The Earliest Chemical Dust Evolution in the Solar Nebula

Chem. Erde 62 (2002), 1–45 Urban & Fischer Verlag http://www.urbanfischer.de/journals/chemerd

INVITED REVIEW The Earliest Chemical Dust Evolution in the Solar Nebula Frans J. M. Rietmeijer* Institute of Meteoritics, Department of Earth and Planetary Sciences, University of New Mexico, Albuquerque, NM, USA Received: 15. 1. 2002 · Accepted: 24. 1. 2002

Abstract The astrophysical aspects of the evolutionary path from dust to planets have been known for a long time but it was only recently that it could be followed by astronomical observations (Hubble Space Telescope). It was also recently that the mineral and chemical properties of dust around young stars in different stage of stellar evolution could be determined (the Infrared Space Observatory). Dust around these stars, including Vegatype stars that serve as an analog for the solar nebula, was identified as pure Mg-silicates (forsterite; enstatite), Fe-bearing and pure-Fe amorphous ‘silicates’, silica, Fe-metal, Feoxides (possibly Fe-sulfides) amorphous carbon and polycyclic aromatic hydrocarbons. They are the same phases that make up the aggregate and cluster interplanetary dust particles (IDPs) collected in the lower stratosphere. Recent vapor phase condensation experiments showed that the original condensates were mostly amorphous, chemically ordered, metastable eutectic ‘FeSiO’, ‘MgSiO’ and ferromagnesiosilica ‘silicate’ dust from which the observed non-carbon mineralogy could have evolved during hierarchical dust accretion in the solar nebula. The hypothesis of hierarchical dust accretion uses the size distributions for the surprisingly limited number of non-chondritic dust types in aggregate and cluster IDPs as a measure of relative time. It predicts the accretion of gradually larger, relatively younger, dust aggregates with increasingly diverse chemical and mineral properties of increasingly larger crystalline grains that evolved from initially mostly amorphous dust. This early chemical and mineral dust evolution can be traced in the collected aggregate and in larger cluster IDPs and in even larger aggregate meteoroids that burn up during atmospheric entry flash-heating but whereof the resulting meteors contain infor* Corresponding address: Frans J. M. Rietmeijer, Institute of Meteoritics, Department of Earth and Planetary Sciences, University of New Mexico, Albuquerque, NM 87131-1116, USA Tel.: ++1-505-277 4204; Fax: ++1-505-277 8843; e-mail: [email protected] 0009-2819/02/62/01-001 $ 15.00/0

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mation on the chemistry, grain size and texture of the original dust. These aggregate particles were protected against post-accretion, thermal or aqueous dust modification of the original presolar dust and the evolved mineralogy and chemistry during cold-storage inside icy protoplanets such as comet nuclei. The interplanetary dust particles provide ground truth to the properties and modification of the presolar dust in dense molecular clouds wherein stars, such as our sun, were born. Key words: Solar system: early evolution, Solar nebula, Comets, Circumstellar Dust, Interplanetary Dust, Chemistry, Mineralogy

“If one is sufficiently lavish with time, everything possible happens” [Herodotus, 485 – (425–420)]

1. Introduction This quote from the Greek philosopher Herodotus serves as a poignant reminder to a discussion of dust evolving in circumstellar dust nebulae, the solar nebula and icy protoplanets. These are environments wherein thermal energy for dust modification may only be available intermittently, e.g. in comets only during perihelion. The situation is different in meteorite parent bodies where the decay of 26Al provided energy for processing of accreted dust (Kerridge and Matthews, 1988). I submit that interplanetary dust particles [IDPs] from icy protoplanets in the solar system could evolve despite a lack of sustained heat sources because of the unique chemical and mineralogical properties of their smallest dust constituents that were formation via kinetically controlled vapor phase condensation. The IDPs will provide information on the types of dust and processes in the evolving solar nebula but which was overprinted, and partially lost, in the collected meteorites during parent body modification. Just before sunrise and just after sunset we are reminded that the inner solar system is a dusty place when looking at the “Gegenschein” and Zodiacal light that are caused by sunlight reflecting off dust orbiting the sun. Dust with a mass of several nano- or picogram is the most abundant material at a distance of 1 AU from the sun. Cometary meteoroids can be up to several grams to kilograms but they will not survive deceleration in the Earth’s atmosphere (Rietmeijer, 2000a, 2001a). Yet, stony meteoroids, such as those delivering the ordinary and carbonaceous chondrite meteorites, ranging from several grams and up to many kilograms will survive atmospheric entry (Kerridge and Matthews, 1988; Rietmeijer, 2002). Because of its unique density stratification, the Earth’s atmosphere is a very gentle decelerator for micrometer-sized dust (Brownlee, 1985). The peak flash-heating temperature in an IDP during deceleration is a function of its size, density, entry angle and entry velocity (Love and Brownlee, 1991, 1994). It provides a narrow window wherein incoming meteoroids will survive as IDPs without complete melting or evaporation. No IDP escapes thermal alteration and seemingly unheated IDPs will show subtle thermal effects. The slow-down from orbital velocity (km/s) to terminal velocity and atmospheric settling rates (cm/s) results in a concentration factor of 106 for IDPs in the middle and lower stratosphere compared to the incoming flux (Brownlee, 1985). This concentration

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means that the collection of IDPs becomes a simple matter of dragging a flat plate through the lower stratosphere. About ten IDPs, 10–15 mm in size, will be collected on a silicone-oil coated, 30-cm2 inertial-impact, flat-plate collector mounted underneath the wings of a high-flying aircraft travelling at ~800 km/h for 35–40 hours between 17–19 km altitude (Zolensky et al., 1994). This collection mode yields a sample of the total, lower stratospheric, dust loading at the time of collection and individual IDPs need to be extracted from among many different types of stratospheric dust that may dominate the collected dusts. For obvious reasons collection flights are not scheduled shortly after a major volcanic eruption. The NASA curator provides a preliminary identification of the particles on the collectors in the Cosmic Dust Catalogs (Rietmeijer and Jenniskens, 2000, for a complete catalog listing as of mid-2000). A combination of visible light-optical properties (e.g. color, luster), particle morphology (sphere, fragment, aggregate) and qualitative chemical composition is sufficient for a reliable classification of >95% of the collected stratospheric dust into three major categories. These categories are (1) interplanetary dust, (2) natural terrestrial dust (mostly volcanic ash and condensed aerosol) and (3) anthropogenic dust that is typically but not exclusively from space-related activities and aircraft operation (Mackinnon et al. (1982). The stratospheric dust collection offers excellent research opportunities to monitor the residence times of various types of volcanic ash (Mackinnon et al., 1984; Rietmeijer, 1988a) and space-related debris, e.g. solid-fuel rocket effluents mostly from the US Space Shuttle (Rietmeijer and Flynn, 1996) and investigate meteoric dust (Rietmeijer, 2001b).

2. Why study interplanetary dust particles? About 4.6 Gyrs ago the sun formed in a fragment of a dense molecular cloud that contained a lot of dust of unknown age and origins. By the time the process had run its natural course the solar system’s main-sequence star, our sun, was surrounded by the terrestrial planets, asteroids, giant gas planets and leftover presolar dust. In the outer solar system this dust co-accreted with (water) ice particles into icy protoplanets wherein the dust remained unmodified during storage at temperatures of 10 to 50 K until environmental conditions in these protoplanets were changed by orbital perturbations. These perturbations also placed these icy protoplanets in Earth-crossing orbits that made it possible for the embedded debris to enter the Earth’s atmosphere. The surviving IDPs still contain a record of the chemical and mineralogical properties of the presolar common dust, as well as the earliest mineralogical and chemical activity in the solar nebula, prior to asteroid and planet formation. The IDPs represent the transition from interstellar, presolar dust to circumstellar dust. They used to be an oddity among the collected extraterrestrial materials but this situation changed dramatically with the emergence of the new field of astromineralogy that identifies dust formation and mineral evolution around young evolving stars (Bouwman et al., 2001; Meeus et al., 2001, among many others). The transition from dust to planets begins with the formation of a dense, dust-rich, core of a few solar masses inside a giant molecular cloud such as the Orion Nebula. Eventually this rotating core collapses to form a central protostar with a surrounding disk embedded in an envelope of gas and dust that is falling into the protostar that will grow

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and develop bi-polar outflows of dust and gas along the rotational axis of the whole system. At some point the in-fall of dust and gas ceases, the bi-polar outflows stop and in a flattened disc of dust and gas, i.e. a solar nebula, dust accretion will begin the process of planet formation. Protostars with an optically thick disk of gas and dust, so-called young stellar objects, include the pre-main-sequence T-Tauri and Herbig Ae/Be stars. These stars are the final stage before developing into young main-sequence stars, e.g. βPictoris and Vega-type stars, that can serve as analogs for the solar nebula (Bouwman, 2001). For the first time using the Infrared Space Observatory (ISO) data was it possible to identify the mineralogical composition of the dust around Herbig Ae/Be and Vegatype stars. These ‘astro’-minerals are remarkably few in type and simple in composition. I will discuss the connection between ‘astro’-minerals and the oldest dust in aggregate IDPs with a record of dust condensation in circumstellar environments that is preserved in the surviving common presolar dust. This dust became the starting material for the mineralogical and chemical evolution of dust that was in the solar nebula. This evolution is characterized by the formation of increasingly larger dusts with an increasingly complex and diverse chemistry and mineralogy.

3. Interplanetary Dust 3.1. Potential sources Dust is produced by (1) catastrophic collisions among asteroids in the asteroid belt located between Mars and Jupiter, among near-Earth asteroids (NEAs), Earth-crossing and Earth-approaching asteroids and (2) meteoroid impacts on these asteroids and the surfaces of planets and satellites with a tenuous or no atmosphere. The most efficient process to produce IDPs is ice sublimation from active comet nuclei during perihelion, dormant icy NEAs (Rietmeijer, 2000a) and Jupiter-family comets such as the ‘asteroid’ Chiron with its intermittent comet-like activity (Bailey et al., 1994). The infrared (IR) reflectance spectra and albedos for small solar system bodies define distinct classes that can be correlated with specific meteorite classes and groups (Bell et al., 1989). One notable exception is the group of spectral P- and D-class bodies with organic materials and possibly ice at the surface. They occur in the outer asteroid belt (Bell et al., 1989) and include most NEAs, Kuiper belt objects (KBOs), that were formed beyond Pluto at ~ 45 AU, and active Oort cloud comets, e.g. comet Hale-Bopp (Hartmann et al., 1987). Oort cloud comets were formed between Uranus and Neptune and ejected into the Oort cloud, a shell of icy protoplanets that surrounds the solar system at a distance of ~ 104 AU. Many chondritic IDPs are spectroscopic P- and D-class materials with spectral reddening due to the carbon-rich materials in these particles (Bradley et al., 1992, 1996). It links them to the most ‘primitive’ icy-protoplanets in the solar system. In this context, ‘primitive’ means that beyond compaction the accreted dust in the icy protoplanets (Shearer et al., 1998) were not modified by the same aqueous and thermal processes that affected the meteorite parent bodies in the asteroid belt (Zolensky and McSween, 1988; Brearley and Jones, 1998). It cannot be ruled out that aqueous processes will not modify dust in icy protoplanets, such as active comet nuclei. Ruling out this possibility leads to a prejudgment of potential IDP sources and their post-accretion evolution.

