has been suggested that cluster IDPs are more isotopically primitive (Messenger, 2000), but their compositions, mineralogy, and petrography appear identical to those of smaller noncluster CP IDPs. Despite their large sizes, cluster IDPs appear less strongly heated than is predicted by atmospheric entry heating models, implying that they are captured from low-speed asteroidal or Kuiper belt orbits (Thomas et al., 1995; Liou et al., 1996; Flynn, 1996; Brownlee et al., 1995). Since (asteroidal) chondritic meteorites contain nanodiamonds, it is not unexpected that cluster IDPs also contain nanodiamonds, assuming they are from asteroids. Irrespective of the origins of cluster IDPs, the absence or depletion or nanodiamonds in smaller carbon-rich chondritic IDPs, some or all of which may be cometary, suggests that nanodiamonds are heterogeneously distributed throughout the solar system and that they may actually be more abundant in asteroids than in
comets. Therefore, it is possible that most meteoritic nanodiamonds formed within the inner solar system in the vicinity of asteroid accretion and not in a presolar environment as is widely believed, although xenon and tellurium isotopes indicate that at least a small fraction of them must be of presolar origin (Hoppe and Zinner, 2000). A solar system origin could explain the anomalously high abundance of nanodiamonds relative to other types of presolar grains in meteorites (Hoppe and Zinner, 2000). The recent detection by the Infrared Space Observatory (ISO) of nanodiamonds formed in situ within the accretion disks of young stars conﬁrms that nanodiamonds could indeed have formed in the inner solar system (Van Kerckhoven et al., 2002).
Glass with Embedded Metal and Sulfides
GEMS are perhaps the most exotic class of primitive meteoritic materials yet encountered
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C C (a)
(b) 100 nm
Relict FeS (c)
Relict forsterite (d)
Figure 8 Transmission electron micrographs of GEMS within CP IDPs. (a) Bright-ﬁeld image of GEMS embedded within amorphous carbonaceous material (labeled ‘‘C’’). The inclusions are FeNi metal (kamacite) and Fe sulfides. (b) Dark-ﬁeld image. Bright inclusions are metal and sulfides; matrix is magnesium silicate glass. (c, d) Dark-ﬁeld images of GEMS with ‘‘relict’’ sulﬁde and forsterite (Mg2SiO4) inclusions. Source: Bradley et al. (1999a).
(Figure 8). They are 0.1–0.5 mm spheroids that are ubiquitous throughout the matrices of CP IDPs. Because their bulk compositions are approximately chondritic they can be viewed as a discrete class of picogram-mass chondritic meteorites. Since they may also contain small amounts of carbon it is possible that they are carbonaceous chondrites (Bradley, 1994b; Brownlee et al., 2000). Despite their submicrometer dimensions, picogram masses, and unique nanometer-scale heterogeneity, their bulk compositions are similar to those of kilogram-mass chondritic meteorites and kiloton-mass asteroids. Terminology used to describe GEMS has evolved as increasingly sophisticated instruments have been used to analyze them, particularly those capable of highest spatial resolution and light element (carbon and oxygen) analyses. GEMS have been described as ‘‘tar balls,’’ ‘‘granular units,’’ ‘‘microcrystalline aggregates,’’ ‘‘unequilibrated aggregates,’’ and ‘‘polyphase units’’ (Bradley, 1988; Rietmeijer, 1989, 1997; Klo¨ck and Stadermann, 1994). They are mineralogically unequilibrated in that they contain nanometer-sized inclusions of FeNi alloy (kamacite) and iron-rich sulﬁde (pyrrhotite) embedded in oxygen-rich, low-iron magnesium silicate glass. Figure 8 shows a bright- and several darkﬁeld images of GEMS. Typically, GEMS are
found embedded within amorphous carbonaceous matrix (Figure 8a). In Figure 8b the bright inclusions are 2–100 nm diameter FeNi metal (kamacite) and pyrrhotite nanocrystals and the uniform gray matrix is magnesium-rich silicate glass. Some GEMS contain deeply eroded ‘‘relict’’ pyrrhotite, forsterite, or enstatite grains toward their cores (Figures 8c and 8d). Experiments with irradiated mineral standards and observations of the surfaces of lunar soil grains exposed to the solar wind indicate that the mineralogy and petrography of GEMS were shaped primarily by exposure to ionizing radiation and the exposure occurred ‘‘prior to’’ the accretion of the host CP IDPs (Bradley, 1994b; Bradley et al., 1996a). The most intriguing property of GEMS is their similarity to ‘‘amorphous silicate’’ grains that are ubiquitous through interstellar space. The presence of silicate grains in the ISM is revealed by spectral features at B1,030 and B525 cm1 (9.7 and 19 mm) corresponding to the Si–O stretching and Si–O–Si bending modes of silicates (Millar and Duley, 1980; Aitken et al., 1989). The features are observed both in absorption and emission along various lines-of-sight and they generally lack ﬁne structure, which suggests that the silicates are predominantly amorphous (i.e., glasses). The size range of the grains inferred from extinction is between 0.1 and 0.5 mm (Kim et al., 1994).
