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The crystal structure of seeligerite, Pb3IO4Cl3, a rare Pb-I-oxychloride from the San Rafael mine, Sierra Gorda, Chile
- L. Bindi, M. D. Welch, P. Bonazzi, G. Pratesi, S. Menchetti
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- Mineralogical Magazine / Volume 72 / Issue 3 / June 2008
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- 05 July 2018, pp. 771-783
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The crystal structure of seeligerite, Pb3IO4Cl3, from the San Rafael mine, Sierra Gorda, Chile, was solved in the space group Cmm2, and refined to R = 3.07%. The unit-cell parameters are: a = 7.971(2), b = 7.976(2), c = 27.341(5) Å, V = 1738.3(6) Å3 and Z = 8. The crystal structure consists of a stacking sequence along [001] of square-net layers of O atoms and square-net layers of Cl atoms with Pb+ and I+ cations located in the voids of the packing. As is typical of cations with a stereoactive lone-pair of electrons, Pb2+ and I5+ adopt strongly-asymmetrical configurations. Pb2+ cations occur in a variety of coordination polyhedra, ranging from anticubes and monocapped anticubes to pyramidal ‘one-sided’ coordinations. I5+ is coordinated by a square of four oxygen atoms: I1 and I3 exhibit a ‘one-sided’ coordination, whereas I2 has square-planar coordination.
The TEM investigation has revealed additional superlattice reflections (which were not registered by X-ray diffraction (XRD)) in the hk0 diffraction pattern of seeligerite based upon a 0.158 Å-1 square net, which can be interpreted as arising from a 20-cation super-sheet motif (12.6 Å x 12.6 Å), likely related to a further level of Pb-I order superimposed upon the 8-site motif identified by XRD.
12 - Howardite–eucrite–diogenite clan
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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- Atlas of Meteorites
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- 11 November 2021
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- 24 March 2013, pp 272-296
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Summary
Introduction
The howardite, eucrite and diogenite (HED) meteorite clan is a suite of igneous rocks from a differentiated asteroid generally thought to be 4 Vesta. According to the Meteoritical Bulletin, as of June 2014, and not accounting for pairing, there were 1246 members of the HED clan, including 284 howardites (16 falls), 797 eucrites (34 falls) and 372 diogenites (11 falls) (www.lpi.usra.edu/meteor/metbull. php). This more than doubles the number listed in the Catalogue of Meteorites [12.1].
The three groups that form the HED clan are:
(i) Howardites, regolith breccias, a physical mixture of fragments of eucrites and diogenites;
(ii) Eucrites, fine-grained basalts with ∼40 vol.% plagioclase and ∼60 vol.% pyroxene (pigeonite); they are subdivided on the basis of chemistry and texture into gabbroic or basaltic, cumulate and non-cumulate;
(iii) Diogenites, orthopyroxenites (usually with <15 vol.% olivine).
Figure 12.1 is a family tree showing the relationship between the different HED groups [12.2, 12.3]. Modal compositions of eucrites and diogenites are given in Table 12.1.
There is mineralogical, chemical and textural diversity among the HEDs [12.3–12.6]. Some of the earliest descriptions subdivided them into monomict or polymict, depending on whether or not they comprise clasts of just one chemical group (monomict) or several groups (polymict) [12.7]. The monomict subgroup includes diogenites (both olivine-bearing and olivine-free) plus all non-polymict eucrites. It contains both brecciated and unbrecciated members. The polymict subgroup is a compositional and textural continuum of regolith and surface breccias that includes components of cumulate eucrites, basaltic eucrites, diogenites and howardites [12.2, 12.8]. Meteorites of the HED clan are described as polymict when they contain more than 90% by volume of a single component (e.g., polymict eucrites contain >90% eucritic material).Howardites are the HED clan meteorites that contain <90% by volume of any single component (see Figure 12.1). They are part of a continuous sequence of polymict breccias from polymict eucrites to polymict diogenites [12.9, 12.10]. Mineralogy is important for distinguishing between polymict eucrites, polymict diogenites and howardites. The boundary of 10% (by volume) clast content that divides the howardites from the polymict subgroups was based on the amount of orthopyroxene detectable by X-ray powder diffraction [12.11].