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3.2. Asteroidal and cometary dust Dust from the asteroid belt travels in (near) circular orbits in the ecliptic plane with a velocity of 5.5 km/s; less for dust from KBOs (Flynn, 1996). Dust from NEAs and Oort cloud comets in non-circular orbits has a higher orbital velocity, viz. 5.5–11 km/s when the perihelion distance is >1.2 AU and >11 km/s when perihelion occurs at <1.2 AU (Flynn, 1989). Active comets are the principle sources for the annual meteor showers but their high entry velocity generally prohibits dust survival. Yet, some Leonid meteoroids from comet Temple-Tuttle with a record high entry velocity of 72 km/s that are not expected to survive showed solid debris at the end of the meteor trail (Rietmeijer, 2001a). This debris will not retain its original morphology but survive as quenched, molten silicate spheres (Rietmeijer and Jenniskens, 2000). Due to Earth’s gravitational focussing effect, the average entry velocity of asteroidal debris is ~11 km/s while dust entering at >20 km/s is designated as cometary dust (Flynn, 1994; Love and Brownlee, 1991). It was assumed (Sandford, 1987; Bradley, 1988) that fluffy (porous) chondritic IDPs [Figure 1] dominated by anhydrous minerals are cometary debris with high entry velocities. It was further assumed that chondritic IDPs with a smooth surface (Rietmeijer, 1998a, fig 7; Schramm et al., 1989, fig. 1) dominated by layer silicates are asteroidal dust with low entry velocities. Joswiak et al. (2000) measured the size and density of IDPs for which the peak heating temperature was determined via monitoring of the He-release profile during step-wise heating. They found entry velocities ranging from <11 to 29 km/s using the correlation between the entry velocity and peak-heating temperature for IDPs of known density and size (Love and Brownlee, 1994). The anticipated separation into smooth asteroidal and fluffy cometary IDPs was not observed (Joswiak et al., 2000). Classification of chondritic IDPs into porous and smooth particles is clearly not desirable. Rietmeijer (1998a) submitted that the group of chondritic particles only includes aggregate IDPs and non-aggregate IDPs. Minor reorganization of the constituents will collapse a fluffy aggregate into a compact particle with a smooth surface that yet remains recognizable as an aggregate IDP (Bradley, 1988, fig. 7). The transition from an aggregate IDP to a non-aggregate IDP will require significant aqueous or thermal processing, or both. Non-aggregate chondritic IDP will be indistinguishable from hydrated matrix material in the undifferentiated CI and CM carbonaceous chondrite meteorites of asteroidal origins (Tomeoka, 1991; Rietmeijer, 1998a). 3.3. Chondritic aggregate and non-aggregate IDPs The meteor entry velocities in the Meteor Observation and Recovery Program define two groups of asteroidal meteors with average velocities of 17.6 and 23.3 km/s, and two groups of cometary meteors (Halliday et al., 1996). The latter include low-velocity (30 km/s) cometary meteors and (typical) cometary meteors at > 35.5 km/s. The entry velocity (30.5 km/s) for Taurid meteors from comet Encke, an almost extinct short-period comet, is in the low-velocity cometary meteor group that probably also includes many NEAs (Rietmeijer, 2000a). When collisions in the outer asteroid belt remove IR-class P and D objects from their resonance zones in this belt they become Earth-approaching and Earth-crossing asteroids with high, orbital velocities of asteroidal and low-velocity cometary meteors. Thus, IDPs with a comet-like entry velocity might be from low-

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velocity comets and NEAs, some of which, e.g. 2201 Oljato, show intermittent cometlike activity in the form of a coma due to ice sublimation. It is unlikely that much of the solid debris from fast-moving active comets will survive atmospheric entry. This is a good reason to study meteor showers associated with known comets. The meteor trails contain information on the chemistry and structure of the incoming, evaporating debris that is complementary to the information in the collected IDPs (Rietmeijer, 2001a, 2002). Mackinnon and Rietmeijer (1987) concluded that the mineralogy and textures of chondritic aggregate IDP are uniquely different from any materials in our meteorite collections implying that these IDPs come from different parent bodies than meteorites. Small bodies, such as meteorite-dropping meteoroids, when ejected from the outer asteroid belt will feel the gravitational pull by the giant planets, in particular Jupiter, which prevents most of them from reaching the inner solar system. But low-mass dust particles from the outer asteroid belt will not be gravitationally perturbed. Their orbit, and in fact all dust orbits in the solar system, is affected by Poynting-Robertson drag caused by the solar wind. In this manner, dust-sized meteoroids from all sources slowly spiral toward the sun and ultimately merge with zodiacal dust cloud. For a 10-µm sized IDP this average solar system sojourn time is 104 to 105 years (Brownlee, 1985; Flynn, 1996). Interplanetary dust in space is a rather representative sample of all dust-producing solar system bodies. But because of gravitational focusing effects and thermal interactions with the atmosphere, the collected IDPs are a biased sample of (unmelted) dust particles ranging from ~10 to ~50 µm in size, including low-density aggregate and cluster IDPs, with entry velocities <35 km/s (Rietmeijer, 2000a).

4. Assumptions 4.1. Hierarchical dust accretion Rietmeijer (1998a) formulated the hypothesis of hierarchical dust accretion that will serve as a framework to discuss the chemical and mineral evolution of dust during IDP accretion. I envision an epoch characterized by cycles of presolar dust accretion, presolar dust modification in (very) small protoplanets, fragmentation of the protoplanets and introduction of processed dust among the solar nebula dust population that included original presolar dusts and presolar dust processed in the nebula (e.g. by irradiation processes). The processed dust might also have included condensed dust from vapors generated in the inner part of the solar nebula during a T-Tauri phase and refractory evaporation residues. Repeated occurrences of these cycles ultimately depleted the original presolar dust although some dust with unique isotopic compositions has survived.

Fig. 1. Scanning electron microscope image of a porous (fluffy) chondritic aggregate IDP collected in the lower stratosphere placed on a nucleopore-filter (background) showing platy silicate grains embedded in a matrix of (partially fused) principal components. Its energy dispersive spectrum at the bottom shows approximately chondritic proportions for the elements. A high bremsstrahlung background indicates the presence of carbonaceous matter. Please note that the Cu-peak is an instrumental artifact from spurious X-rays emitted from materials inside microscope. Courtesy of the National Aeronautics and Space Administration, Particle W7029B13 (NASA number S-8227575)

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The earliest protoplanets were centimeter-sized to meter-sized protoplanets (‘pebbles’) with variable dust to ice grain ratios wherein dust evolved into larger grains. Grain growth is evident around Herbig Ae/Be stars where dust grains >50 µm and up to ~1 mm dust grains included Mg-silicates (forsterite, enstatite), amorphous Mg,Fe-‘silicates’, Fe-oxide (or Fe-metal), amorphous and crystalline silica when C/O < 1 and amorphous carbon, graphite(?) and polycyclic aromatic hydrocarbons (PAHs) when C/O > 1 (Molster et al., 1999; Bouwman et al., 2001; Meeus et al., 2001). The mass fraction of forsterite and enstatite crystals increases with heliocentric distance into the cooler regions of some of these stars which might be the first evidence for dust processing in cometsized bodies around young stars (Bouwman, 2001). Using this new data Nuth et al. (2000a) proposed that the relative ages of comets around young stars can be traced by differences in the relative proportions of crystalline to amorphous Mg-silicates and the proportions of the volatile and dust components among comets. 4.2. Icy protoplanets Aggregate particles are debris from icy protoplanets from beyond the ‘snow line’ that was located somewhere in or close to the outer asteroid belt in the solar nebula. The ‘snow line’ is the isotherm in the solar nebula where water vapor condensed to ice. There is no fundamental distinction between comet nuclei and icy asteroids. These icy protoplanets are mostly rubble piles of (sub)meter-sized pebbles (Beech and Nikolova, 1999) and boulders of several hundreds of meters to kilometers in size held together gravitationally, dirty ice (‘Whipple glue’) or both (Gombosi and Houpis, 1986). They probably span the full range of homogeneous ‘dirty snowballs’ to ice-poor agglomerations of boulders and pebbles. They will be unpredictable mixtures of these end-members determined by the accretion history and locations and post-formation evolution, e.g. the number of perihelion passages. Comet-like behavior is an accident of orbit that allows icy protoplanets to reach the solar system inside Jupiter’s orbit where insolation is sufficient for ice sublimation. Their debris includes IDPs, pebbles and boulders that could deliver CI meteorites (Campins and Swindle, 1998; Rietmeijer, 2001a, 2002). The smallest dust constituents in the collected aggregate IDPs are as close as we can get to the properties of the original dust in the dense molecular cloud fragment. 4.3. Atmospheric entry flash heating The original properties of this > 4.56 Ga-old dust might be modified during residency in icy protoplanets but dynamic pyrometamorphism on atmospheric entry is the most dramatic, thermal event that affected all collected IDPs. The maximum temperatures reached during this final, 5 to 15 s flash heating event are between ~300 °C and ~1,000 °C, with an average of about 750 °C, well below the melting point at ~1,500 °C for chondritic IDPs (Rietmeijer, 1998a; Joswiak and Brownlee, 1998; Joswiak et al., 1998, 2000). Thermal modification of pre-entry minerals and chemical composition in unmelted IDPs includes localized melting and mass loss. Complete melting leads to the formation of quenched-melt, Mg,Fe(±Ca)- and Mg,Ca,Al-silicate, and Fe,Ni,S- and Fespheres (Rietmeijer, 2000a). The spheres have unmelted counterparts among the nonchondritic silicate and sulfide IDPs (Mackinnon et al., 1982). Most spheres are not flashmelted aggregates but melted non-chondritic dust fragments with proper ratios of frag-

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ment to sphere diameters predicted by ablation models for atmospheric entry of IDPs (Rietmeijer and Jenniskens, 2000). The spherical IDPs could be surviving non-chondritic debris from high velocity (>35 km/s) icy protoplanets. In this review I will not discuss spherical dust particles but concentrate on the chemical and mineral properties of surviving dust whereby I will occasionally refer to thermal modification when relevant to the flow of this presentation.

5. Additional Information Sources There are excellent reviews describing the collection and curation procedures, the physical, mineralogical and chemical properties, including isotopic compositions and the extraterrestrial signatures. They also include images of IDPs and their constituents. Mackinnon and Rietmeijer (1987) and Zolensky et al. (1994) discuss the collection and curation procedures. For details on the general physical IDP properties I refer the reader to Brownlee (1978, 1985) and Sandford (1987). For the petrology and mineralogy I recommend the papers by Mackinnon and Rietmeijer (1987) and Rietmeijer (1992a, 1998a). Flynn (1994), Flynn et al. (1996) and Thomas et al. (1996) offer bulk particle and trace element chemistry, including isotopic compositions. While listed for specific information most of these papers also discuss other aspects of IDPs. Rietmeijer (1998a) provides a complete bibliography of original papers prior to 1998.

6. Analytical Techniques The techniques currently used for laboratory analyses of IDPs collected in the Earth’s lower stratosphere are listed as of December 2001 in Table 1. The review papers by Sutton (1994), Zolensky et al. (2000), Stephan (2001) and Rietmeijer (1998a) offer general information and details. The other references are mostly first-reported applications of a particular technique. Different papers and abstracts may refer to the same basic technique, e.g. infrared spectroscopy, but with different names specifying the uniqueness of the particular technique that was used. It is noteworthy that the most references are mostly abstracts. It reflects an explosive development in analytical techniques with unprecedented spatial and mass resolution.

7. Classification of Interplanetary Dust Particles 7.1. Particle chemistry Among many other authors, Anders and Grevesse (1989) noticed a strong correlation between elemental abundances in the solar photosphere and in the hydrated, carbonaceous (type I) meteorites Orgeuil and Alais. The highly volatile elements H, He, C and N are higher in the solar photosphere than in these altered meteorites. The CI, or chondritic, composition is accepted as the solar nebula bulk composition. Thus, a CI composition for aggregate IDPs indicates an extraterrestrial origin because aggregates of terrestrial crustal dust are most unlikely to obtain this composition by random

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Table 1. Techniques currently used for the chemical analyses of interplanetary dust particles. References Light elements Hydrogen Li, Be, B isotopic compositions

TOF-SIMS; EELS SIMS; TOF-SIMS

Elemental Carbons, organic carbons, hydrocarbons Mapping of bonding states [C-O; C-H, STXM; C-XANES; O-XANES C-C]; identification amorphous and poorly EELS; TEM graphitized carbons, organic carbons Crystallographic ordering amorphous carbon IR microscopy; Raman spectroscopy; to graphite TEM Hydrocarbons constituent (H, C, O, N, P) species; polycyclic aromatic hydrocarbons LM/MS; µL 2MS and alkylated derivatives Structural evolution of hydrocarbons to Raman spectroscopy; C-XANES elemental carbons C/O ratio Combined C-XANES/O-XANES C-H feature of aliphatic hydrocarbons TOF-SIMS; IR micro-spectrometry Synchrotron FTIR; C-XANES; C-H2 stretching modes; C=O bonds of organic carbonaceous matter TOF-SIMS Nitrogen EELS; TOF-SIMS Stable Light Element isotopes Quantitative measurements and mapping for H, C, N, O, Mg

2, 21 17, 21, 22

8, 25

16, 17, 25

5, 9, 25 8, 25 8 4, 21 6, 8, 21 21, 25

Ion microprobe; TOF-SIMS; NanoSIMS

20, 21, 25

Step-heating mass spectrometry; TOF-SIMS; laser gas-extraction Synchrotron X-ray microprobe

10, 15, 21

Electron microprobe; AEM, TOF-SIMS

17, 21, 22, 25

10

Na to Br, including REE Most elements between S and Mo Halogens

Laser gas extraction Synchrotron X-ray microprobe PIXE; INAA SXRF TOF-SIMS

24, 25 22, 25 21

Surface chemistry Sulfur, Chlorine Halogen salt phases

Scanning Auger microprobe TEM

12 17

Minor element mapping and abundances K, Ca in thin sections atomic number (26) e.g. Fe, Ni, Zn, Br, Sr

STXM K and Ca edges SXRF microtomography

8, 25 23

Noble gases Abundances and isotopic compositions of He, Ar, Ne

Major element composition Major (rock-forming) elements Na through Ni Minor and trace element composition Cu, Zn, Ga, Ge, Se

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Table 1. Continued. References Mineralogy Identification of minerals and amorphous phases; crystallographic textures Compositional states amorphous materials, e.g. oxidation states Large Fe-sulfide/oxide IDPs IDP classification Constituent identification in ultrathin sections Physical properties Surface imaging, e.g. mineral growth features Grain size, shape, textural relationships Porosity Density (combined techniques)

TEM; FEG-TEM; HRTEM; AEM; Electron diffraction, XRD EELS; ELNES, C-XANES

17, 22, 25

Synchrotron X-ray Diffraction SEM, EDS, FTIR; IR microspectrometry Micro-FTIR, C and O XANES

7 1, 19, 25

AFM TEM; SEM; light optical imaging Imaging techniques Weighing, ATEM, SynchrotronXRD, neutron activation ananlysis