Mineralogy These grains are believed to have originally formed in the atmospheres and outﬂows of oxygen-rich post-main-sequence AGB stars (Mathis, 1993; Henning, 1999). With the exception of carbon, most of the condensed rockforming elements in ISM are associated with these silicate grains (Snow and Witt, 1996). Immediately prior to the collapse of the solar nebula, most of the condensed atoms in the solar system were carried within these interstellar amorphous silicate grains. The physical and chemical properties of interstellar silicates inferred from astronomical observations match the exotic properties of GEMS. For example, it has been proposed that polarization of starlight caused by alignment of IS amorphous silicates in the galactic magnetic can be explained by nanometer-sized inclusions of superparamagnetic FeNi metal (Jones and Spitzer, 1967). Other properties of GEMS including their bulk compositions and size distribution match those of IS amorphous silicates (Bradley et al., 1997). The only way to prove that GEMS are indeed presolar interstellar amorphous silicate grains is to measure nonsolar isotope abundances. But if GEMS truly are interstellar amorphous silicates, it is not entirely clear that they should have nonsolar isotopic compositions, because grains undergo considerable processing during their 108–109 years lifetimes in the ISM such that the chemical and isotopic compositions of most grains are likely homogenized (Seab, 1987; Mathis, 1993). It is signiﬁcant that, without exception, all of the presolar grains so far identiﬁed in meteorites are highly refractory minerals capable of withstanding processing in the ISM and incorporation into asteroidal parent bodies. Although the glassy silicate structures, chondritic (solar) compositions, and irradiation effects in GEMS are consistent with considerable grain processing, relict grains indicate that some GEMS may retain a memory of their circumstellar origins (Figures 5c and 5d). The presolar interstellar origin of two GEMS was conﬁrmed by measurements of their oxygen isotopic compositions using the new-generation NanoSIMS ion microprobe (Messenger et al., 2003). Both exhibit 16O abundances signiﬁcantly different from those of solar system abundances, indicating that they contain preserved circumstellar silicate components, and conﬁrming that they are interstellar ‘‘amorphous silicates.’’ It is pointed out that the other 40 GEMS analyzed (even in the same IDP sections) have perfectly normal isotopic compositions and that it is therefore not possible to conclude whether they are solar system materials or isotopically homogenized interstellar grains. The observation that some GEMS are presolar indicates that at least some of them may indeed be the
long-sought interstellar ‘‘amorphous silicates’’ (see Flynn, 1994b; Goodman and Whittet, 1995; Martin, 1995). In addition to GEMS, there are other types of polycrystalline aggregates in CP IDPs. Rietmeijer (1997, 1998) described course-grained and ultraﬁne-grained ‘‘polyphase units’’ composed of glass, Mg–Fe olivine and pyroxene, oxides, sulfides, and metal. Other submicrometer grains described as ‘‘equilibrated aggregates’’ also have bulk compositions that are approximately chondritic (Bradley, 1994a). But in contrast to GEMS that have highly unequilibrated compositions and mineralogy, ‘‘equilibrated aggregates’’ contain iron-bearing olivine and pyroxene grains with equilibrated Fe/Mg ratios and iron sulfides embedded in aluminosilicate glass. Their textures and petrography suggest that they have an igneous origin (e.g., formation by collisional shock melting). ‘‘Reduced aggregates’’ are yet another component of chondritic IDPs. These carbon-rich aggregates contain FeNi metal, FeNi carbide, and FeNi sulﬁde crystals embedded in carbonaceous matrix (Bradley, 1994a). Some of the metal grains are rimmed with a thin (o10 nm) rim of graphitic carbon (Bradley et al., 1984b). Since all of these aggregates have undergone an accretional event and processing ‘‘prior to’’ accretion of the IDPs in which they reside, they are older than the IDPs. Therefore, the study of aggregates in CP IDPs provides a window back to times predating the accretion of interplanetary dust. Whether the accretion occurred in the early solar system or presolar interstellar environments is unknown at this time.
CS IDPs are low-porosity objects composed predominantly of hydrated layer lattice silicates (clays) (Figures 1b and 9). IR spectra indicate that silicates are the most abundant minerals in CS IDPs, some of them contain carbonates, and that the presence of hydrous silicates (clays) is a fundamental distinguishing property (Sandford and Walker, 1985; Germani et al., 1990; Bradley et al., 1992). Transmission electron microscopy studies of electron-transparent thin sections conﬁrm the presence of hydrous layer lattice silicates and carbonates. Their mineralogical and petrographic similarity to the ﬁne-grained matrices of type CI and CM carbonaceous chondrites is unmistakable but there are important differences. Whereas serpentine is the dominant layer lattice silicate in CI and CM chondrites, smectite is the dominant layer silicate in CS IDPs, suggesting that there were signiﬁcant differences in the parent
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RB12A44−1 10.8 Å 7A 7 5 11 7 7 7 7
Figure 9 Lattice-fringe images of cronstedite and tochilinite in CS IDP RB12A44 (see also Figure 1b). (a) Mixture of iron-rich serpentine cronstedite (7.3 A˚ spacing) and tochilinite (5.4 and 10.8 A˚ spacings). (b) Unit cell intergrowth of cronstedite and tochilinite. (c) Pseudo-rectangular tochilinite ‘‘prototube’’ nucleated in a cronstedite plate. (d) Tochilinite tube. Source: Bradley and Brownlee (1991).