3 - Ordinary chondrites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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- 24 March 2013, pp 71-167
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Summary
Introduction
The ordinary chondrites (OC) are the largest class of meteorites: from a total of 49,407 records in the Meteoritical Bulletin (update of June 2014), 43,101 (87%) are ordinary chondrites [3.1]; see Table 1.2. Of course, this does not take pairing of desert meteorites into account. Nonetheless, a similar percentage obtains when only records for meteorites observed to fall are considered: of the 1264 observed falls, 870 are OC, some 70% [3.2].
The OC class is subdivided on the basis of chemistry into three groups, the H, L and LL chondrites, for high total iron, low total iron and low total iron plus low metallic iron [3.3–3.7]. On the basis of their mineralogy, a small group of about five OC (including Burnwell) has been recognized as being extremely low in FeO, but with total iron and iron metal even higher than H-group chondrites (‘HH chondrites’) [3.8]. There are also several meteorites (e.g., Tieschitz) that are transitional between groups [3.9]. There are two groups of iron meteorites with silicate inclusions associated with OC: silicates in IIE irons are related to H-group chondrites [3.10, 3.11], whilst silicates in IVA irons have affinities with LL-group meteorites [3.12, 3.13].
As outlined in Chapter 1, the three main ordinary chondrite groups are subclassified into petrologic types that are a measure of the extent of thermal processing experienced by a meteorite. All three groups exhibit the full spectrum of petrologic types from 3 to 6 (Figure 3.1 and see Table 1.2). The most unequilibrated chondrites of petrologic type 3 are known as the UOC (for unequilibrated ordinary chondrites) [3.14, 3.15], and these can be further divided into petrologic subtypes from 3.0 to 3.9 [3.16]. The UOC are the least abundant of all the OC. Even subtracting the effect of large pairing groups of desert OC by considering only meteorites observed to fall, the most abundant meteorites are still H5 and L6 (Figure 3.1) with UOC being ∼3% of all falls [3.2].
The three chemical groups of OC have similar mineralogies, but different mineral chemistries. Although each group is an assortment of olivine, pyroxene and plagioclase, with accompanying Fe-Ni-metal and metal sulphides (Table 3.1), each group also exhibits a specific range of silicate and metal compositions [3.9, 3.17–3.19].
17 - Martian meteorites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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- 24 March 2013, pp 352-371
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Summary
Introduction
This class of meteorites is a group of magmatic rocks whose relatively young crystallization ages implies derivation from a planetary-sized body rather than an asteroid [17.1], leading several authors to consider that the specimens may come from Mars [17.2–17.4]. The suggestion of a potential Martian origin was strengthened by the discovery of argon with a 40Ar/36Ar ratio compatible with that of Mars’ atmosphere (which is very different from the terrestrial atmosphere). The gas was trapped within impact-melt glass in shergottite EETA 79001 [17.5]; subsequent analysis of other noble gases, nitrogen and CO2 [17.6–17.8] confirmed a strong link with the composition of the atmosphere of Mars (Figure 17.1).
For many years, Martian meteorites were referred to as the SNCs (for Shergotty, Nakhla and Chassigny, type specimens of the three original groups within the class). This was before the Martian origin of the meteorites was completely accepted. Now, however, evidence that the meteorites are from Mars is almost indisputable, and the number of groups has increased, so it is prefereable to drop the acronym and simply refer to the samples as Martian meteorites. The Martian meteorites are all igneous rocks and provide critical constraints on the mineralogical, geochemical and geophysical properties of the Martian crust, mantle and core. There have been many reviews of Martian meteorites; the most complete are by McSween [17.1, 17.11], plus the ISSI special issue [17.12, 17.13]. An extremely valuable resource that maintains an updated bibliography of Martian meteorites is the Mars Meteorite Compendium [17.14]. Much of the information in this chapter has been derived from these sources.