8, 25

3, 8, 11, 13

14, 18 17, 25 17, 25 17, 25

Abbreviations: AEM – Analytical TEM and SEM (ASEM) by energy dispersive spectrometry for all elements from carbon to high atomic number depending on detection limits as a function of incident electron energy and the use of window-less, thin or thick window energy dispersive X-ray detectors. AFM – Atomic force microscope C-XANES – Carbon X-ray near edge structure EDS – Energy dispersive spectroscopy EELS – Electron energy loss spectroscopy ELNES – Electron energy-loss near-edge structure FTIR – Fourier transform infrared spectroscopy INAA – Instrumental neutron activation analysis LM/MS – laser microprobe/mass spectrometry mL2MS – microprobe two-step laser mass spectrometry SEM – Scanning electron microscope TEM – Transmission electron microscope, including field-emission gun (FEG) TEM and high resolution TEM (HRTEM) PIXE – Proton-induced X-ray emission SIMS – Secondary ion mass spectrometry, including Time-of-flight (TOF) SIMS STXM – Scanning transmission X-ray microscopy SXRF – Synchrotron X-ray fluorescence XRD – X-ray diffraction References: (1) Borg et al. (1998); (2) Bradley (1994a); (3) Bradley et al. (1999); (4) Brownlee et al. (2000); (5) Clemett et al. (1993); (6) Flynn et al. (1998); (7) Flynn et al. (2000); (8) Flynn et al. (2001a); (9) Gibson Jr.(1992); (10) Kehm et al. (1998); (11) Keller et al. (2000a); (12) Mackinnon and Mogk (1985); (13) Molster et al. (2001); (14) Nakamura et al. (2000); (15) Nier and Schlutter (1992); (16) Raynal et al. (2000); (17) Rietmeijer (1998a); (18) Romstedt (1998); (19) Sandford (1987); (20) Stadermann (2001); (21) Stephan (2001); (22) Sutton (1994); (23) Sutton et al. (2000); (24) Wallenwein et al. (1987); (25) Zolensky et al. (2000)

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sampling. While it is helpful to recognize collected extraterrestrial dust, a non-chondritic dust composition cannot be used to exclude this particular origin. Other, mineralogical and chemical, indicators for an extraterrestrial particle origin, such as solar flare tracks in silicate minerals, and a cosmic Fe/Ni 20 ratio are reviewed by Rietmeijer (1998a). Stable isotope compositions, e.g. D/H and 15N/14N ratios (Messenger, 2000) identify surviving presolar dust in collected IDPs. The He, Ne and Ar abundances and isotopic compositions also prove an extraterrestrial particle origin and trace solar system histories that could be unique to aggregate IDPs (Pepin et al., 2001). An individual IDP may show any combination, but almost never all, of possible extraterrestrial signatures. In some particles only one diagnostic feature may survive atmospheric entry flash heating (Rietmeijer, 1998a).

7.2. Particle morphology A descriptive IDP classification uses both chemical composition (chondritic or nonchondritic) and particle morphology (aggregate or non-aggregate) (Figure 2). Chondritic aggregates and non-chondritic fragments are found side by side as individual particles on the flat-plate collectors. Non-chondritic fragments and (rare) refractory aggregate IDPs often have chondritic aggregate material at their surface as proof that they are extraterrestrial. Independent confirmation of an extraterrestrial origin for fragments is often lacking but the oxygen isotopic compositions confirm this origin for refractory aggregates (McKeegan, 1987). This IDP classification is desirable over earlier classifications that introduced genetic factors by combining primary and secondary dust properties (Rietmeijer, 1998a). This classification (Figure 2) highlights an important property of dust from icy protoplanets, namely that it is characterized by a surprisingly simple mineralogy for constituents < 100 µm in size in aggregate and cluster IDPs (Table 2) Brownlee (1978) noticed that a cluster IDP would break apart during inertial-impact in the collectors thereby exposing its individual non-chondritic fragments and aggregate-IDP constituents. These particle impacts appear as optically dark, isolated areas of very high dust concentration on a collector (Rietmeijer, 1998a, fig. 25). These macroscopically distinct debris clusters led to a proposed minimum cluster IDP size of ~ 60 µm. The simplest interpretation of the interrelationships among the particles is that cluster IDPs are agglomerates of chondritic aggregate IDPs and non-chondritic IDPs. But, when individual aggregate IDPs and (fragmented) cluster IDPs occur on the same collector it raises the question whether they are debris from two fundamentally different sources or whether they reflect two different windows of opportunity in the survival of fractal aggregate particles of different size and density (Rietmeijer, 1998a). The total exposure time (35–40 h) for most collectors was in reality accumulated during individual slots over a period of a few weeks or even months. During these periods debris from many sources, such as sporadic meteors (Zodiacal dust) and annual meteor showers, are commingled at the collection altitudes (Rietmeijer, 2000a). This important issue might be addressed by targeted stratospheric dust samplings of only a few hours using the large area collectors or by near-nucleus analyses and collection of ejected dust by missions to active comets (the Stardust and Rosetta missions) or comet nucleus sample return missions.

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Fig. 2. Chemical and morphological classification of interplanetary dust particles collected in the Earth’s lower stratosphere (columns 2 and 3) and mineralogical identification of aggregates and fragments (column 4). Modified after Rietmeijer (2000b, 2002).

Table 2. Chondritic and non-chondritic IDPs, on average 10–15 µm in size, from icy protoplanets as collected in the Earth’s lower stratosphere (Rietmeijer, 1998a, 1999a). Chondritic Aggregate IDPs

Non-chondritic IDPs (chondritic aggregate material adheres to surface)

A matrix of principal components (100– ~ 1,000 nm) with embedded variable amounts of ~5 µm-sized Mg,Fe- and Ca,Mg,Fe-silicates, Ni-free and low-Ni pyrrhotite, iron oxides, and amorphous materials (1) Silicate IDPs, mostly Mg,Fe-silicates and Mg(Fe),Ca,Al-silicates (2) Sulfide IDPs, mostly Ni-free and low-Ni pyrrhotite; rare pentlandite (3) Refractory, Ca,Ti,Al-rich IDPs (4) (rare) plagioclase-like IDPs (5) IDPs that are admixtures of variable amounts of IDPs types 1–4 Cluster IDPs ~60–100 µm in size are mixtures of variable amounts of aggregate IDPs and nonchondritic IDPs (1) – (5). Cluster IDPs also contain Fe-oxide fragments (Thomas et al., 1995) that could be original phases or oxidized Fe-sulfides during atmospheric entry.

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8. Petrographic Evidence for Hierarchical Dust Accretion 8.1. Mineralogy 8.1.1. Aggregate IDP matrix

Aggregate IDPs consist of a matrix aggregates with embedded grains that are both ~5 µm in size. The matrix consists of unique, generally spherical entities, known as principal components (PCs) that are ~100 to ~1,000 nm in diameter (Figure 2; Table 2), 1. Carbonaceous PCs of refractory hydrocarbons, amorphous (often-vesicular) carbons, poorly graphitized and pre-graphitic turbostratic carbons. They are often fused into contiguous patches and clumps (Thomas et al., 1993, fig. 2), 2. Carbon bearing ferromagnesiosilica PCs of ultrafine (2 to ~50 nm) platy Fe,Mgolivines, Fe,Mg-pyroxenes, Fe,Ni-sulfides, Fe-oxides embedded in a carbonaceous matrix of refractory hydrocarbons and amorphous carbons. They contain sulfur and minor Al, Ca, Cr, Mn, Ni, and traces of phosphorous and zinc (Rietmeijer, 1989, figs. 2, 3), and 3. Ferromagnesiosilica PCs (Rietmeijer et al., 1999a, fig. 9; Rietmeijer, 1998a, figs. 32, 34; Bradley, 1994a, fig. 1B): (I) Coarse-grained PCs with a (Mg,Fe)6Si8O22 (smectite dehydroxylate) bulk composition and Fe/(Fe+Mg) (fe) = 0 – 0.33. These Mg-rich PCs consist of a coarse-grained (10 – 410 nm) Mg,Fe-olivine and Mg,Fe-pyroxene plus amorphous aluminosilica material (± traces Ca, Mg, Fe), (II) Ultrafine-grained PCs with a (Mg,Fe)3Si2O7 (serpentine dehydroxylate) bulk composition and fe = 0.3 – 0.83 (modal fe = 0.67). The Fe-rich units consist of an amorphous matrix with embedded Mg,Fe-olivine, Mg,Fe-pyroxene, Fe,Ni-sulfides, magnetite, and metal (kamacite) grains (<50 nm). Magnetite is presumably a product of terrestrial oxidation; sulfur contents are variable. Matrix domains of mostly Fe-rich ferromagnesiosilica PCs occasionally have associated grains, <100 to ~500 nm in size, of Mg,Fe(Ca)-silicate and Fe,Ni-sulfide of unknown origin. They could be an accretion feature or in situ post-accretion dust processing. The matrix is primarily a mixture of carbonaceous PCs, C-bearing ferromagnesiosilica PCs and (pure) ferromagnesiosilica PCs. The composition, size and (estimated) mass of the PCs resemble the particles in the coma of comet Halley (Rietmeijer, 1998a) wherein they were identified as (1) CHON (carbon, hydrogen, oxygen, nitrogen), (2) mixed carbon-silicate and (3) silicate particles (Jessberger et al., 1988; Fomenkova et al., 1992). These similarities offer a direct link between matrix aggregates and dust in active comets or in general dust in the dirty-ice of icy protoplanets. 8.1.2. Embedded non-chondritic grains

These grains include various silicate minerals: silica (SiO2; tridymite) (Figure 3), olivine [(Mg,Fe)2SiO4], Ca-free Mg,Fe-pyroxene [(Mg,Fe)2Si2O6, mostly enstatite, Mg2Si2O6], Ca-pyroxenes [Ca(Mg,Fe)Si2O6, viz. pigeonite, Ca < (Mg,Fe), and diopside, Ca (Mg,Fe)] and Ni-free and low-Ni Fe-sulfides (mostly pyrrhotite, Fe 7S8) and low-Ni Fe-metal (kamacite) grains (Rietmeijer, 1998a,1999a). In what could to be a subtype of aggregate IDPs enstatite occur as elongated crystals such as whiskers, rods and blades. Micrometer-sized amorphous ‘silicate’ grains in aggregate IDPs have a (1) plagioclase (NaAlSi3O8 – CaAl2Si2O8), (2) alkali-feldspar [(Na,K)AlSi3O8] or (3) Mg,Fe-bearing

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Fig. 3. Scanning electron microscope image of a large euhedral SiO2 crystal embedded in the matrix of aggregate IDP W7010*A2 (see, Rietmeijer, 1989). Reproduced by courtesy of D.S. McKay (NASA/JSC, Houston).

aluminosilica composition [Figure 4]. The irregular shape suggests coagulation and fusion of aggregates of PCs of the same type. A similar process formed the contiguous patches of carbonaceous material with embedded C-free grains in carbon-rich aggregate IDPs (Thomas et al., 1993). In this case, they represent fused matrix aggregates dominated by CHON-type PCs. These grains of processed PCs represent the first evidence for dust processing via Ostwald ripening to lower the surface free energy. Within petrological context of the host aggregates it is unlikely but not impossible that the fused amorphous grains are due to localized melting during atmospheric entry flash heating. 8.2. Grain size 8.2.1. Aggregate and cluster IDPs

The size of the constituent non-chondritic dusts (Mg,Fe olivine, Mg,Fe(±Ca)pyroxene, Fe,Ni-sulfides and oxides) vary in a non-random manner. The typically nanometer-sized grains in matrix aggregates (~5 µm in size) are mixed with the PCs that are ~100 nm in size and larger. In aggregate IDPs (10–15 µm in size) these grains are on average ~5 µm in size, while in cluster IDPs (>60–100 µm in size) these non-chondritic

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Fig. 4. The Al2O3 – SiO2 diagram showing two distinct compositional fields for amorphous Mg,Fe-bearing aluminosilica grains in aggregate IDPs (from Thomas et al., 1989) plus the compositions for non-stoichiometric plagioclase (triangles), alkali-feldspar (diamonds) and layer silicates (squares)in five other aggregate IDPs. The compositions of amorphous aluminosilica ‘restite’ in Mg-rich ferromagnesiosilica PCs after crystallization of olivine and pyroxene would occupy lowAl2O3 portion of this diagram for SiO2 >90 wt. %. Modified after Rietmeijer (1992a, 1998a)

grains are >10 µm in size and mixed with aggregate IDPs (Table 2). I interpret the systematic recurrence of the same dust types with increasingly larger size to support a sequence of hierarchical dust accretion that started with the PCs into matrix aggregates to aggregate IDPs and to cluster IDPs. The relative proportions of the constituents are the result of temporal and spatial variations of their availability in the solar nebula accretion regions. Hierarchical dust accretion with dust recycling through the nebula and very small (icy) protoplanets will, as accretion proceeded, produce protoplanets that contain increasingly larger dust grains that are increasingly more mineralogically and chemically diverse and with the initially amorphous dust steadily decreasing in favor of crystalline dust. 8.2.2. Super clusters