bodies of the IDPs and meteorites (Zolensky and McSween, 1988; Brearley and Jones, 1998; Germani et al., 1990). In contrast to the layer silicates in CI/CM chondrites, those in CS IDPs are poorly crystallized, the predominant basal spacing is B10 A˚, and they are compositionally more heterogeneous on a scale of B50 nm than the ﬁne-grained matrices of carbonaceous chondrites (Germani et al., 1990). Most of the layer silicates in CI and CM chondrites are serpentine with an B7 A˚ basal spacing and they formed mostly from crystalline silicates like olivine and pyroxene. The layer silicates in CS IDPs formed in situ by aqueous alteration of silicate glasses and their compositions plus the B10 A˚ spacing suggest that they are smectites. A variety of other minerals have been reported in CS IDPs. They include the anhydrous crystalline silicates diopside (Figure 4d), enstatite, fassaite, and forsterite; amorphous silicates (glasses) with variable amounts of magnesium, calcium, aluminum, and iron; the sulfides pyrrhotite, troilite, and nickel sulﬁde; the oxides magnetite and chromite; and a phosphide, schreibersite. The sulﬁde mineralogy of CS IDPs differs from that of CP IDPs. Whereas low-nickel pyrrhotite ([Fe,Ni]1xS) is the dominant sulﬁde in CS IDPs (Dai and Bradley, 2001), nickel-rich sulfides with compositions ranging from low-nickel pyrrhotite compositions up to and including pentlandite ([Fe,Ni]9S8) are more abundant in CS IDPs (Zolensky and
Thomas, 1995). A ‘‘low-nickel pentlandite’’ has been reported in a hydrated CS IDP (Tomeoka and Buseck, 1984). The composition and crystal structure of the low-nickel pentlandite are similar to those of the cubic spinel-like sulﬁde identiﬁed in CS IDPs (Dai and Bradley, 2001). To conﬁrm whether they are one and the same mineral requires more study. CS IDPs are, on average, signiﬁcantly enriched in carbon relative to CI chondrites and contain disordered carbonaceous material similar to that found in CP IDPs. Carbon abundances vary from 5% to 420% with an average of B13% (Keller et al., 1994). The hydrated silicate mineralogy of CS IDPs indicates that they are derived from parent bodies in which aqueous alteration has occurred. High-nickel sulfides (e.g., pentlandite) are also consistent with formation during asteroidal parent-body aqueous alteration because Ni-rich sulfides form at relatively high oxygen fugacities (Godlevskiy et al., 1971). Asteroids are the logical parent bodies, since it is well established that aqueous alteration is an important parentbody process within at least some regions of the asteroid belt (Brearley and Jones, 1998). The mineralogy of several CS IDPs provides a direct connection to the asteroids. Tochilinite, an ordered mixed-layer mineral containing magnesium, aluminum, iron, nickel, sulfur, and oxygen, identiﬁed in CS IDP RB12A44, has been found in only one other class of meteorites, the type CM carbonaceous chondrites (Bradley
Intensity (arb. units)
The optical properties of chondritic IDPs have been measured in the IR and visible spectral regions (Figures 3, 10, and 11). Most IR measurements have been acquired in transmission over the 2–25 mm IR region using microscope spectrophotometers equipped with globar sources (Figure 3). More recently,
(a) Intensity (arb. units)
high-brightness synchrotron light sources have proved to be ideal for spectral microanalysis of IDPs and even subcomponents of IDPs (e.g., GEMS and sulfides). Synchrotron sources can deliver an apertured beam as small as B3 mm diameter spot 4100 brighter than globar sources. The 2–25 mm region includes the important B10 and B20 mm silicate features. Almost all chondritic IDPs exhibit a dominant B10 mm silicate feature, the position and shape of which have been used to classify particles as ‘‘pyroxene,’’ ‘‘olivine,’’ or ‘‘layer silicate’’ IDPs. The 10 mm feature of pyroxenerich CP IDPs typically consists of principal bands at 9.1–9.4 mm (1,064–1,099 cm–1) and 10.5–10.75 mm (930–953 cm1) that are consistent with monoclinic pyroxene. Olivine-rich CP IDPs exhibit an intense band at 11.2–11.3 mm (885–892 cm1) with less intense bands at 10.1, 10.75, and 11.9 mm (840, 930, and 990 cm1). Hydrated CS IDPs usually produce a single featureless band at 9.7–9.8 mm (see Sandford and Walker, 1985; Bradley et al., 1992). The 10 mm feature of chondritic IDPs has been compared with the 10 mm feature of astronomical silicates. No particular IDP IR class
(a) 9.7 µm
and Brownlee, 1991). Similarly, unit cell-scale intergrowths of serpentine and saponite observed in CSIDP W7013F5 are also found within the ﬁne-grained matrices of type CI chondrites (Keller et al., 1992). The presence of these distinctive secondary mineral assemblages provides direct petrogenetic links between some IDPs and speciﬁc classes of meteorites and thus conﬁrms that some IDPs collected in the stratosphere do indeed have an asteroidal origin (Rietmeijer, 1996). But the scarcity of these IDPs suggests that CS IDPs sample a broad range of hydrous parent bodies and that materials with CM and CI mineralogy are not abundant among the hydrous dust-producing asteroids.