Collection of meteorites from Antarctica and hot deserts has dramatically increased the inventory of Martian meteorites. As of June 2014, and not accounting for pairing, there were 132 Martian meteorites listed in the Meteoritical Bulletin [17.15], only five of which are observed falls.
Shergottites (after Shergotty) are the largest group of Martian meteorites, 107 specimens in total at June 2014 (again, not accounting for pairing). Originally, just a single group of basalts [17.1], shergottites are now divided into one of three subgroups: (i) basaltic shergottites (meteorites with a volcanic origin derived from a fractioned magma); (ii) lherzolitic shergottites (meteorites with a plutonic origin and a cumulate texture) [17.16]; (iii) olivine–phyric or picritic shergottites [17.17, 17.18] (meteorites with olivine–porphyritic textures).
Contents
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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8 - Winonaite–IAB–IIICD Clan
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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- 24 March 2013, pp 225-235
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Summary
Introduction
Winonaites are a group of primitive achondrites with chondritic chemistry and reduced mineralogy. According to the Meteoritical Bulletin (www.lpi.usra.edu/meteor/metbull. php), not accounting for pairing, there are currently (June 2014) 26 winonaites, only one of which (Pontlyfni) was observed to fall. The type specimen, Winona, was found in 1928 and originally classified as a mesosiderite [8.1]. The group was also known as the forsterite chondrites [8.2] and IAB chondrites [8.3]. The last name was suggested because of the similarities with silicate inclusions in IAB irons, which are so strong that Mount Morris (Wisconsin) was considered to be part of the silicate fraction of the Pine River IAB iron meteorite [8.4]. It was proposed that winonaites and IAB irons had a common parent body [8.5], and also that the IAB and IIICD iron meteorites formed a single group (the IAB complex) [8.6]. Winonaites are now classified as part of the Win–IAB complex (or Win–IAB–IIICD) clan of primitive achondrites (Figure 1.1) [8.7]. Note, however, that the link between IIICD meteorites and winonaites is weaker than that between IAB meteorites and winonaites [8.5].
Three main criteria are used to classify winonaites: their reduced mineralogy (Fa<7), oxygen isotope composition and the abundance and distribution of metal and troilite [8.8]. Zag (b) is listed as a winonaite in the Meteoritical Bulletin database (and has an oxygen isotopic composition typical of winonaites), but has a mineralogy (olivine Fa19.4; pyroxene Fs25.7) that suggests a relationship with the ungrouped achondrite Divnoe [8.9].
Mineralogy
The modal mineralogy of winonaites is generally similar to that of chondrites, with a mineral chemistry between ordinary and enstatite chondrites [8.5, 8.8]. Figure 8.1 shows the composition of olivine and pyroxene in winonaites (and IAB silicates); modal abundances of minerals are given in Table 8.1. Low-Ca pyroxene is the main mafic silicate phase. Other silicates occur in lesser abundances: olivine, calcic pyroxene (mainly diopside) and plagioclase [8.8, 8.10–8.13].
There is a wide range of plagioclase compositions among winonaites, but only slight variation within an individual meteorite: plagioclase in Pontlyfni is An7–9, in Tierra Blanca it is An12–15 and in Y-75300 it is An25 [8.5, 8.8, 8.10]. Fe-Ni metal (0.1–27 vol.%) and troilite (0.2–28 vol.%) have wide abundance ranges, although this might be the influence of terrestrial weathering [8.5, 8.8]. Accessory phases include chromite, daubréelite, alabandite, schreibersite, graphite, K-feldspar, apatite and whitlockite [8.11, 8.12].
11 - Aubrites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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- 24 March 2013, pp 261-271
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Summary
Introduction
Aubrites (or enstatite achondrites) are coarse-grained breccias with a reduced mineralogy dominated by almost FeO-free enstatite. They have a variable metal content, and their mineralogy and oxygen isotope composition is very similar to enstatite chondrites (EC) [11.1–11.4], although the relationship between aubrites and EC is unclear [11.5]. According to the Meteoritical Bulletin (www.lpi.usra.edu/meteor/metbull.php), as of June 2014, and not accounting for pairing, there were nine observed aubrite falls and 62 finds, including six meteorites classified as anomalous aubrites. The type meteorite, Aubres, fell in France in 1836.