Hierarchical accretion predicts super clusters that will be mixtures of cluster IDPs and non-chondritic dusts >60–100 µm in size, eventually the formation of larger giant clusters that are mixtures of ‘super clusters’ and commensurately sized non-chondritic fragments, and so on. In fact, pebbles in icy-protoplanets might be accretions of very large clusters after mechanical rearrangement of constituent dusts when pore spaces collapsed after ice sublimation. The pebbles and boulders might be anhydrous proto-CI material or hydrated CI material, which is consistent with the meteor studies that report incoming, structurally weak CI-like meteoroids with cometary orbits (Rietmeijer, 2000a). It seems possible that the 50 to 150 µm-sized Fe,Ni-sulfide grains and the Mg-rich olivines, 50 to

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~500 µm, in the matrix of undifferentiated meteorites represent the youngest step in the accretion sequence. If so, the hypothesis provides a conceptual framework for the interrelationships among meteorite matrix, very small to comet-like icy protoplanets, solar nebula and the common presolar dust (Rietmeijer, 2001a). Is there evidence for incoming super and giant aggregate IDPs that could not survive atmospheric entry flash-heating when the meteoroid size-density product (Love and Brownlee, 1991) is a key controlling parameter? 8.3. Meteors The answer to this question may be obtained from the observations of meteor showers. Aggregate IDP density ranges from 0.4 to 3.4 g/cm3 (mean = 2 g/cm3; see Rietmeijer, 1998a). The mass of super and giant clusters is increasingly determined by the massive Mg-rich silicate (3.2–3.5 g/cm3) and Fe-sulfide (pyrrhotite: ~4.6 g/cm3) fragments. Their size-density product becomes unfavorable for survival compared to cluster IDPs with the same entry parameters and for collapsed pebbles the probability of survival is further diminished. Rietmeijer (2000a) submitted that the window of opportunity for unmelted survival of meteoroids shows a gap between cluster IDPs and cm-sized meteorites. This gap would include the hypothesized clusters that become meteors and fireballs during deceleration in the atmosphere. The resulting light curves contain information on meteoroid size, texture and chemistry. A continuous light curve is an indication for a texturally homogenous body while a “humped” light curve supports a heterogeneous, composite meteoroid with a dustball component and a massive grain. Rietmeijer (2001a) calculated that a Leonid meteor with two components of (2.4–2.5) × 10–4 g (Murray et al., 2000) is a composite meteoroid with a 575 µm-sized dustball (ρ = 1 g/cm3) and a massive, 525 µm-sized, Mg-rich silicate grain (ρ = 3.3 g/cm3) or a 460 µm-sized pyrrhotite grain. This meteoroid might be a giant cluster. It suggests that millimeter-sized aggregates can be found among the dust in comet Temple-Tuttle. Continued analyses of shower meteors (Jenniskens et al., 2000) will reveal similarities and differences among solid debris from known comets and other small bodies as well as finding that this debris has similar and complementary chemical and physical properties to those of aggregate IDPs (Rietmeijer, 2001a).

9. Origins of Dust 9.1. Solar nebula dust condensation Cosmochemistry embraces the axiom that all solids initially formed by fractional condensation from a cooling, hot (up to 2000 K) chondritic gas that produced a sequence of stoichiometric minerals (Grossman and Larimer, 1974, for a review) albeit that in recent studies the requirement for a chondritic gas composition was relaxed. This condensation model is inconsistent with the meteorite record that has yet to produce unambiguous evidence for equilibrium condensation. Moreover, the sizes of meteoritic minerals believed to have condensed are too large to be primary condensates. I submit that solar nebula condensation was not a defining process for the constituents of aggregate and cluster IDPs that represent dust from the outer solar nebula where the tempera-

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ture was ~10 K during the transition from presolar, molecular cloud dust to solar nebula dust. The formation of principal components is the key to appreciate the chemical and mineral dust evolution. The PCs have a non-chondritic bulk composition, which is also true for all other dusts in matrix aggregates, aggregate and cluster IDPs. Ironically or significantly, depending on one’s point of view, the petrology of aggregate and cluster IDPs support that accretion was limited to a small number of different, non-chondritic dusts. In fact, a bulk CI composition of aggregates arose only some time after solar nebula dust accretion had begun. It doesn’t mean that solar nebula condensates might not be present in these aggregate particles. For example, the elongated enstatite single-crystals (Bradley et al., 1983) show that vapor phase condensation occurred in aggregate IDP parent bodies. Morphologically similar wollastonite (CaSiO3; a pyroxenoid mineral) whiskers occur in cavities inside the Allende carbonaceous chondrite meteorite (Miyamoto et al., 1979). This unique morphology requires an active screw dislocation during growth anchored on a surface such as a cavity wall. Bradley et al. (1983) interpreted the enstatite whiskers as evidence for equilibrium solar nebula condensation. They are a secondary condensation feature and grew inside cavities in aggregate IDP parent bodies, e.g. consolidated ice-free pebbles. High temperatures necessary to produce a vapor that could condense enstatite whiskers would probably require large ice-free protoplanets such as asteroids. The subtype of aggregate IDP with whiskers may be asteroidal debris. A search for the properties of condensed solids in extraterrestrial materials should focus on the smallest grains, i.e. the PCs in the matrix aggregates that were formed far removed from the regions of vapor phase condensation in the inner part of the solar nebula. 9.2. Principal components 9.2.1. Processed vapor phase condensates

Two of the three PCs are chemically pure solids but are they ‘pristine’ or processed dust? The chemically and crystallographically diverse dust >50 µm and up to ~1 mm in size around young stars (Molster, 2000; Bouwman et al., 2001; Meeus et al., 2001) is clearly larger than the individual PCs and the amorphous ‘silicate’ grains of fused PCs in aggregate IDPs. It is unknown if different ‘astrominerals’ co-occur in polymineral grains such as the coarse-grained ferromagnesiosilica PCs and the carbon-bearing ferromagnesiosilica PCs. The latter are mixtures of carbonaceous and ferromagnesiosilica dusts that are not to be confused with ferromagnesiosilica PCs enclosed in contiguous patches of carbonaceous materials in some aggregate IDPs (Bradley et al., 1999; Rietmeijer, 1998b). The unfortunate but unavoidable sampling bias in the collected aggregate and cluster IDPs does not allow us to follow the process of grain fusion to sizes comparable with those of astro-minerals. Laboratory dust condensation experiments may be limited to carbon and ‘silicate’ vapors. I am unaware of dust analog experiments using chemically complex C,H,O,Nvapors that would be relevant astrophysical and cosmochemical carbonaceous matter that also contain light elements such as D/H ‘hot spots’ (McKeegan et al., 1985), 15N/14N domains (Messenger, 2000) and oxygen (Flynn et al., 2001a). Condensed carbons from pure carbon and C-H2 vapors included buckyballs, buckytubes, amorphous and poorly graphitized carbons and graphitic ribbons (Rotundi et al., 1998). Non-equilibrium carbon vapor condensation was evidently kinetically controlled.

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9.2.2. Mg-rich ferromagnesiosilica PCs • 9.2.2.1. ‘Silicate’ vapor condensation experiments: Kinetically controlled, non- equilibrium condensation experiments showed that gas to solid condensation in vapors with Al-SiO-O2-H2, Mg-SiO-O2-H2 and Fe-SiO-O2-H2 compositions is not chaotic. The condensed mixed ‘AlSiO’, ‘FeSiO’ and ‘MgSiO’ grains (<<50 nm in diameter) have distinct, chemically ordered compositions. This behavior in the binary systems Al2O3 – SiO2 (Rietmeijer and Karner, 1999), MgO – SiO2 (Rietmeijer et al., 2002a) and the pseudo-binary system FeO/Fe2O3 – SiO2 (Rietmeijer et al., 1999b) is possible because they have at least two eutectic points to constrain a metastable eutectic point. The mixed grain compositions invariable match those of the metastable eutectics that are possible in these binary phase diagrams (Nuth et al., 1998, 2000b; Rietmeijer et al., 2000). These invariably amorphous, chemically ordered, metastable eutectic solids are ‘dissipative’ structures defined by the efficiency of thermal energy dissipation in the condensation experiments. Pure oxide grains typically co-condensed with the metastable eutectic grains. The invariably crystalline oxide grains are tridymite (a high-temperature SiO2 polymorph), Fe3O4 (magnetite and γ-Fe2O3, maghémite) and periclase (MgO), but corundum (Al2O3) was not found (Rietmeijer and Nuth, 2000a). A similar kinetically controlled condensation experiment in a Mg-Fe-SiO-H2-O2 vapor showed that none of the solids, despite a “Mg,Fe,SiO” gas composition, had mixed ternary compositions (Figure 5) but only pure end-member and the same metastable eutectic, magnesiosilica

Fig. 5. Ternary diagram MgO-FeO-SiO2 (oxide wt. %) with the compositions (open squares) of condensed metastable eutectic ‘MgSiO’ (Rietmeijer et al., 2002a) and ‘FeSiO’ (Rietmeijer et al., 1999b) ‘silicate’ dust and the crystalline oxides. Despite a ternary metal-oxide vapor composition (filled circle) kinetically controlled condensation in Mg-Fe-SiO-H2-O2 vapors does not produce mixed, ferromagnesiosilica condensates. Modified after Rietmeijer et al. (1999a).

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and ferrosilica compositions of the binary systems (Rietmeijer et al., 1999a). The lack of condensed ferromagnesiosilica solids is related to the complete mixing along the MgO – FeO/Fe2O3 side of the ternary diagram where there will be no preferred metastable eutectic compositions. It follows that the Mg-rich ferromagnesiosilica PC compositions cannot be obtained by condensation alone. • 9.2.2.2. Post-condensation agglomeration: Although they are among the smallest constituents in aggregate IDPs they are not original circumstellar dust condensates. Condensed dust processing had occurred prior to the accretion of aggregate IDPs. It is unknown whether this proposed dust processing occurred in the molecular cloud or in the solar nebula. The Mg-rich ferromagnesiosilica PCs compositions that are narrowly defined in the ternary diagram Mg-Fe-Si (element wt %) (Figure 6) resulted from agglomeration of nanometer-sized magnesiosilica and ferrosilica condensates along experimentally defined mixing lines (Rietmeijer et al., 1999a). In summary,

Fig. 6. The Mg-Fe-Si (element wt %) diagram with the mixing lines (dashed) connecting the amorphous, metastable eutectic condensed dust compositions (filled circles) that constrain the range of Mg-rich ferromagnesiosilica PCs compositions. The diamonds are the compositions for three random PCs including “Big Guys” (BG; see Table 4). The two lines originating from the Si corner (solid square) represent the maximum lowest possible Mg/(Mg+Fe) ratio, mg = 0.65 in these PCs. The lower line represents the average ratio. The mixing lines connect the metastable eutectic smectite (Sm-d) and serpentine (S-d) dehydroxylate and the Si-bearing Mg-oxide compositions to both low-FeO smectite dehydoxylate compositions. Fe2+,Fe3+ symbolizes variations in the Fe2+/ Fe3+ ratios. Another mixing line connecting the S-d and greenalite (G-d) (open circle) metastable eutectic dust compositions is located in between the stoichiometric olivine (Fo-Fa) and pyroxene (En-Fs) compositional mixing lines. Modified after Rietmeijer et al. (1999a).

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1. The condensed amorphous, metastable eutectic, ‘MgSiO’ dust has three preferred compositions: (I) smectite dehydroxylate, Mg6Si8O22, (II) (rare) serpentine dehydroxylate, Mg3Si2O7, and (III) a low-SiO2 ‘MgSiO’ composition, 2. The condensed amorphous, metastable eutectic, “FeSiO’ dust also has these preferred compositions whereby the ferrosilica smectite dehydroxylate composition is variable as a function of the Fe2+/Fe3+ ratio, 3. The Mg-rich ferromagnesiosilica PCs compositions are defined by mixing of this condensed nanometer-sized dust, and 4. The amorphous metastable eutectic condensates and Mg-rich, Mg/(Mg+Fe) (mg) = 1.0–0.65, ferromagnesiosilica PCs do not have a stoichiometric mineral composition. Massive amorphous Mg-rich ferromagnesiosilica PCs are expected to evolve in Orich circumstellar environments from ferromagnesiosilica agglomerates of amorphous metastable condensates without the necessity of an external energy source. Once formed, they can be processed via hydration or hydrogenation to form stoichiometric layer silicates such as saponite (a smectite mineral), Mg6Si8O22 + 2H2 = Mg6Si8O20(OH)4, or into an assemblage of forsterite, enstatite and silica, Mg6Si8O22 = Mg2SiO4 + 2 Mg2Si2O6 + 3 SiO2, by thermal processing (Rietmeijer et al., 1999a). Initial phase stability for the nmsized grains is dependent on the surface free energy of the reaction products. 9.2.3. Fe-rich ferromagnesiosilica PCs • 9.2.3.1. Processed dust but how? Iron-rich ferromagnesiosilica PCs have an amorphous matrix with embedded nanocrystals and the bulk compositions range from mg = 0.65 to 0.15. The exact matrix composition is cause for confusion. Rietmeijer (1998a) reported nanometer-sized Mg,Fe-olivines and Mg,Fe-pyroxenes with slightly variable mg-ratios, Fe-sulfide with variable Fe/S ratios and Fe-oxides (<30 nm) in an amorphous ferromagnesiosilica matrix. Bradley (1994a) coined the acronym GEMS, viz. glass with embedded metal and sulfides, for similar-sized units with a pure magnesiosilica matrix without embedded silicates in other aggregate IDPs. Actually amorphous GEMS matrix has a ferromagnesiosilica composition with average mg = 0.67 and contains Fe-oxide nanocrystals (Joswiak et al., 1996). There seems to be no fundamental difference between (ultrafine-grained) Fe-rich ferromagnesiosilica PCs and GEMS but they could have different modes of formation. The GEMS formed in interstellar space as the result of irradiation by highly energetic H and He nuclei mixing an enstatite crystal and an Fe,Nisulfide crystal whereby the crystal lattice of both minerals was destroyed which allowed mixing of the resulting amorphous ‘silicate’ and iron materials (Bradley (1994a). Apart from other problems of this GEMS formation (Martin, 1995) the requirement of an enabling external energy source (energetic nuclei) might suggest that GEMS formation is restricted to specific environments. Or, it is limited to GEMS with a pure-Mg matrix wherein metal and sulfide grains are relics of the primary phases along with partial sulfur loss from the mixture. The corollary is that GEMS are either rare or from specific sources. • 9.2.3.2. Post-condensation agglomeration: The most Fe-rich unit of Mg-rich ferromagnesiosilica PCs and the most Mg-rich unit among the Fe-rich ferromagnesiosilica PCs have the same Mg/(Mg+Fe) ratio, viz. mg = 0.65. The Fe-rich ferromagnesiosilica PCs