Emissivity (arb. units)
Emissivity (arb. units)
9 10 11 12 Wavelength (µm)
9 10 11 12 Wavelength (µm)
Figure 10 Comparison of the 10 mm Si–O stretch bands of a ‘‘GEMS-rich’’ IDP and astronomical silicates. Left: (a) CP IDP L2008V42A, proﬁle derived from transmittance spectrum; (b) comet Halley (Campins and Ryan, 1989); (c) comet Hale-Bopp (Hanner et al., 1997); (d) late stage Herbig Ae/Be star HD163296 (Sitko et al., 1999). Right: (a) GEMS (in IDP L2011*B6); (b) Elias 16 molecular cloud (Bowey et al., 1998); (c) trapezium molecular cloud (Hanner et al., 1995); (d) pre-main sequence T Tauri YSO DI Cephei (Hanner et al., 1998); (e) post-main sequence M-type supergiant m-Cephei (Aitken et al., 1988). Source: Bradley et al. (1999a).
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14 20 Reflectance (%)
5 0 300
600 700 500 Wavelength (nm)
10 Reflectance (%)
W7030A5 8 6 4 2 0 300
500 600 700 Wavelength (nm)
wavelength range (Figure 11). Interpretation of the indigenous reﬂectance characteristics of IDPs can be complicated because of spurious scattering effects from large mineral grains, secondary magnetite formed on the surfaces of some IDPs during atmospheric entry, and other small particle light scattering artifacts. In general, chondritic IDPs are spectrally dark objects with o15% reﬂectivity over the 400–800 nm range. Most anhydrous CP IDPs dominated by enstatite, forsterite, and GEMS exhibit spectral characteristics similar to those of smaller, more primitive solar system objects (e.g., P and D asteroids). Carbon-rich CP IDPs are spectrally red with a redness comparable to the comet-like outer asteroid Pholus (Binzel, 1992). Hydrated CS IDPs that contain layer lattice silicates exhibit spectral characteristics similar to carbonaceous chondrites and main-belt C-type asteroids (Bradley et al., 1996b).
Figure 11 Reﬂectance spectra of (a) CSIDP W7030A15 and (b) CP IDP W7030A5. Source: Bradley et al. (1996b).
consistently matches the B10 mm feature of solar system comets or silicate dust in the ISM (Sandford and Walker, 1985). However, the B10 mm features of CP IDPs composed mostly of GEMS and submicrometer enstatite and forsterite crystals generally resemble those of comets and late-stage Herbig Ae/Be stars in support of the hypothesis that some CP IDPs are of cometary origin (Figure 10). The IR spectral features of individual subcomponents of chondritic IDPs have also been measured. A measurement of the B10 mm silicate feature of GEMS produced a broad featureless band at B9.7 mm that matches the spectra of interstellar molecular cloud dust, providing further evidence in support of the controversial hypothesis that GEMS are interstellar amorphous silicates (Figure 10) (Bradley et al., 1999a). This hypothesis was recently conﬁrmed by the discovery of GEMS with oxygen isotopic compositions indicating that they are indeed presolar silicates (Messenger et al., 2003). Some chondritic IDPs exhibit a broad feature at B23.5 mm and a similar broad feature is seen in the IR spectra of young stellar objects. Detailed mineralogical analyses of the IDPs in conjunction with IR spectroscopy of mineral standards established that iron sulfides are responsible for the B23.5 mm feature (Keller et al., 2001). Reﬂectance spectra have been collected from chondritic IDPs over the visible 450–800 nm
COMPOSITIONS Major Elements
The bulk compositions of several hundred IDPs like those shown in Figure 1 have been measured using electron beam X-ray energydispersive spectroscopy (EDS), synchrotron X-ray ﬂuorescence (SXRF), proton-induced X-ray emission (PIXE), and instrumental neutron activation analysis (INAA). The benchmark standard of comparison is with CI chondritic meteorites. The CI meteorites are considered to be the most chemically primitive class of meteoritic materials, because their bulk compositions, more than any other class of meteorites, closely correspond to the composition of the solar corona (see Chapter 1.03). Within a factor of 2–3, the element ratios for most chondritic CP and CS IDPs match those of the CI chondrites. (Carbon is an exception with abundances 5 higher than CI chondrites (Keller et al., 1994).) The compositions of CP IDPs are chondritic (solar) on a scale of less than 1 mm, indicating that they are mineralogically heterogeneous on a submicrometer scale (Bradley et al., 1989). CS IDPs are less heterogeneous, presumably as a result of in situ aqueous alteration (Germani et al., 1990). IDPs dominated by a single mineral grain (e.g., forsterite or pyrrhotite) typically have nonchondritic compositions reﬂecting the composition of the grain. Other stratospheric particles identiﬁed as IDPs include the so-called refractory IDPs rich in the elements calcium, aluminum, and titanium (Zolensky, 1987).