Mineralogy
Aubrites formed under highly reducing conditions and contain a variety of minerals rare in other extraterrestrial rocks (with the exception of EC). The following descriptions are based on the studies of [11.1, 11.2, 11.5, 11.6] and references therein. Enstatite, with very low iron content (FeO <1.0 wt.%), is the most abundant phase (75–98 vol.%); other components present in very variable abundances include forsterite (0.3–10 vol.%), diopside (<3 vol.%), plagioclase (0.3–16 vol.%), kamacite (<4 vol.%) and troilite (<7 vol.%). Figure 11.1 shows the restricted compositional range of pyroxene and olivine in aubrites.
Enstatite grains are millimetres to centimetres in size and are essentially homogeneous in composition and almost FeO-free (En99–100). Diopside (Wo40–46) occurs both as isolated grains and as exsolution lamellae in enstatite and is also almost FeO-free (Figure 11.1). The high-calcium pyroxene is characterized by a very low Cr2O3 content (<0.1 wt.%), lower that in other achondrite groups (Figure 11.2). Minor amounts of FeO-bearing orthopyroxene (En77.8–80.6Wo0.4–1.0) are found within matrix and in clasts [11.8]. Forsteritic olivine (Fa0.2) occurs in the matrix as distinct grains up to 5 mm in size, and also within pyroxene grains. Olivine is also Cr2O3-poor (≪0.1 wt.%) with CaO ∼0.07 wt.% [11.8]. Plagioclase is generally albitic, with a narrow compositional range, An1.8–8.2 [11.2, 11.6]. Plagioclase in impact melt clasts from Norton County has a range of compositions (An0–92.3), implying that the grains crystallized rapidly [11.9]. Fe-Ni metal occurs in a variety of morphologies: as individual cm-sized nodules, as micronsized inclusions within silicates, and as interstitial grains between silicates [11.2, 11.10].
Foreword and acknowledgements
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Why have we produced an atlas of meteorites? All three of the authors trained, at some period of their careers, as geologists. And in doing so, we all used the textbook An Atlas of Rock-Forming Minerals in Thin Section by W. S. MacKenzie and C. Guilford (published by Longman in 1980, and reprinted many times since then). Understanding of the formation of terrestrial rocks and the processes these have experienced is enhanced and facilitated by study of thin sections of material; the same is true for meteorites. Textures, mineralogy and mineral chemistry are all revealed by optical study, enabling classification of meteorites into their different classes and groups. Meteorites differ from terrestrial rocks in containing significant quantities of opaque minerals, especially iron–nickel metal – a phase which dominates the mineralogy of iron meteorites. In such circumstances, thin-section work is neither feasible nor useful, and polished mounts for examination under reflected light are appropriate. So we have attempted to produce an atlas of meteorites – but have omitted the rider ‘in thin section’ – to assist with the recognition, identification and classification of meteorites.
We have not been able to include every meteorite (there are about 60, 000 of them!). And we probably haven't included your favourite meteorite, for which we apologize. But we have tried to produce images of all meteorite-‘type’ specimens, and representatives of sequential petrologic types, textures, shock stages and weathering categories. We have also tried to produce images of each thin section at the same magnification in plane- and cross-polarized transmitted light, and plane-polarized reflected light, to enable different features to be highlighted under different illuminations. This has not always been possible – for many of the older specimens, uncovered thin sections were not available, and only covered sections, thick sections or mounts could be photographed. In some cases, where material was not available for loan, curators sent their own images for inclusion; this was tremendously helpful, and we have acknowledged the photographer as appropriate.
We have drawn heavily on published resources, and, we hope, have referenced them wherever appropriate. The SAO/NASA Astrophysical Data System (ADS) was an essential tool for this project, without which we could not have produced the bibliographies for each chapter.