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could also not have formed by dust condensation alone. They are modified agglomerates of metastable eutectic dusts. Rietmeijer (1999b) proposed a mixing model along experimentally defined mixing lines, i.e. (1) mixing of amorphous Mg-smectite dehydroxylate dust or Mg-rich ferromagnesiosilica PCs with condensed Si-poor, metastable eutectic ferrosilica dust ‘or Fe-oxides (Figure 7). The orientations of the mixing lines and their intersections with the mixing line for the serpentine dehydroxylate condensates allow for the formation of Fe-rich ferromagnesiosilica PCs with a wide range of mg-ratios. However, the intersecting mixing lines for the condensed metastable eutectic dust favor a modal composition with mg 0.65. Compact, chemically homogeneous, amorphous Ferich ferromagnesiosilica PCs formed by post agglomeration fusion of the metastable dust aggregates.

Fig. 7. A slightly modified version of the Mg-Fe-Si (element wt %) diagram in Figure 6 with the mixing lines (dashed) that constrain the compositional range and minimum mg ratio (solid line origination in the Si-corner) for Mg-rich ferromagnesiosilica PCs (diamonds). The average composition for the Fe-rich ferromagnesiosilica PCs in aggregate IDPs (open triangle) is defined by the intersection of the metastable eutectic, serpentine dehydroxylate (S-d) – greenalite (G-d) (open circle) mixing line and the metastable eutectic smectite dehydroxylate (Sm-d) – Si-poor ‘FeSiO’ mixing line. The range of compositions of Fe-rich ferromagnesiosilica PCs (solid squares) (Rietmeijer, 1998a, table 19) is entirely defined by mixing lines between the Mg-rich ferromagnesiosilica PCs and the S-poor ‘FeSiO’ metastable eutectic dust condensates or the apex of the diagram representing condensed Fe-oxides or sulfides after post-condensation sulfidation. GEMS formation required enstatite and Fe sulfide along the mixing line between the apex and enstatite. The preferred metastable composition is represented by the Fe-rich ferromagnesiosilica PC on the minimum mg-line for Mg-rich ferromagnesiosilica PCs. Modified after Rietmeijer (1999b).

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10. Dust evolution 10.1. Matrix aggregates 10.1.1. Carbonaceous matter

Although Principal Components are the smallest entities in the matrix of aggregate IDPs, they are the result of very early dust processing. There is ample evidence for microto millimeter-sized dust around young stars that includes amorphous dust with a pure-Fe or ferromagnesiosilica composition more than a few wt. % FeO and crystalline, pure Mgsilicates (Bouwman et al., 2001; Meeus et al., 2001). The amorphous ‘silicate’ dust might chemically resemble the Mg-rich ferromagnesiosilica PC (Figure 8). It might also resemble the Fe-rich ferromagnesiosilica PCs when they formed by mixing as Rietmeijer (1999b) has suggested. The proportions of PCs (Figure 9) in matrix aggregates with a density as low as 0.1 g/cm3 (Rietmeijer, 1998a) reflect their availability as a function of accretion time and heliocentric distance. It is intuitively plausible, but maybe wrong, that the relative propor-

Fig. 8. The Mg-Fe-Si (element wt %) diagram with the compositions of the condensed metastable eutectic ‘MgSiO’ and ‘FeSiO’ dust (filled circles) and the co-condensed crystalline oxides (solid squares). The compositions of amorphous ‘silicates’ detected around young stellar objects stars are probably constrained within the dashed area by mixing of the condensed dust. Silica denotes that SiO2 in the condensation and thermal annealing experiments, in aggregate and cluster IDPs, and in dust around young stellar objects can be both amorphous and crystalline. Crystalline Enstatite and Forsterite occur around young stellar objects, the collected IDPs and in thermally annealed, condensed magnesiosilica dust (Rietmeijer et al., 2002b). Diamonds indicate the selected Mg-rich ferromagnesiosilica PCs including “Big Guy” (BG). The condensation experiments predict the presence of Mg-rich magnesiosilica dust and MgO. Solids with these compositions have not yet been detected around young stellar objects.

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Fig. 9. Ternary diagrams depicting the accretion hierarchy of matrix and matrix aggregates to aggregate IDPs and to cluster IDPs. The ultimate aggregate particle make-up will plot somewhere within the appropriate triangle depending on the availability of the accreting constituents as a function of time and location of accretion. Modified after Rietmeijer (1998a) and Rietmeijer and Nuth (2000b).

tions of PCs are different in Oort cloud and Kuiper belt comets. Missions to icy protoplanets will be necessary to assess the degree of the original chemical heterogeneity in the outer solar nebula as a function of accretion time and location. For example, carbon in comet Halley is overwhelmingly sequestered in CHON particles with an idealized formula, C100H80N4S2O20 (Kissel and Krueger, 1987) and with lesser amounts in mixed particles. There are two implications for aggregate IDPs, viz. (1) carbon and nitrogen, and most H and some fractions of O and S, occur in carbonaceous and mixed PCs and (2) these element abundances, notably of C and N, are fixed in aggregate IDPs early during

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Table 3. Si-normalized atomic ratios for ultrafine-grained (ufg) and coarse-grained (cg) ferromagnesiosilica PCs, aggregate IDPs, cluster IDP L2008#5 and bulk CI chondrite (sources Rietmeijer, 1998a; (*): Schramm et al., 1989) and proto-CI material (Rietmeijer and Nuth, 2000b). There are no bulk compositions available for the mixed and carbonaceous PCs. Aggregate matrix

C/Si Mg/Si Fe/Si S/Si Al/Si Ca/Si Na/Si Ni/Si Mn/Si Ti/Si K/Si

ufg units

cg units

Zero 0.64 0.46 0.2 0.09 0.027 No data 0.032 0.007 No data No data

zero 0.48 0.18 0.0 0.09 0.04 No data 0.001 0.004 No data No data

Aggregate IDPs

Cluster IDP

Proto-CI

Bulk CI

3.7 (*) 0.85 0.63 0.35 0.063 0.048 0.049 0.037 0.015 0.0022 No data

1.1 0.7 0.7 0.4 0.1 0.07 0.2 0.04 0.01 No data No data

0.52 0.80 0.67 0.35 0.07 0.04 0.05 0.03 0.006 0.002 0.003

0.7 1.06 0.90 0.46 0.085 0.07 0.06 0.05 0.01 0.002 0.01

accretion. Indeed, some amounts of oxygen (Flynn et al., 2001a) and sulfur (Rietmeijer (1998a) is associated with carbon and carbonaceous matter in aggregate IDPs. The Sinormalized carbon content in aggregate IDPs, C/Si = 3.7 (Table 3) with matrix aggregates that could range from C-free, pure ‘silicate’ to pure carbon materials is less than in the dust in the coma of comet Halley, C/Si = 4.4 (Jessberger et al, 1988) (Figure 10). The difference could be a primary feature, viz. the abundance of carbonaceous PCs was lower in the accretion regions of aggregate and cluster IDPs than where comet Halley’s nucleus accreted. Or it could be a secondary feature reflecting parent body modification, dust erosion and space weathering, sublimation during atmospheric entry, or a combination of these processes. 10.1.2. Carbonaceous dust loss

Dust erosion and weathering begins in the coma and continues as space weathering during solar system sojourn. Both processes involve interactions with UV radiation, solar wind and solar flare particles, among others, that will destroy a volatile “organic glue” and other highly volatile materials, e.g. low-temperature Na- and K-minerals, in orbiting debris form icy protoplanets (Rietmeijer, 2001a). In a competing process hydrocarbon polymerization will increase the structural integrity of dust aggregates and friable pebbles. The loss of extremely volatile materials from cometary meteoroids via boiling and sublimation during atmospheric entry was observed in Leonid meteors. Spurny et al. (2000), Campbell et al. (2000) and Murray et al. (1999) detected ‘nebulous meteors’ or “diffuse comet-like structures” at altitudes between 200 km and ~140 km where there are hardly enough molecules to collide with the incoming meteoroids. The nebulous meteor transit to a ‘sharp’ ablation phase with a well-defined meteor head and a sharp trail. Other Leonid meteors showed jet-like structures about 0.5–1.0 km long caused by ejection of

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Fig. 10. Ternary diagram of the three principal components showing the extent of mixing in matrix aggregates as a function of availability at the time and location of accretion of matrix aggregates. The dashed lines reflect my qualitative estimate of the relative proportions inferred from the chemical data available for aggregate IDPs. The position of the open square reflects the modal composition determined in this particle by the author. The dashed line near the apex denotes CHON-like particles as seen in comet Halley. They have not yet been found among the collected stratospheric IDPs, possibly the result of thermal modification during their long history. The filled circles represent the relative proportions in the coma of comet Halley detected by the PIA and PUMA-1 instruments (Langevin et al., 1987) and both PUMA instruments (Fomenkova et al., 1992). The scatter of these points could reflect variations in detected dust compositions as a function of time or it might reflect different instrument detection capabilities and data reduction procedures. Modified after Rietmeijer and Nuth (2000b).

small meteoroid fragments away from the main core of the decelerating meteor (Taylor et al., 2000; Spurny et al., 2000). Both features support the degradation of an extremely volatile “glue” that is holding the structural entities together in the incoming dust from comet Temple-Tuttle some of which survived deceleration at these extreme high altitudes. Dynamic pyrometamorphism when decelerating in the lower, denser atmosphere will further erode surviving less volatile “glue’ or cause hydrocarbon polymerization, carbonization and graphitization. The products of these processes in aggregate IDPs are present in the form of amorphous and poorly graphitized carbons and refractory hydrocarbons with variable (residual?) amounts of H, O, N and S (Rietmeijer, 1998a, 2000a). 10.1.3. Aggregate carbon contents

The average bulk carbon content of aggregate IDPs is ~2 to 3 times its CI abundance but the carbon content in 50 different aggregate IDPs shows a range from ~1 to 47 element wt % with the highest value corresponding to approximately 13 times the CI

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abundance (Thomas et al., 1996). Thomas et al. (1993) suggested a correlation between the bulk carbon content and the olivine to pyroxene ratio but without due consideration to grain size, i.e. relative position of the constituents in the accretion hierarchy, there is no basis for this suggestion. But how should one interpret the carbon data for aggregate and cluster IDPs? The complex history of carbonaceous PCs precludes a simple assessment of a primary or secondary origin of the wide range of carbon contents among the individual aggregate IDPs. The measured average carbon content in the collected IDPs is a lower limit for the most refractory initial carbonaceous matter or the pyrolysis products of more volatile materials that survived the various thermal modification processes. The range of the carbon, and N, H, O and S contents is caused by variable amounts of compositionally different carbonaceous PCs, initially high ratios of carbonaceous to ferromagnesiosilica PS, or both. In this interpretation the surviving carbon and carbonaceous matter might reflect the nature of the original materials in icy protoplanets. Before conclusions of this kind can be considered, future studies of the carbon abundances in aggregate and cluster IDPs should include detailed characterizations of the petrology of carbonaceous and mixed PCs and the relative proportions of matrix PCs. At this time the common interpretation is that aggregate IDPs are ‘carbon enriched’ relative to the CI abundance. It might be germane to consider the possibility that the CI abundance [Table 3] is ‘carbon depleted’ and ask the question if it is possible that the lower CI abundance is a dilution effect from hierarchical dust accretion? 10.1.4. Hydrogen and nitrogen isotopic compositions