Table 1 Mean atomic element/Si ratios for stratospheric micrometeorites versus those of various chondritic meteorite classes. Ca
Chondritic All CS CP Coarse
IDPsb 1.75 1.32 2.39 1.31
4.17 4.49 3.98 3.81
0.052 0.051 0.056 0.043
0.980 0.824 1.015 1.203
0.075 0.082 0.070 0.075
0.356 0.341 0.417 0.231
0.052 0.021 0.047 0.125
0.015 0.014 0.016 0.013
0.697 0.742 0.705 0.585
0.027 0.032 0.024 0.019
Chondritic CIc CMd Ld
meteorites (bulk) 0.70 7.64 0.057 0.35 4.38 0.029 0.02 3.49 0.046
1.040 1.023 0.928
0.083 0.088 0.067
0.444 0.201 0.099
0.061 0.070 0.050
0.013 0.012 0.011
0.868 0.804 0.594
0.048 0.048 0.032
Chondritic meteorites (fine-grained matrices) CIe NA NA 0.016 0.920 NA NA 0.038 0.957 CMe
C and O analyses were done for only 30 IDPs. b IDP data from Schramm et al. (1989). c CI chondrite average; Palme and Jones (Chapter 1.03). CM and L chondrite averages calculated from Jarosewich (1990). e CI and CM matrix compositions from McSween and Richardson (1977).
The major-element compositions of 200 chondritic IDPs were measured by EDS (Table 1 and Figure 12). All of the particles were identiﬁed as extraterrestrial because they have approximately chondritic compositions or consist predominantly of a single mineral grain-like forsterite or pyrrhotite (commonly found within chondritic IDPs): 37% of the particles are CS IDPs, 45% are CP IDPs, and 18% IDPs composed predominantly of a single mineral. Table 1 summarizes the compositions of the IDPs. Within a factor of 2 the abundances of oxygen, magnesium, aluminum, sulfur, calcium, chromium, manganese, iron, and nickel are approximately chondritic. CP IDPs are a closer match to CI carbonaceous chondrites than CS IDPs, and they are closer to CI bulk than to CI or CM matrix. Anhydrous CP IDPs are the only known meteoritic materials that have a composition at the nanometer scale that is similar to CI bulk. Despite the compositional similarities between CP and CS particles there are signiﬁcant differences. While CP IDPs are a close match to CI abundances (they are a closer match to CI bulk than to CI or CM matrix), the CS group shows systematic magnesium and calcium depletions and a stoichiometric ‘‘excess’’ of oxygen. The mean Mg/Si ratio for CP IDPs is 6% below the CI mean but the Mg/Si ratio of CS IDPs is 25% below the CI mean. The Ca/Si ratio shows a large range with a mean for all IDPs that is depleted by 15% relative to CI. Like Mg/Si, there is a clear difference between the CP and CS particles, with the former containing normal calcium and the latter depleted in calcium (see also Fraundorf et al., 1982b). These element patterns are consistent with the presence of hydrous layer silicates in CS IDPs and the loss of magnesium and calcium by formation of secondary Mg–Ca carbonates on
Figure 12 CI chondrite-normalized element to silicon ratios for CS and CP IDPs. The solid line represents frequency of CS IDPs and the dotted line frequency of CP IDPs. Numbers in upper right of each histogram are the number of CS and CP IDPs, respectively, with element to silicon ratios 43 CI. CS IDPs are systematically depleted in calcium and magnesium while CP IDPs are only slightly depleted in calcium, aluminum, sulfur, and iron relative to CI (vertical dotted line). Source: Schramm et al. (1989).
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IDPs Average Logarithmic average
Compositions the parent bodies. Similar magnesium depletions in the ﬁne-grained matrices of CI (and CM) meteorites also have been attributed to leaching during aqueous alteration. Thus, the magnesium compositions of the smooth group of chondritic IDPs suggest that they too have been processed by aqueous alteration, which is an important clue regarding their origin. Conversely, CP IDPs are not signiﬁcantly depleted in either magnesium or calcium suggesting that they have not been exposed to aqueous alteration. The lack of hydrous minerals in most CP IDPs supports this assertion. The mean Al/Si ratio relative to CI varies among IDPs by an amount similar to that seen for Mg/Si. Again there is a systematic difference between CP and CS IDPs with the latter being enriched in aluminum. CS particles contain secondary layer lattice silicates (clays) with a high percentage of aluminum (Germani et al., 1990). The S/Si ratio shows a large range although there is no systematic difference between CS and CP particles. Although it is depleted from the CI ratio by 30%, it is still higher than any other chondrite group except CIs. Sulfur is the most volatile major element in IDPs and its measured abundance is complicated by the potential for sulfur loss by frictional heating during atmospheric entry and possible contamination of IDPs from stratospheric sulfate aerosols. Because pyrrhotite (FeS) is the major carrier of iron in chondritic IDPs, iron correlates strongly with sulfur (Figure 12). The correlated depletion of iron and sulfur is likely due to exclusion of singlemineral FeS-dominated grains from the data set. Fe/Si is depleted by 20% relative to CI chondrites and there is no signiﬁcant difference in the Fe/Si ratio between CP and CS particles. Iron and aluminum are correlated in CS IDPs but not in CP IDPs. The average Fe/Al value for CS particles is 9.05, which is further from the solar system value of 10.50 (see Chapter 1.03) than the 10.13 Fe/Al mean value for CP IDPs. This same Fe–Al correlation was seen in point count areas analyses of a CS IDP that contains abundant layer lattice silicates. The correlation in CS IDPs almost certainly reﬂects
the abundance of aluminum- and iron-containing layer lattice silicates. The C/Si ratio in chondritic IDPs is systematically higher than all classes of chondritic meteorites. The mean carbon abundance is B10 wt.% versus 3.22 wt.% for CI (see Chapter 1.03). Nitrogen has been detected in chondritic IDPs but as yet not quantiﬁed, although Keller et al. (1995) report that the C/N ratio is approximately chondritic. Electron energy-loss spectra show that nitrogen is carried in amorphous carbonaceous material and that it is heterogeneously distributed as ‘‘hot spots.’’ There is indirect evidence that the nitrogen is associated with PAHs (Section 220.127.116.11).