4 - Enstatite chondrites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
Enstatite (E) chondrites are highly reduced meteorites, and as their group name implies, have a mineralogy dominated by iron-free orthopyroxene (enstatite) as the main silicate phase. Because of their reduced nature, iron occurs within enstatite chondrites mainly as metal and sulphides, rather than silicates, and enstatite chondrites have a high abundance of unusual opaque minerals. Not accounting for pairing, as of June 2014, there were 547 records for enstatite chondrites in the Meteoritical Bulletin, of which 17 are observed falls [4.1]. Around 50% have been classified in terms of subgroups (Table 1.2).
Enstatite chondrites are divided into two subgroups on the basis of their bulk iron content [4.2, 4.3]: the EH group (∼30 wt.% total iron) and the EL group (∼25 wt.% total iron). They are also distinguished by metal and sulphide compositions [4.3]. Like the ordinary chondrites the enstatite chondrite group is subdivided by petrologic type, from 3 to 7, but the distribution of petrologic types is not the same between EH and EL chondrites (Figure 4.1). Although there are several unequilibrated (type 3) enstatite chondrites, it is difficult to classify them into petrologic subtypes because of their thermal and shock histories [4.4, 4.5].
Most of the EH chondrites have been heavily shocked, generally to S4, whereas EL chondrites have a lower mean shock level of S2. Several E chondrites are impact melt breccias [4.6], and based on sulphide composition, EH4, EH5 and EH melt rocks can be divided into high- and lowtemperature classes [4.7]. An intermediate grouplet, currently composed of the meteorites Y-793225 (EH6-an) and QUE 94204 (EH7-an), has been suggested [4.8].
Mineralogy and texture
The relative abundance of different components in enstatite chondrites is given in Table 4.1. The chondrule-to-matrix ratio in each group is very similar, but the average size of chondrules in EH chondrites is ∼200 μm; that of EL chondrites is ∼600 μm [4.9, 4.10]. In both groups chondrules are bigger than metal grains [4.10].
Calcium, aluminium-rich inclusions
Calcium, aluminium-rich inclusions (CAI) are much less abundant in enstatite chondrites than in carbonaceous or ordinary chondrites, comprising <1 vol.% of the meteorites (Table 4.1).
Frontmatter
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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5 - Rumurutiite and Kakangari-type chondrites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
The three major chondrite classes (carbonaceous, ordinary and enstatite) are supplemented by two additional sets of meteorites: the rumurutiites (R chondrites) and Kakangari (K) chondrites. Other than being primitive chondrites, the R and K meteorites have no generic relationship to each other, but to save on paper, we consider them together in this chapter. The type specimen, and only observed fall, of the R chondrites is Rumuruti. It fell in Rumuruti, Kenya in 1934 [5.1], but was not described until 1994 [5.2]. Not accounting for pairing, as of June 2014, there were 152 rumurutiites, most of which have been recovered from northern Africa [5.3]. Rumuruti chondrites are not subdivided into groups, but are probably from a single body. They exhibit a range of petrologic types, from 3 to 6 (Figure 5.1), implying that the parent body has been subject to thermal modification. It has been argued that R chondrites may be part of an OC “super-clan” [5.4], but this interpretation has been rejected in favour of designating the meteorites as a separate class [5.2, 5.5].
Three K-type chondrites are currently (June 2014) listed in the Meteoritical Bulletin [5.3]. The type specimen, Kakangari, fell in Tamil Nadu State, India in 1890 [5.1]. Kakangari has had a varied classification history, first as a unique chondrite, designated a K chondrite [5.6, 5.7], a forsterite chondrite [5.8] then associated with Adelaide and Bench Crater as a CK (for Kakangari) chondrite [5.9]. Since 1989, however, Kakangari has been regarded as belonging to no known chondrite group [5.10], and now forms its own grouplet [5.11] along with Lewis Cliff 87232 (originally classified as a possible CR2 [5.3]) and Lea County 002 [5.12]. This last was originally classified as an anomalous chondrite [5.3], then thought to be a K chondrite [5.11]. There are, however, sufficient differences between Lea County 002 and Kakangari/LEW 87232 for its classification to be questioned, and it might be a weathered CR2 chondrite [5.13]. Because of this doubt, Lea County 002 will be excluded from further consideration. Kakangari and LEW 87232 are both petrologic type 3.