The isotopic compositions of hydrogen and nitrogen provide the hard evidence for surviving presolar dust in aggregate and cluster IDPs. In many IDPs the D/H ratios fall within the range among terrestrial rocks but in other aggregate IDPs and most cluster IDPs these ratios are in excess of the terrestrial values and are unarguably extraterrestrial. The D/H ratios show a considerable range among different fragments of the same aggregate particle. This range is larger than found in comets and meteorites and the highest D/H values in aggregate and cluster IDPs, up to δD = 24,800‰, overlap this ratio of molecules in cold interstellar clouds (Messenger, 2000). The highest (or “excess’ relative to terrestrial values) D/H values occur as ‘hot spots’ in organic matter that probably includes polycyclic aromatic hydrocarbons and associated with layer silicates as with water of hydration (Messenger and Walker, 1997; Keller et al., 2000b). The 15N/14N enrichments co-occur with excess D/H values in the same fragments of aggregate and cluster IDPs but they occur in different carrier phases. High 15N/14N ratios in only two IDPs were clearly associated with organic carbonaceous matter (Messenger and Walker, 1997). The highest excess values, up to δ15N = 480‰, imply a formation temperature of ~20K, also consistent with an origin in cold molecular clouds (Messenger, 2000). The excess D/H and 15N/14N ratios found in contiguous patches of organic matter (Bradley et al., 1999, fig. 1A; Keller et al., 2000b; fig. 3) were originally carried in individual carbonaceous PCs that were subsequently fused together. Some of these isotopic compositions are associated with the early accreting carbonaceous PCs in aggregate IDPs that could subsequently become incorporated in the younger cluster IDPs during hierarchical dust accretion. Messenger (2000) found that H and N isotopic anomalies are more common and have a much wider, ‘less equilibrated’ range among the aggregate-IDP fragments of cluster IDPs than among individual aggregate

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IDPs. I submit that the aggregate-IDP in the cluster IDPs sampled a much wider range of “isotopic anomalies” as a function of time and location in the solar nebula, that included both pristine presolar dust and isotopic compositions associated with hydrated dust processed in parent bodies. The latter dust could be mixed and ferromagnesiosilica PCs or amorphous ‘silicates’ that were hydrated in the dirty-ice, pebbles and boulders of icy protoplanets. Linking individual isotopic anomalies to specific host phases will be a roadmap to trace post accretion dust evolution in these rubble piles wherein hydro-cryogenic alteration could not erase the aggregate texture of the accreted dust. 10.2. Evolved dust constituents in aggregate IDPs 10.2.1. Amorphous ‘silicates’

Equilibrium condensation models for the solar nebula predicting stoichiometric minerals, e.g. olivine, pyroxene, plagioclase and feldspar, remains an unproven mode of formation. I submit that the silicate mineralogy could evolve when amorphous ferromagnesiosilica PCs fused into increasingly larger (non-chondritic) amorphous ‘silicate’ grains during hierarchical dust accretion. The mostly <100 up to ~500 nm, Mg,Fe(Ca)silicate and Fe,Ni-sulfide grains that are intimately associated with the matrix aggregates in some aggregate IDPs (Figure 9) could indicate the earliest accretion of non-chondritic dusts. We know very little about these tiny grains but the sulfides include thin (10–30 nm) euhedral plates of low-Ni pentlandite (Ni < 3 at %) up to ~800 nm in size (Tomeoka and Buseck, 1984). The petrological context of these grains in relation to the ferromagnesiosilica PCs leaves the possibility that they might be associated with irradiation processes such as invoked for GEMS formation. Either way the silicate and sulfide grains show the first modification of PCs. The aggregate IDPs contain variable proportions of matrix aggregates, Mg,Fe(±Ca)silicates, amorphous ‘silicate’ grains, Fe,Ni-sulfides and refractory mineral aggregates that are on average ~3 to 5 µm in size. Mg-rich ferromagnesiosilica PCs in the matrix of many aggregate IDPs are an assemblage of a Mg-rich olivine, a Mg-rich pyroxene (both with the same mg-ratio of the bulk PC) plus an amorphous aluminosilica (± Ca) material (Figure 11). This aluminosilica material is a ‘restite’ phase after co-crystallization of Mg,Fe-silicates in an amorphous, Si-saturated material with the smectite dehydroxylate composition of fused condensed dust (Rietmeijer et al., 1999a). Irregular, amorphous patches of fused Mg-rich ferromagnesiosilica PCs form micrometer-sized ‘silicate’ grains in aggregate IDPs such as “Big Guy” (Table 4) (Rietmeijer, 1998a). Other amorphous ‘silicate’ grains are Fe,Mg-bearing, aluminosilica material with variable Al2O3/SiO2 ratios (Figure 4) and almost pure amorphous silica and K-bearing silica (Rietmeijer, 1998a). In the large ‘silicate’ grains the trace and minor elements in PCs become concentrated to detectable levels in increasingly larger ‘silicate’ grains in aggregate IDPs. The 3 µmsized amorphous “Big Guy” ‘silicate’ grain consists of fused Mg-rich ferromagnesiosilica PCs (Rietmeijer (1998a; fig. 45). It is a Si-rich oxygen-deficient non-stoichiometric material [Table 4] that fits in the sequence of “anhydrous biopyriboles”, viz. pyroxene – ‘anhydrous amphibole’ – smectite dehydroxylate, which is unique to many crystalline and amorphous solids in aggregate IDPs (Rietmeijer, 1999a; Joswiak et al., 1999). I use the “Big Guy” composition to show how processing of small constituents

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Fig. 11. Transmission electron micrograph of a coarse-grained Mg-rich ferromagnesiosilica PC in chondritic aggregate IDP L2011A9 showing coexisting olivine (ol), pyroxene (pyr) crystals and an amorphous ‘silica’ ‘restite’ phase contain Al and Ca that was present in the amorphous precursor. Reproduced from Rietmeijer et al. (1999a, fig. 9) by courtesy of the American Astronomical Society.

could produce a variety of larger amorphous units that upon crystallization produces a number of minerals that are larger than the precursors and with increasingly complex mineral assemblages, “Big Guy” ⇒ plagioclase + Mg,Fe-pyroxene + Mg,Fe-olivine + aluminosilicate + silica, viz. 2.5 Mg2.5Fe0.5Ca0.4Al1.4(Si8O21.5)amorphous ⇒ CaAl2Si2O8 + 5 Mg0.8Fe0.2SiO3 + 1.25 Mg1.8Fe0.2SiO4 + 0.75 Al2SiO5 + 11 SiO2 [Eq. 1].

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Table 4. Calculated structural formulae normalized to Si = 8 afu (Rietmeijer, 1999a) for two large, irregularly shaped, amorphous grains: ‘Big Guy’ in aggregate IDP L2011A9 and in IDP U2015B9, a fragment from cluster IDP U2015*B (author, unpublished data.). IDP L2011A9

SiO2 Al2O3 MgO FeO CaO MnO Na2O NiO Oxygen calc.

IDP U2015B9

wt %

Structural formula

wt %

Structural formula

67.1 9.7 14.2 5.3 3.2 0.35 Not detected 0.1 —

8.00 1.40 2.50 0.50 0.40 0.04 — 0.01 21.55

69.3 0.2 26.9 0.2 2.5 Not detected 1.0 Not detected —

8.00 0.11 4.63 0.04 0.32 — 0.42 — 21.37

The minerals produced by this reaction were identified in aggregate and cluster IDPs (Rietmeijer, 1998a) with the exception of aluminosilicate minerals although but amorphous aluminosilica grains abound. The resultant olivine, mg = 0.9 is more Mg-rich than pyroxene, mg = 0.8 which is different when they crystallize within the closed system of a Mg-rich ferromagnesiosilica PC. Thermal annealing studies showed that crystallization of Mg-rich olivine and Mg-rich Ca-free and Ca-rich pyroxenes in these PCs could occur during flash heating up to 1,000 °C (Joswiak and Brownlee, 1998) or at lower temperatures as a function of the prevailing time-temperature regime, dust size and composition (Rietmeijer et al., 2002b). Equation 1 highlights that neoformation of amorphous ‘silicate’ materials via fusion of smaller amorphous ‘silicate’ grains is a prerequisite for silicate evolution in aggregate and cluster IDPs. As a result of this post-condensation dust modification pure silica, either amorphous or crystalline, will be a common phase in these IDPs. The conditions for this feature were preset during kinetically controlled dust condensation in a Mg-Fe-SiO-H2-O2 vapor. Comet Halley’s dust is Si-rich compared to the CI standard (Jessberger et al., 1988) as would be the case for a mixture of Mg-rich and Fe-rich ferromagnesiosilica PCs with the former being slightly more abundant. During the earliest stages of hierarchical dust accretion this metastable dust would have been susceptible to accommodate environmental changes in a series of kinetically controlled dust modification reactions. 10.2.2. Silicate minerals

The Mg,Fe(±Ca)-silicate minerals in aggregate IDPs show considerable clustering at Mg-rich compositions with high abundances of pure enstatite [Mg2Si2O6] and forsterite [Mg2SiO4] for grains ranging from 0.1 to 1.0 µm up to several micrometers (Rietmeijer, 1998a; figs. 36, 41). A conclusion that aqueous alteration of the parent bodies of aggregate IDPs caused Fe-loss from olivines, Ca-free and low-Ca pyroxenes (Zolensky and Barrett, 1994) is premature based on the available data. Olivine [MgFeSiO4] composi-

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tions range from mg = 0.52 to 1.0. The Ca-free [MgFeSi2O6] pyroxenes, including enstatite whiskers (see above), and low-Ca clinopyroxene (<25 mole % Wo; pigeonite) compositions range from mg = 0.46 to 1.0. The micrometer-sized Ca-rich clinopyroxene [Ca(Mg,Fe)Si2O6] grains are mostly Fe-free augite, diopside and Ti-diopside (formerly fassaite).They may contain up to 4 wt % Al2O3 and show evidence for amorphization due to interactions with energetic particles in space (Rietmeijer, 1999a). The mg ratios for the pyroxenes and olivines are mostly within the range, mg = 1.0 – 0.65), for the Mg-rich ferromagnesiosilica PCs. In what could be a compositionally distinct subtype of aggregate IDPs, Ca-poor pyroxenes and olivines (mg >0.95) contain up to 5.1 MnO wt %. Klöck et al. (1989) concluded that this MnO content is consistent with equilibrium condensation in a cooling solar nebula but they did not provide supporting crystallographic evidence, e.g. a whisker shape. Papike (1998) suggested that the Mn/Fe ratio in the solar nebula varied systematically as a function of heliocentric distance. It suggests that the MnO-rich silicates are from a different accretion zone than the silicates in most other aggregate IDPs. This is another example how chemical and mineralogical properties of specific dust types can be used to constrain the processes and environments in the solar nebula as a function of time and heliocentric distance during hierarchical dust accretion. Time is the underlying theme of hierarchical dust accretion that uses grain size as a parameter. It is then unfortunate that we do not know the grain size of each of the analyzed olivine and pyroxene grains. It is available from back-scattered electron images for six out of 22 IDPs. They are aggregate IDPs L2005F39, L2005C37 and L0205Z17 (Zolensky and Barrett (1994; figs. 1, 2) and another aggregate IDP and an unusual massive, polycrystalline IDP without matrix material (Christoffersen and Buseck, 1986a) that contain rare olivine and Ca-free pyroxene with mg < 0.65. The massive IDP is not a proper aggregate IDP because it lacks the PC aggregate matrix. It is a dense monomineralic aggregate that is consistent with expectation of variable proportions of accreting dust. The silicate grain sizes in these IDPs are within the range for embedded silicates in aggregate IDPs. Without the benefit of grain size data there are potential problems as is demonstrated by the IDPs L2005E36 and L2005Z17 with a zoned enstatite grain (20 × 15 µm) and low-Ca pyroxene grains, 15 × 10 µm to 45 × 30 µm (Zolensky and Barrett, 1994). These large grains suggest that these IDPs are fragments from compact, low-porosity cluster IDPs with minimum dimensions of ~25 µm. 10.2.3. Sulfides

It is accepted that Fe,Ni-sulfides in the solar nebula are secondary minerals from sulfidation of metallic iron and/or Fe-oxide by H2S (Zolensky and Thomas 1996). Iron oxides are common among circumstellar dust (Rietmeijer, 1992b). Whether sulfidation was a sustained or episodic phenomenon during solar nebula dust accretion is unknown. Larger Fe,Ni-sulfide grains could evolve from compact agglomerations of condensed Feoxides, or sulfides, or both. The Fe,Ni-sulfide grains are several micrometers, and up to ~ 8 µm in size. They occur as rounded grains, compact clusters of rounded grains and as euhedral pseudo-hexagonal plates (Rietmeijer, 1998a). They are mostly troilite, pyrrhotite with Ni contents ranging from zero to within the range of pentlandite, Ni/Fe (at %) = 0.5–0.7 (Zolensky and Thomas, 1996) and some sulfides have up to 30 wt. % Ni. The crystallographic data support troilite, low-Ni pyrrhotite with a hexagonal superstruc-

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ture and cubic pentlandite (Tomeoka and Buseck, 1984). Rare Fe-rich sphalerite (ZnS) (Rietmeijer, 1998a) and Mn-sulfide (Stephan, 2001) grains are a few micrometer in size. Comet Halley had only a few FeS-dominated grains (Jessberger et al., 1988) among the small grains in its coma. This observation is consistent with the aggregate IDP data wherein most sulfide grains are in the micrometer range that was not sampled by the mass spectrometers at comet Halley. Sun-grazing comet Ikeya-Seki 1965fs contained Co, V, Ni and Cu that would be indicative for probably large Fe-sulfides in this comet (Rietmeijer, 1988b). Both observations, based on very different data, suggest that Fe(Ni)sulfides similar to those in aggregate IDPs occur in icy protoplanets. 10.2.4. Refractory minerals