Most chondritic IDPs have ‘‘chondrite-like’’ trace-element compositions (Arndt et al., 1996). Abundances in individual chondritic IDPs generally scatter from B0.3 CI to B3 CI and enrichments are more common than depletions (Flynn and Sutton, 1992a–c). Volatile elements tend to be enriched relative to CI meteorites (Ganapathy and Brownlee, 1979; Sutton, 1994). Enrichments of bromine measured in some IDPs probably reﬂect stratospheric contamination (Van der Stap et al., 1986; Flynn, 1994a; Flynn et al., 1996), and zinc depletions probably reﬂect loss of zinc by heating during atmospheric entry heating. Some low-zinc IDPs are also depleted in other volatile elements (copper, gallium, germanium, and selenium) (Flynn and Sutton, 1992a; Flynn et al., 1992). The most important trace-element trends in chondritic IDPs are illustrated in Figure 13. Element ratios for two different elements are plotted on the x- and y-axes and the reference lines are where CI-normalized element ratios are 1. Nickel and chromium do not show a trend and are scattered about the CI reference lines and average Cr/Fe and Ni/Fe are similar to CI (Figures 13a and 13b). Calcium is depleted in most of the IDPs in accordance with Schramm et al. (1989) (Figure 13b), and titanium appears to be enriched (Figure 2c).
Figure 13 Trace-element ratios in IDPs. Data from SXRF analyses are plotted on ‘‘three-element’’ diagrams. Element ratios are normalized to bulk CI abundances: (element/Fe)sample/(element/Fe)CI also denoted ‘‘element/Fe/CI.’’ CI composition lies at the point element/Fe/CI ¼ 1 on each plot. Averages, assuming data are normally distributed (open squares) and assuming the data are log normally distributed (open diamonds), are also shown. Plots (a)–(c) exhibit the behavior of some more refractory elements chromium, calcium, and titanium with respect to nickel, while (d) and (e) show the behavior of zinc (relatively volatile) with respect to nickel (relatively refractory) and selenium (relatively volatile). Source: Kehm et al. (2002).
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Volatile trace elements are plotted in Figures 13d and 13e. There are both enrichments and depletions of zinc (Figure 13d). Figure 13e shows the relationship between selenium and zinc. Low-zinc IDPs tend to also have low selenium and selenium deﬁciencies, like zinc deﬁciencies, and are more common in chondritic IDPs than selenium enrichments.
Because a single 10 mm IDP can contain several tens of thousands of submicrometer grains, and the ion microprobe has traditionally measured isotopic composition on a scale of B10 mm, large isotopic anomalies in ‘‘individual’’ grains within IDPs may not be recognized because they were averaged out on a scale of 10 mm. With this caveat in mind, the hydrogen, carbon, nitrogen, oxygen, magnesium, and silicon isotopic compositions of chondritic IDPs have been measured. Esat et al. (1979) ﬁrst measured the magnesium isotopic compositions of four chondritic IDPs as well as calcium in one IDP. They found that the magnesium compositions were very close to normal isotopic composition but their normalized isotopic ratios appeared to show nonlinear effects of 3–4%, which at that time was near the limit of detection. The isotopic composition of calcium was found to be normal (solar) within 2%. Esat et al. recommended that it might prove useful to measure individual B1 mm components of IDPs. Signiﬁcant enrichments and depletions in D/H have been found in IDPs (Zinner et al., 1983; McKeegan et al., 1985; McKeegan, 1987). The carrier of the D/H anomalies is believed to be the carbonaceous matrix within IDPs and, since the highest D/H enrichments (up to B30,000%) are found in large cluster IDPs that likely maintain a thermal gradient during atmospheric entry (i.e., their interiors remain cool), the carrier is presumably organic (Messenger, 2000). Enrichments of 15N (up to d15N ¼ 1,280%) have also been found. Many but not all of the 15N-enriched particles also show D/H enrichments, but the converse is not true (Stadermann et al., 1989). Both the D/H and 14N/15N often vary signiﬁcantly within a given particle, in some cases displaying a pronounced ‘‘hot spot.’’ The D and 15N anomalies have been attributed to organic materials produced by ion–molecule reactions in cold interstellar molecular clouds (Messenger, 2000), although the same processes might equally work in the outer fringes of the solar nebula. PAHs, believed to be important constituents of interstellar molecular clouds, were found
in only two (of seven measured) IDPs that had large D anomalies. D and N enrichments have been observed in both the CP and CS particles. A new type of ion microprobe, the NanoSIMS, has made it possible to measure the isotopic compositions of chondritic IDPs on a scale of 0.5–1 mm (Floss and Stadermann, 2003; Messenger et al., 2003). Six silicate grains identiﬁed by Messenger et al. in nine chondritic IDPs have isotopic compositions conﬁrming their presolar origins. Three of the grains exhibit elevated 17O/16O ratios and solar 18O/16O ratios consistent with origins in red giant and asymptotic giant branch stars, one is 16O rich consistent with formation in a low-metallicity star, and two of uncertain stellar origin are 16O depleted. One of the grains is a forsterite crystal and two others are GEMS. Floss and Stadermann (2003) measured carbon, nitrogen, and oxygen enrichments in two IDPs using the NanoSIMS (Figure 14). Two 17O-enriched presolar grains (of unknown mineralogy) but with isotopic compositions similar to those of the presolar silicates were observed as well in a region with a modest but signiﬁcant depletion in 13C (d13C ¼ 75%) and spatially associated with a nitrogen ‘‘hot spot’’ with d15N ¼ 1,280%. Although hints of depletions of 13C (with large errors) have been reported previously (McKeegan, 1987), the NanoSIMS measurements provide the ﬁrst indication of correlated carbon and nitrogen isotope anomalies.