Much of the information in this chapter about R chondrites is drawn from the review paper by Bischoff et al. (2011) [5.14]; similarly, information about K chondrites comes mainly from Weisberg et al. (1996) [5.11].
1 - Introduction
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Solar System history started some 4567 million years ago with the collapse of an interstellar molecular cloud to a protoplanetary disk (the solar nebula) surrounding a central star (the Sun). Evolution of the Solar System continued through a complex process of accretion, coagulation, agglomeration, melting, differentiation and solidification, followed by bombardment, collision, break-up, brecciation and re-formation, then to varying extents by heating, metamorphism, aqueous alteration and impact shock. One of the key goals of planetary science is to understand the primary materials from which the Solar System formed, and how they have been modified as the Solar System evolved. The last two decades have seen a greater understanding of the processes that led to the formation of the Sun and Solar System. Advances have resulted from astronomical observations of star-formation regions in molecular clouds, the recognition and observation of protoplanetary disks and planetary systems around other stars, and also from advances in laboratory instrumentation that have led to more precise measurements on specific components within meteorites, e.g., refinement of chronologies based on shortlived radionuclides. Results from meteorites are important because meteorites are the only physical materials available on Earth that give direct access to the dust from which the Solar System formed. Study of meteorites allows a more complete understanding of the processes experienced by the material that resulted in the Earth of today.
Naming of meteorites
Meteorites are pieces of rock and metal, almost all of which are fragments broken from asteroids during collisions. They fall at random over the Earth's surface, and have also been identified as components within lunar soils [1.1, 1.2] and on Mars’ surface [1.3]. Meteorites are named from their place of find or fall, traditionally after a local geographic feature or centre of population. However, where large numbers of meteorites are found within a limited area, this convention is not possible to follow. The recovery of meteorites from desert regions has resulted in a name–number nomenclature that combines geographic and date information. Antarctic specimens collected by government-funded expeditions are given a year–number combination with a prefix recording the icefield from which they were retrieved (e.g., Allan Hills 84001), whereas meteorites collected in hot deserts are simply numbered incrementally by region (e.g., Dar al Gani 262).
2 - Carbonaceous chondrites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
The class of meteorites known as carbonaceous chondrites was originally an assemblage of three groups that were associated on the basis of relatively high abundances of carbon and water [2.1]. Now, however, the class comprises eight individual groups, most of which are not noticeably enriched in either carbon or water [2.2]. They are considered collectively as a class on the basis of their mostly unfractionated bulk chemical composition relative to that of the Sun, and the individual groups are now thought not to be genetically related or to come from the same parent object. CI chondrites are the most primitive of the eight groups, and have a composition identical to that of the solar photosphere for all but the most volatile of elements (Figure 2.1).
Each carbonaceous chondrite group is assumed to originate from a separate parent body; associations of groups (clans, or supergroups) are thought to be of parent bodies that formed in similar regions of the solar nebula at similar times [2.2]. Figure 2.2 is a family tree for carbonaceous chondrites, showing where individual groups might cluster together in clans [2.2]. Association into clans tends to be based on commonalities in chemistry of the clan members, rather than mineralogy or texture. The numbers and names of specimens in each group are recorded and regularly updated at the Meteoritical Bulletin website [2.5].
Carbonaceous chondrites are composed of several distinct components [2.6–2.8]: fine-grained darkmatrix, chondrules, high-temperature inclusions and opaque minerals. There is a variation in chemistry and relative abundance of these components among the different carbonaceous chondrite groups.
Petrologic types and subtypes are applied to carbonaceous chondrites (Figure 2.2 and Table 2.1). CI, CM and CR chondrites include only petrologic types 1 and 2; their finegrained matrices are phyllosilicate-rich, indicating that the meteorites have been altered by aqueous fluids. McSween [2.7] suggested that the trend from type 2 to type 1 was one of increasing aqueous alteration. The other carbonaceous chondrite groups show little evidence for extensive aqueous alteration, and the major matrix minerals are olivines and sulphides [2.6]. It is mostly the CK chondrites which have specimens with petrologic type 4 and above.