Refractory oxides occur as aggregate particles, ~5 to ~15 µm, that consist (almost) entirely of mostly small (<500-nm) hibonite [CaAl11Ti0.5Mg0.5O19], gehlenite [Ca2Al2SiO7] and perovskite [CaTiO3] grains. They also contain diopside [CaMgSi2O6], Ti-diopside, corundum [Al2O3], spinel [MgAl2O4], melilite (Na,Mg,Fe-bearing gehlenite) and plagioclase [CaAl2Si2O8] (Rietmeijer, 1998a; Zolensky, 1987). Compositionally similar diopside and Ti-diopside also occur as micrometer-sized grains. The grain sizes of hibonite, gehlenite, perovskite and corundum indicate that some fraction may have been present at the time of matrix aggregate accretion, although they are not yet found to be associated with PCs. It suggests that these grains accreted elsewhere into rare refractory aggregates. They could be either rare presolar dust or residual dust after high-temperature processing in the innermost solar nebula followed by transport into outer solar nebula accretion regions of icy protoplanets. 10.3. Cluster IDPs Aggregate and cluster IDPs are mostly unconsolidated aggregate particles that are not fundamentally different except in the overall size of the accreted constituents. Cluster IDPs are randomly variable mixtures of aggregate IDPs and non-chondritic IDPs (Table 2). The mineralogical and major element compositions of the aggregate-IDP fragments can not be distinguished from those of individual aggregate IDPs. Their >10 µmsized, non-chondritic dust fragments include olivine, Ca-poor and Ca-rich pyroxenes and Fe,Ni-sulfides (Thomas et al., 1995) that are not obviously different from their smaller counterparts in individual aggregate IDPs. For the first time stoichiometric plagioclase (calculated by the author form data in Thomas et al. (1995; table 1) is present among the fragments in cluster IDPs. Smaller, but non-stoichiometric, amorphous and crystalline, plagioclase and feldspar grains occur in aggregate IDPs. We know little about cluster IDPs and their non-chondritic fragments. Flynn et al. (2001b) found that five euhedral-shaped sulfide IDPs (10 to 25 µm in size) had an average Se content of 110 ppm which is higher than in 30 to 50 µm-sized pyrrhotite grains in the Orgueil CI meteorite, viz. Se 65 ppm. Unfortunately we do not yet know the Se content of the individual sulfides in aggregate IDPs. Particle U2071H9 is another massive 10 µm-sized sulfide IDP that has chondritic material at its surface. The sulfide contains significant ppm levels of Co, V and Mn and K (Rost et al., 1999; table 2). These elements, but not K, are typical minor elements in Fe-sulfides. The identical ranges of the bulk carbon contents for aggregate and cluster IDPs, 1 to 44 el. wt % (Thomas et al.,

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1996), are consistent with the fact that this element accreted very early as carbonaceous and mixed PCs in the matrix of aggregate IDPs. The anomalous D/H and 15N/14N ratios (Messenger, 2000) indicate that cluster IDPs incorporated aggregate-IDP fragments from a large and evolved dust population that was depleted by the accretion of matrix aggregates. Compact aggregate and cluster IDPs exist. Textural re-arrangement of aggregate IDPs after ice sublimation or (partial) hydration in small icy protoplanets would promote the formation of compact aggregates. The compact silicate (Christoffersen and Buseck, 1986a) and refractory-mineral rich (Christoffersen and Buseck, 1986b) aggregate IDPs suggest that particle morphology may also be determined by the mineralogical make-up of aggregates. Such compositionally-extreme aggregates are a primary accretion feature. Their incorporation as fragments in cluster IDP might lead to a correlation between a porous or collapsed morphology and cluster IDP bulk composition. In the absence of detailed characterizations of mineralogy, texture and particle mineralogy, bulk compositions for aggregate particles will be of limited value for tracing dust processing in the solar nebula.

11. Chemical evolution 11.1. General considerations The relative proportions the constituents in matrix aggregates and in aggregate and cluster IDPs may be highly variable and form aggregate of a single dust type. The latter aggregates are rare but the apparent scarcity could reflect a sampling bias caused by the window of opportunity for atmospheric entry survival. Accretion of a limited number of non-chondritic dusts will not result in a unique bulk composition for aggregate but large variations are not likely either among aggregate IDPs (Schramm et al. 1989, fig. 12) and cluster IDPs (Table 3). It was appreciated that aggregate IDP bulk compositions are chondritic within a factor of 2 for major rock-forming elements. Compositional spikes, or correlated lower-than-CI abundances, occur for one to three major elements. Deviations from the CI composition were not fully appreciated. They are accretion-related features with information on dynamic processes in the solar nebula. Mapping such anomalies will provide information on the evolution of individual dusts and processes operating in the nebula such as the extent of dust transportation. Refractory minerals that make up the calcium-aluminum-rich inclusions in undifferentiated meteorites (Brearley and Jones, 1998) are a prime example. They are high-temperature condensates or fractionation residues from the innermost hot region of the solar nebula. Finding these minerals with high abundances for the refractory elements Ti, Al and Ca in aggregate IDPs with a matrix of PCs would indicate dust exchange between the innermost solar nebula and its cold outer regions. Rock-forming element abundances in aggregate and cluster IDPs can trace the effects of post-accretion modification processes when the accretion properties are understood. For example, initially high Ca/Si ratios in anhydrous aggregate IDPs are believed to be significantly decreased as a result of parent body aqueous alteration resulting in the formation of hydrated IDPs with lower Ca/Si ratios (Schramm et al., 1989). Zolensky and Barrett (1994) found no evidence for the selective removal of calcium during aqueous

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alteration. They reported identical Ca/Si ratios in anhydrous and hydrated IDPs that were both uniformly below the chondritic Ca/Si ratio. The question that needs to be addressed is whether the low Ca content of particles analyzed by Zolensky and Barrett (1994) is a primary accretion feature that sampled a low abundance of Ca-containing phases. When using Si-normalized element abundances there is a more pressing issue that pertains to the exact silicon abundance in the solar nebula and element dilution during continued accretion of aggregate particles.

11.2. Accretion 11.2.1. Dilution effects

The mass of aggregate and cluster IDPs are primarily caused by the proportions of massive non-chondritic constituents rather than by the low-density matrix aggregates. In this regard the ferromagnesiosilica PCs are not critical to the overall aggregate and cluster IDP bulk composition in the presence of co-accreted Mg,Fe-olivine, Mg,Fe,Ca-pyroxene, amorphous ‘silicate’, plagioclase, alkali-feldspar, silica and sulfide grains and fragments. These grains contribute dramatically to the mass of a particle and its Al, Mg, Si, Ca, Ti, Fe, Ni and Mn abundances relative to the matrix. The matrix aggregates in C-rich aggregate IDPs are dominated by carbonaceous PCs that carry all C and N but the co-accreting with massive non-chondritic dust will dilute the carbon content of the bulk aggregate IDPs. This dilution effect of the bulk composition of aggregate particles becomes more effective with continued hierarchical dust accretion, although it will be most dramatic in the earliest stages. The putatively carbon ‘enrichment’ of aggregate and cluster IDPs is an artifact of early accretion. It will disappear and approach the CI abundance based on 109 to 1012 times more massive CI carbonaceous chondrites. 11.2.2. The silicon abundance

• 11.2.2.1: Predicting the anhydrous mineralogy: Cosmochemists often use Si-normalized element abundances, e.g. Mg/Si(meteorite, IDP) that are then normalized to the CI ratio, viz. [Mg/Si(meteorite, IDP)]/[Mg/SiCI] to compare different samples. Understanding the variations in the silicon content during hierarchical accretion will be critical when comparing normalized element abundances among IDPs and with meteorites. In terms of igneous petrology (Carmichael et al., 1974) a chondritic rock will be classified as ultrabasic (SiO2 <45 wt %) but its average silica content, SiO2 33 wt % (McSween, 1979), is lower than in these terrestrial rocks. A chondritic rock is olivine-normative with not enough silica to react all Mg,Fe-olivine to Mg,Fe-pyroxene whereby the Feo/Fe2+ ratio determines the olivine to pyroxene ratio. The calculated normative mineralogy shows that Fe,Ni-sulfides or metal (~12%) are the only major phases besides silicates (>80%) (Rietmeijer, 1987) which is broadly consistent with the observed relative abundances among collected nonchondritic IDPs (Rietmeijer, 2000a; table 12). The silicon content in aggregate and cluster IDPs is a ‘free parameter that will vary as a function of ferromagnesiosilica dust condensation and accretion history. It is affected by the ratios of (1) carbonaceous to ferromagnesiosilica PCs, (2) matrix aggregates to non-chondritic dust, (3) silicate to sulfide dust and (4) aggregate-IDP fragments to non-chondritic dust in cluster IDPs. The contribution of condensation-related silicon carried in matrix aggregates will diminish during continued hierarchical dust accretion. But will the bulk composition of the

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accreted anhydrous dust in icy protoplanets at the end of hierarchical dust accretion be CI-like? • 11.2.2.2. Proto-CI or CI bulk compositions: Magnesiosilica samples that served as analogs for comet dust formation and evolution (Rietmeijer et al., 2002a, b) were subjected to hydration in the laboratory. This hydration experiment showed that most of the silica that was expelled from the starting material precipitated as compact silica masses. When comparing Si-normalized element abundances for dust from icy protoplanets with normalized CI abundances this result could be significant because the chondritic composition is based on fully hydrated CI meteorites. Based on this experimental result Rietmeijer and Nuth (2000b) substituted the mass of water in CI meteorites with an equal mass of silica to calculate the anhydrous proto-CI composition (Table 3). In the absence of aqueous alteration the Si-normalized proto-CI element ratios should be used for comparing the bulk compositions of aggregate and cluster IDPs. Using the CI values introduces a discontinuity due to secondary processing of accreted dust in CI parent bodies. 11.3. Major and minor element variations Most Al, Ca, Ti and a fraction of Mg is present in Si-free refractory grains and a fraction of Al and Ca also occurs in the silicate minerals gehlenite and plagioclase. Refractory oxide aggregates >10 µm are not found among the fragments of cluster IDPs and individual aggregate IDPs. The existence of larger refractory oxide aggregates may be inferred from the observed co-evaporation of Ca, Ti and Al in meteor trails, suggesting the presence of refractory host minerals in incoming meteoroids (Rietmeijer, 2000a). If it is possible to determine the size of these host minerals from the meteor light curves, such as the composite Leonid meteoroids (Murray et al., 2000), the meteoroid trajectory and velocity will then constrain their source. Iron sulfides and oxides carry substantial amounts of iron, and probably all Ni, while Mg plus some amount of Fe is present in Mg,Fe(±Ca)-silicates. Thus, the bulk Mg/Si and Fe/Si ratios will be sensitive to the sulfide to silicate and the olivine to pyroxene ratios of aggregate and cluster IDPs. The volatile element abundances in collected IDPs are the amounts that survived the thermal interactions they may have experienced during their long history prior to collection in the Earth’s atmosphere. For example, loss of sulfur during flash heating is well documented observationally (Rietmeijer, 1998a) and experimentally (Greshake et al., 1998). Potassium and sodium are highly volatile elements that could form their own lowboiling point minerals but for which there is no mineralogical evidence in the collected aggregate particles. Sodium in comet Halley occurs in hydrated silicate particles (Rietmeijer et al., 1989). The alkali element concentrations in ferromagnesiosilica PCs are below the detection limit of the energy dispersive systems used for the analyses of rockforming elements (Table 3) (assuming they were not lost due to thermal interactions of the incident electron beam with the sample during chemical analysis). Sodium, potassium and calcium were concentrated to detectable (by analytical TEM) amounts in amorphous ‘aluminosilicates’ and non-stoichiometric plagioclase and alkali-feldspar minerals in aggregate IDPs (Table 4). They were further concentrated in the silicate fragments of cluster IDPs that include almost pure silica grains (see, Thomas et al., 1995; table 1). Phosphorous and calcium show a similar concentration history. Thomas et al. (1995) did not provide sufficient petrological data to assess the host phases of the elements in cluster

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fragments. Phosphorous is probably present as calcium phosphate (Stephan, 2001). Despite insufficient petrological data it is still possible to assess the mineralogy of some fragments. For example, olivine and pyroxene were identified in fragment #13 of cluster IDP L2008#5 (Thomas et al., 1995). Other studies showed that olivine and pyroxene are typically Na- and K-free and Al2O3 is only present in diopside, plagioclase and alkalifeldspar (Rietmeijer, 1998a, 1999a). I submit that Na and K in this fragment reside in alkali-feldspar [(Na,K)AlSi3O8]. Similarly, the composition of fragment #39 suggests that it is dominated by a clinopyroxene plus a non-stoichiometric plagioclase [NaAlSi3O8 – CaAl2Si2O8] material. My assessment of the mineralogical make-up of these fragments is speculative. It is consistent with general database on aggregate particles. A high sodium content in comets such as Hale-Bopp with a distinct Na-tail suggests that their dust in part resembles collected cluster IDPs (Rietmeijer, 1999c).