Rajan et al. (1977) ﬁrst measured the noble gases in chondritic IDPs and found solarwind 4He concentrations comparable to those observed in lunar soil grains. The measured concentrations were consistent with the B104 years calculated exposure ages of small particles in solar orbit. Hudson et al. (1981) measured 20Ne/22Ne in 13 combined IDPs and observed a ratio of 1373, which is within the range of solar-wind neon. Nier and Schlutter (1990) measured 3He/4He and 20Ne/22Ne in 16 individual IDPs. The average helium content was 0.02770.01 cm3 STP g-1 (in the same range reported by Rajan et al., 1977), and the average 3 He/4He ratio of 15 of the IDPs was (2.470.3) 10–4 (one IDP had a 3He/4He ratio of (1.4570.3) 103). But using stepwise heating, Pepin et al. (2000) measured 3He/4He ratios up to 40 the solar-wind ratio in several cluster particles, which they attribute to either cosmic-ray-induced spallogenic reactions during prolonged exposures of the IDPs in space or irradiation of the IDPs on their parent-body
Compositions Benavente (L2036-G16)
19 Benavente (L2036-G16) 1,500 1,400 1,200 1,000 800 600 400 200 0 −200 −300 −600
1,500 1,400 1,200 1,000 +1,280‰
800 600 400 200
0 −200 −300 17O
Image area: 10 × 10 µm2
Image area: 10 × 10 µm2
Figure 14 (a) d O image of portion of Benavente (L2036-G16), showing a 17O-rich subgrain within the IDP. The grain is B300 nm2 in size. The extremely anomalous O isotopic composition indicates that this grain is of presolar origin. (b) d15N image of a portion of Benavente (L2036-G36) showing a strongly 15N-enriched portion of the IDP. The ‘‘hotspot’’ is B0.6 mm 1.8 mm in size. The bulk IDP is also 15N-enriched with an average d15N of about 230%. Data courtesy of C. Floss and F. Stadermann, Washington University.
105 SW Pyx
D10 H21 C3 H14 L20 C21 H8 F32 I3 F24 I24 F12 D3 J13 B19 B1 M20 Plag O16 G8 I14 F9 G14 I8 G19 E15 G13 G11 H25 G2
Figure 15 Noble-gas elemental ratios in IDPs compared with CI meteorites and solar-wind (star) noble gas compositions are also plotted. Closed and open diamonds represent ‘‘unheated’’ IDPs and Zn-depleted IDPs, respectively. Square and circle represent lunar mineral separates (Signer et al., 1977) and planetary bulk CI chondrite (Jeffery and Anders, 1970), respectively. Data courtesy of K. Kehm. See also Kehm et al. (2002).
regoliths prior to their release into the interplanetary medium. The average 20Ne/22Ne value for 10 of the 16 IDPs measured by Nier and Schlutter (1990) was 12.070.3 and the average 21Ne/22Ne value for three of the 16 IDPs was 0.03570.006. Noble gas ratios (4He/36Ar versus 20Ne/36Ar) in 31 IDPs are plotted in Figure 15. Plotted uncertainties are 1s. Also plotted are the noble gas elemental composition of the CI carbonaceous chondrite Orgueil (planetary) and the
solar wind (SW). The observed elemental ratios in Figure 15 indicate that the IDPs contain solar-wind noble gases diffusively fractionated either in solar orbit or by heating during atmospheric entry with helium and neon depleted with respect to argon. (An observed correlation between helium and zinc abundances in the IDPs suggests that it is more likely that helium is lost by frictional heating during atmospheric entry (Flynn and Sutton, 1992a; Kehm et al., 2002).