There have been several thorough reviews of the mineralogy and petrology of carbonaceous chondrites [e.g., 2.6, 2.9–2.11] from which much of the information in this chapter is drawn.
7 - Brachinites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
Brachinites are achondrites that formed early in Solar System history. They consist almost exclusively of olivine; if they were terrestrial rocks, they would be classified, petrographically, as dunitic wehrlites [7.1]. It is not completely clear whether brachinites are primitive achondrites or more differentiated cumulate meteorites [7.2]. According to the Meteoritical Bulletin, not accounting for pairing, there are currently (June 2014) 37 brachinites, none of which was observed to fall (www.lpi.usra.edu/meteor/metbull.php). The type specimen, Brachina, was found in South Australia in 1974, and originally thought to be a second chassignite [7.3]. It was recognized as a unique meteorite almost a decade after it was found [7.4], and the brachinite group was proposed following the recovery of several additional meteorites [7.5–7.8]. The olivine-rich achondrite LEW 88763 is recorded as a brachinite in the Meteoritical Bulletin, but has characteristics that suggest it might be a unique achondrite [7.9]. The Tafassasset meteorite has also been grouped with the brachinites, but it too may be unique [7.10].
Mineralogy
The most complete descriptions of brachinites are given by [7.1, 7.11], from which much of the following information is derived.
The main mineralogical component of brachinites is olivine, with variable amounts of high-Ca pyroxene (augite or diopside), iron sulphides, chromite (0.5–2 vol.%), traces of orthopyroxene, minor phosphates (mainly chlorapatite and whitlockite) and Fe-Ni metal (Table 7.1). The composition of olivine and pyroxene is shown in Figure 7.1. Olivine compositions range from Fa29 to Fa36, with molar Fe/Mn ratio of 52–77 and CaO content generally around 0.10 wt.% [7.7, 7.11]. Clinopyroxene composition is Wo38–45 with mg# of 79–84; orthopyroxene, when present, has Wo2–4 and mg# of 71–74 [7.1, 7.9–7.12].
Chrome-rich spinels are Ti-poor in comparison with those of equilibrated ordinary chondrites [7.1, 7.11] and have molar Cr/(CrþAl) ratios ranging from 0.73 to 0.83 [7.11]. Sulphides are mainly troilite with 0.4–3% Ni, whilst the rare metal grains are taenite containing 19–55 wt.% Ni and 1–2 wt.% Co [7.1, 7.5, 7.11]. Brachina and EET 9940n both contain up to 10 vol.% plagioclase (An14–40), whilst the other brachinites contain almost none. The same two meteorites have up to 0.26 wt.% CaO in olivine [7.5, 7.12, 7.14]. A plot of CaO vs. Cr2O3 in olivine distinguishes brachinites from other primitive achondrites, as well as from HEDs (Figure 7.2).
Institutions that provided specimens
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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15 - Iron meteorites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
The most comprehensive systematic study of iron meteorites was published in 1975, and still remains the most authoritative and best illustrated text on the subject [15.1], despite advances in classification and understanding of chemical composition [15.2]. Only a few iron meteorites with representative textures are included here. Although [15.1] is out of print, it is available on-line (evols.library.manoa.hawaii. edu/handle/10524/33750). The number of iron meteorites listed in the Meteoritical Bulletin (June 2014), and not accounting for pairing, was 1097 (www.lpi.usra.edu/meteor/metbull.php). Only 49 of these (∼5%) are observed falls. Iron meteorites are now classified into 11 groups on the basis of composition [15.2], with a large number of specimens (∼20%) that are not affiliated to any specific group. An additional 292 meteorites were classified as iron meteorites but are now regarded as related to either primitive achondrites or ordinary chondrites (see Table 15.1).