12. Post-accretion Dust Modification Once hierarchical dust accretion had ended, the accreted dust was subject to a variety of post-accretion dust modification processes that are thermal or aqueous in nature. Initial dust modification was governed by (1) the inherited ‘chemical free energy’ of condensed dusts, (2) the ultrafine size of metastable amorphous solids, and (3) the unequilibrated nature of the evolved crystalline solids (minerals). These unique properties are critical and given sufficient activation energy and time the initial condensates will become arranged into the least-energetic physical and chemical configurations. The resulting kinetically controlled reactions may not lead to stable configurations, and continued dust modification in energy-starved environments will proceed earlier and at lower temperatures towards thermodynamic equilibrium. The mineralogical properties of aggregate and cluster IDPs suggest that the solid materials in icy protoplanets did not reach thermodynamic equilibrium.

12.1. Thermal processing 12.1.1. Icy protoplanets

Thermal modification of silicate mineral and amorphous ‘silicate’ dust will neither be efficient or pervasive in icy protoplanets. The black, ‘tar-like’ mantle of processed organic molecules partially covering the surface of the nucleus of comet Halley, which is a feature that is expected to be present on all icy protoplanets, had thermal ‘hot spots’ of up to 125 °C during perihelion (Emerich et al., 1986). Underneath this black mantle at ‘hot spots’ ice-free pockets might exist wherein the conditions are conducive to thermal dust processing, e.g. crystallization of amorphous aluminosilica (± K, Na, Ca) grains. Thermal processing might also occur in the mantle itself. Depending on the thermal conductivity in the mantle and the dust to ice ratio in the underlying dirty-ice, water might have a transient existence and induce hydration reactions leading to layer silicate formation rather than purely thermal processing. The IR spectra of several active comets show the presence of forsterite in the ejected dust and the conventional interpretation is that this silicate mineral predates comet accretion (Nuth et al., 2000a, b). Grain size will

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be an important parameter during dust modification in comets but purely thermal mineral reactions are unlikely to be effective under these conditions. 12.1.2. Solar system sojourn

During solar system sojourn after being released from the parent body dust will interact with energetic nuclei from the solar wind, solar flares and coronal mass ejection events. These interactions are recorded in Mg-rich olivine, Ca-free and Ca-rich pyroxene single-crystals in the form of solar flare tracks caused when the kinetic energy of the impinging particle destroyed the silicate crystal lattice along its deceleration path (Bradley et al., 1984; Rietmeijer, 1998a). Silicate crystals may develop an amorphous rim from the accumulated dissipation of kinetic energy of impinging nuclei (Bradley, 1994a; Flynn, 1996) and small diopside crystals become completely amorphous (Rietmeijer, 1999a). In O-rich environments these interactions may produce a magnetite (Fe3O4) rim on low-Ni, Fe,Ni-sulfide (kamacite) grains (Bradley, 1994b) and resultant release of sulfur. These thermally induced effects might be used to constrain an asteroidal or cometary origin of the host IDP (Sandford, 1987) or identify dust from the Kuiper Belt (Flynn, 1996). Many aggregate IDPs do not show evidence for this particular thermal interaction, because they were shielded inside larger particles or because they were flashheated above ~600 °C which erased the solar flare track record (Sandford, 1987). 12.1.3. Atmospheric entry

Flash heating during atmospheric entry is the final, most dramatic thermal event and yet its dynamic pyrometamorphic effects are surprisingly unappreciated by most workers in the field. Various mineral and chemical modifications are documented in heated IDPs, whereby a narrow (~100 nm) rim of Fe-oxides on Fe-containing constituents in aggregate and cluster IDP is the most readily recognizable feature. In general, this ultra-fast thermal event promotes oxidation reactions but reduction reactions occur in the presence of carbonaceous matter (Rietmeijer, 1998a). Some thermal indicators are less reliable. For example, a claim of zinc loss relies on the assumption that this element had a higher-thanmeasured abundance in the original IDP that contained Zn-bearing phases such as sphalerite and smectite layer silicates (Rietmeijer, 1998a). Localized fusion and (partial) melting of different constituents may obscure original minerals and create new phases with unusual compositions, such as amorphous hedenbergite-like (CaFeSi2O6) domains of melted diopside-sulfide clusters (Rietmeijer, 1998a). Experimental studies (Joswiak and Brownlee, 1998) and modeling studies based on IDP observations (Rietmeijer, 1996) indicate that isochemical crystallization with no substantial diffusion will readily occur in amorphous ferromagnesiosilica PCs. It highlights a need for criteria to determine a primary (pre-entry) or secondary origin for the silicate minerals in Mg-rich and Fe-rich ferromagnesiosilica PCs in aggregate and cluster IDPs.

12.2. Aqueous dust modification in icy protoplanets Lacking sufficient internal energy for pervasive aqueous alteration as in primitive asteroids (Zolensky and McSween, 1988), the environments of icy protoplanets are still able to support a unique form of aqueous dust modification. Hydro-cryogenic alteration is possible due to presence of a few nm-thick film of water at the dust-ice interfaces, which

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is quite efficient to support sub-zero aqueous alteration down to about –50 °C. Hydrocryogenic alteration leads to the formation of layer silicate and salt minerals (Rietmeijer, 1985). It may not cause pervasive aqueous alteration of icy protoplanets. The KOSI (translation: Comet Simulation) experiments showed that transient pockets of liquid water might be possible, albeit barely so, in the dirty-ice of a comet during perihelion (Kömle and Steiner, 1992). The corollary of these experiments is that the conditions for hydro-cryogenic alteration could exist in active comet nuclei and it should be considered as an important dust modification process in active comets. This will be particularly true when hydration could proceed almost isochemically in the case of amorphous dust that already has the appropriate metal-oxide ratios for layer silicates or amphiboles. For example, Fe-smectite forming in Fe-rich “pyroxene glass” (Germani et al., 1990) and non-stoichiometric Ca,K,Na-bearing layer silicates forming in amorphous ‘plagioclase’ [Ca/(Ca+Na) = 1.0–0.0] and ‘alkali-feldspar’ [K/(K+Na) = 1.0–0.06] grains (Rietmeijer, 1998a). The formation of Na,K-bearing, Mg-rich amphibole [arfvedsonite, Na3(Mg,Fe2+)4AlSi8O22(OH)2] (Joswiak et al., 1999) is linked to the presence of amorphous ‘anhydrous biopyribole’ grains and Mg-rich ferromagnesiosilica PCs in aggregate IDPs (Rietmeijer, 1999a). Pervasive aqueous alteration of IDPs requires the presence of water for some (unspecified) period of time at temperatures >0 °C. Ice melting is a costly process that will probably not be effective in energy-poor icy protoplanets. The conditions for aqueous alteration at these temperatures in small rubble pile bodies will be determined by (1) the initial dust to ice ratio, (2) the rate of ice sublimation, (3) the thermal conductivity of dirty ice, the pebbles and boulders, or (4) a combination of these physical properties. Thermal energy might be imparted as dissipated impact energy from collisions among icy protoplanets when ejected into the Oort cloud or evolving in the Kuiper belt. The bottom line is that dust modification, in particular hydro-cryogenic alteration will be possible in thermally active icy protoplanets, which after multiple perihelion passages, might add up to significant dust modification. Ice sublimation ‘drying out’ icy protoplanets might be the more important process that will cause a gradual transition from mostly hydro-cryogenic to purely thermal dust modification. The almost extinct (dormant?) periodic comet Encke might be an example. Its associated annual meteor streams should include mineralogically and chemically-evolved meteoroids.

13. Chemical contamination in the stratosphere The volatile element abundances among individual IDPs may vary for a number of reasons, viz. (1) fixed during hierarchical dust accretion, (2) thermal or aqueous parent body alteration and (3) dynamic pyrometamorphism during atmospheric entry. It is important to establish the origin for other than the already mentioned volatile elements in collected IDPs before we can draw conclusions about the original abundances. The measured abundances of the volatile elements Ga, Ge, Se, Cu, Zn and Br in many IDPs were in excess of their CI values which prompted Flynn et al. (1996) to claim that IDPs are a new type of chondritic extraterrestrial material. This chemical argument confirmed the same conclusion reached previously by Mackinnon and Rietmeijer (1987) who used mostly mineralogical indicators. The observed abundance pattern of increasing element abundances as a function of decreasing condensation temperature (Flynn et al., 1996) is

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consistent with the cosmochemical models predicting increasing solar nebula abundances as a function of increasing heliocentric distance into the regions of icy protoplanets. The measured volatile element abundances are most likely the sum of surviving (pre-entry) volatile elements plus stratospheric contaminant aerosols (Rietmeijer, 1995). An indigenous origin for most of these elements in IDPs is yet to be firmly demonstrated and their host phases are mostly still unknown. Upon reaching terminal velocity, IDPs settle gravitationally (cm/s) through the mesosphere and stratosphere to the collection altitudes. During this time they will encounter condensed and mineral dust aerosols that will cause surface contamination of the IDPs. This process will be most effective in the stratosphere wherein the aerosols are mostly of volcanic and anthropogenic origins, although condensed meteoric aerosols might be present (Rietmeijer, 2002). It is a major problem for assessment of this surface contamination that there is little information on the aerosol compositions and concentrations as a function of time and altitude. Aerosol samplings (Murphy et al., 1998) are rapidly improving this situation. The volatile elements Na, K, Ga, Ge, Se, Cu, Zn, F, Cl and Br seemingly reside at the surface of collected IDPs (Rost et al., 1999; Stephan, 2001). Surface contamination is consistent with the finding of salt minerals of Na, K, Cl and Br (Rietmeijer, 1998a) and the observation of different Br bonding energies for Br on IDPs (Flynn et al., 1996). Sutton et al. (2000) used X-ray fluorescence microtomography, shown on the front cover of this issue, which allows viewing the interior of a collected IDP and clearly casts doubt that volatile elements such as Zn and Br uniquely reside at the surface as aerosol contaminants. It is only possible to trace the origins of these volatile elements in aggregate and cluster IDPs when their host phases are known. For example, rare Sn- and Bi-oxide nanocrystals in aggregate IDPs might be indigenous. Yet, it cannot be ruled out that the volatile Bi-oxides are contaminant aerosol particulates when considering the lower stratospheric volcanic Bi abundances (Rietmeijer, 1998a). Another way to assess volatile element contamination is trying to show that small incoming meteoroids do in fact contain these elements. Thus, Cr, Mn, Cu, Zn, Ga, Ge and Se in micrometeoroids extracted from an aerogel collector exposed in low-Earth orbit on the Mir space station show that these elements are indigenous to some of the incoming dust (Flynn et al., 2001c). I am less pessimistic than Jessberger (1999) about the information that can be extracted from aggregate and cluster IDPs. Minor and trace elements that cannot be tied to specific host phases should be treated cautiously but when tied to a specific host phase they will contain information about the chemical evolution of dust from icy protoplanets. It boils down to the question: What is the information one wants to extract from these particles? And at this point there is still plenty of information that remains to be extracted.

14. Conclusions There are no firm conclusions in the traditional sense of the word. I hope to have shown in which direction research on the chemical and petrological properties of aggregate and cluster IDPs from icy protoplanets should take in the near future. It should include detailed petrological and mineralogical analyses because without them much of the chemical information cannot be assessed appropriately. This research has reached a

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point where astronomical observations and laboratory experiments can be integrated to develop a coherent evolutionary framework. It is a story of hierarchical dust accretion in the solar nebula whereby the energy to support the earliest chemical and mineral evolution was provided during condensation of a small number of chemically-simple, amorphous presolar dust types. These dust compositions can be simulated in kinetically controlled, gas-to-solid condensation experiments and the results predict the dust that is observed around young stars and found in aggregate IDPs. This nanometer-sized, chemically ordered, metastable eutectic dust contains considerable internal “chemical energy” and surface free energy. The extreme non-equilibrium nature of this dust provides the driving force for chemical and mineral dust evolution during hierarchical dust accretion, and all initial thermal and aqueous modification of the evolving dust. The hypothesis of hierarchical dust accretion uses the discrete size distributions for a limited number of non-chondritic dust types as a measure of relative time. It predicts the accretion of gradually larger aggregates with increasingly diverse chemical and mineral properties for crystalline solids that formed from initially amorphous dust. Chemical and mineralogical dust heterogeneity arose during accretion depending on the dust types that were available as a function of time and space during accretion. The story I have developed is essentially my interpretation of observational and experimental data with an occasionally interspersed ‘leap-of-faith” to keep the story moving along. It is not a finished story or a story that can only be told in one particular way. I offer a testable hypothesis for rigorous testing using collected IDPs. Time will tell how well I have told this story connecting the humble beginnings of dust condensation to a complex evolution in icy protoplanets that we can see and visit in our solar system. Aggregate and cluster IDPs, and larger aggregate particles, with unique thermodynamic properties show the transition from a dusty nebula to very young solar system protoplanets wherein these particles were able to experience chemical and mineralogical evolution despite significant internal heating sources. As a result, icy protoplanets in the solar system will be exciting places to visit. Herodotus already said, “If one is sufficiently lavish with time, everything possible happens”. Eventually we will truly understand all the information in extraterrestrial dust entering the atmosphere and collected at their sources.

Acknowledgments I greatly appreciate the encouragement I received from Prof. K. Keil to write this review and his suggestions that greatly improved its presentation. I also thank Prof. G.J. Flynn for his helpful comments. Dr. Steve Sutton of the Advanced Photon Source at the Argonne National Laboratory graciously provided the image shown on the front cover. Jim Karner at UNM assisted with computer graphics. This work was supported by a research grant from the National Aeronautics and Space Administration, NAG5-4441.

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