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Implanted helium is released from IDPs during pulsed stepwise heating in a furnace that mimics frictional heating during atmospheric entry. The helium release proﬁle can be used to estimate the peak frictional heating temperatures and speeds experienced by individual IDPs during atmospheric entry (Nier, 1994). Stepwise heating has been used to distinguish high-speed cometary IDPs from low-speed asteroidal IDPs (Brownlee et al., 1995). In summary, the abundances and isotopic compositions of noble gases in chondritic IDPs are consistent with a solar-wind origin, although fractionations due to cosmogenic spallation reactions and ‘‘secondary’’ diffusive processes are evident (e.g., heating during atmospheric entry). The solar-wind gases are implanted in IDPs during their lifetimes in solar orbit and they may also contain a primordial (pre-accretional) noble gas component.
Chondritic IDPs are an important resource of extraterrestrial materials because they sample a much broader range of primitive solar system bodies than do conventional meteorites and micrometeorites. Technical difﬁculties that limited interest in IDP research have been overcome as a result of rapid advances in microparticle handling and microanalytical instrumentation. The atmospheric entry speeds of IDPs suggest that some hydrated CS IDPs are from asteroids and some anhydrous CP IDPs are from comets (Brownlee et al., 1995). The mineralogy and petrography of CS IDPs clearly indicate that they were derived from hydrous objects where parent-body aqueous alteration occurred. Given their similarity to the ﬁne-grained matrices of CI and CM meteorites, asteroids are the logical sources. In a few cases the mineralogy and petrography of chondritic CS IDPs link them directly to CI- or CM-like hydrous asteroids. The mineralogy and petrography of anhydrous CP IDPs suggest that they are from either anhydrous objects or very low-temperature hydrous objects where parent-body alteration was either minimal or nonexistent. Comets or ‘‘cometlike’’ outer asteroids are the likely sources of CP IDPs. But it is also likely that some CP IDPs are from asteroids and some CS IDPs are from comets. Since studies of IDPs are equivalent to a limited sample return, clariﬁcation of the source(s) of the different classes of IDPs is a high priority of future research. Anhydrous CP IDPs are unique among known natural geological materials in that they are mineralogically heterogeneous at the
nanometer scale, and unique among known meteoritic materials in that they have ‘‘not’’ been subjected to signiﬁcant postaccretional (parent-body) processing. They differ fundamentally from even the most primitive chondritic meteorites and micrometeorites. Some enstatite and forsterite crystals exhibit preserved evidence of condensation from nebular gases. Others have nonsolar isotopic compositions indicating that grain condensation occurred in the atmospheres of other stars. GEMS are perhaps the most enigmatic component of CP IDPs. Although they are cosmically primitive they have been extensively processed by ionizing radiation as free-ﬂoating objects. Material removed by sputtering has been thoroughly mixed and redeposited on grain surfaces producing GEMS with cosmic (chondritic) elemental compositions. The oxygen isotopic compositions of some GEMS establish that they are presolar grains. The ratio of presolarto-solar system components in CP IDPs is unknown. It is possible that some GEMS-rich CP IDPs are relatively well-preserved aggregates of presolar circumstellar and interstellar materials, the ‘‘common stuff’’ of the ISM. Emerging ‘‘nanobeam’’ analytical technologies will undoubtedly play a major role in future interplanetary dust research. Laboratory analytical data from IDPs will increasingly be compared with in situ spacecraft measurements as well as ground-based observational data from dust in space. Comet and asteroid sample return missions like STARDUST, MUSES-C, and Gulliver will undoubtedly provide new insight about stratospheric IDPs (Brownlee, 1996; Zolensky, 2000; Britt, 2003). The STARDUST mission successfully returned a sample of comet Wild-2 to Earth in January 2006. Papers presented at the 2006 Lunar and Planetary Science Conference (Brownlee et al., 2006; Flynn et al., 2006; Ho¨rz et al., 2006; Keller et al., 2006; Sandford et al., 2006; Tsou et al., 2006; Zolensky et al., 2006) described both expected and unexpected results from the preliminary examinations of the dust grains. Enstatite, forsterite and FeNi sulﬁde crystals similar to those found in CP IDPs are also present in comet Wild-2. ‘‘GEMS-like’’ amorphous silicates with metal and sulﬁde inclusions are abundant in the impact tracks in aerogel, but it is unclear at this early stage whether they are relatively pristine GEMS, (impact) modiﬁed GEMS, or melt residues produced during hypervelocity impact and unrelated to GEMS. An unexpected result is the ﬁnding of high-temperature refractory minerals like those found in calcium aluminum inclusions (CAI’s) in meteorites. They include melilite, anorthite, and osbornite (a refractory
References titanium vanadium nitride). Assuming these mineral assemblages are indeed related to CAI’s, one implication is that material formed at high temperatures in the inner nebula was transported outward to the accretion region of comets. In other words, the solar nebula was considerably more turbulent that has been appreciated. There is no doubt that the Stardust samples are a precious new resource of cosmic dust that will be examined using all of the analytical techniques that have been developed for IDPs. The STARDUST samples are already providing insight about the scientiﬁc importance of IDPs collected in the stratosphere.
ACKNOWLEDGMENTS This research is supported by NASA grants NAG5-10632 and NAG5-10696. I gratefully acknowledge discussions with and data from C. Floss, F. Stadermann, and K. Kehm and reviews by G. Flynn and A. Davis.
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Treatise On Geochemistry ISBN (set): 0-08-043751-6