Iron meteorites consist of metallic iron with variable nickel contents (typically between 5 and 15 wt.%), although there are a few meteorites with up to 60 wt.% Ni [15.1]. Iron meteorites are the biggest and heaviest examples of extraterrestrial material that fall to Earth. Because they are so different in appearance from terrestrial rocks, they are readily identified and collected. Prior to the start of regular collection trips to deserts (which have returned several thousand stone meteorite finds), iron meteorites dominated meteorite finds, although accounting for only about 5% of witnessed falls [15.3]. A single iron meteorite can weigh several tonnes. As noted above, one of the most important sources of information about iron meteorites is The Handbook of Iron Meteorites by Buchwald [15.1], from which much of the information given here about the structural classification of iron meteorites is drawn. Recent reviews of iron meteorites [15.4–15.7] have also been valuable sources of information.
Classification of iron meteorites
Structural classification
Iron meteorites are classified into eight subgroups on the basis of metallic structure, an intergrowth of two metals of Fe with Ni. The structure is shown when an iron is polished, and then etched with nital, a solution of nitric acid in alcohol. One metal (kamacite; α-Fe-Ni, also known as ferrite) is body-centred cubic iron with <6 wt.% nickel. The second is taenite (γ-Fe-Ni, or austenite), face-centred cubic iron with >25 wt.% nickel.
13 - Mesosiderites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
Historically, mesosiderites were classified as one of the two major divisions of stony-iron meteorites, the other being the pallasites (Chapter 14). Meteorites within the two divisions have very different compositions, textures and petrogeneses, and it is now thought to be more appropriate to consider mesosiderites and pallasites (as well as iron meteorites) as metal-rich achondrites [13.1]. As of June 2014, and not accounting for pairing, there are currently 202 mesosiderites, only seven of which are observed falls [see Meteoritical Bulletin Database, www.lpi.usra.edu/meteor/metbull.php].
There are several reviews of mesosiderites, from which much of the following information has been drawn [13.2– 13.5]. Mesosiderites are polymict breccias composed of metal and a variety of silicate clasts (up to several cm across) set in a matrix of silicate plus metal. The silicate clasts are both lithic and mineral; the former (up to ∼10 cm across) are igneous in nature, and include basalts, gabbros and pyroxenites. Dunites are less abundant and anorthosites are rare. Mineral clasts are almost all monomineralic, either coarsegrained low-calcium pyroxene or olivine; plagioclase clasts are less abundant [13.6, 13.7]. Metal is unevenly distributed as cm-sized slugs, veins and smaller (submillimetre to millimetre) grains mixed with the silicate clasts.
There have been several schemes for the classification of mesosiderites, the most successful of which were based on the texture of silicate matrix [13.8, 13.9]. Texture, however, is a secondary classification parameter, related to parentbody processing rather than formation environment. The classification scheme now generally adopted is based on silicate mineralogy [13.10]. Mesosiderites are subdivided into three classes (A, B and C) according to plagioclase content and the ratio of orthopyroxene to plagioclase (Figure 13.1). Silicates in class A mesosiderites have a large eucritic component (higher plagioclase and clinopyroxene abundances); those in class B are more diogenitic (containing a greater amount of orthopyroxene), whilst silicates in the only member of class C (RKP A79015) are composed almost entirely of orthopyroxene [13.11]. The three classes are also recognized on the basis of the siderophile composition of the metal component (Figure 13.4 below) [13.12–13.14]. Note, though, that a recent analysis of the metal in RKP A79015 does not distinguish it from class A mesosiderites [13.15]. Modal mineralogies are sufficiently distinct that mineralogical class is readily established; the same parameter allows mesosiderites to be distinguished from howardites and eucrites.
9 - Ureilites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
Ureilites are coarse-grained, ultramafic meteorites. They have igneous textures and are depleted in incompatible lithophile elements, indicating that they have been heated, but they also exhibit primitive, nebula-derived signatures, including variable oxygen isotopic compositions. Along with acapulcoites, lodranites, brachinites and winonaites, ureilites are primitive achondrites and form a bridge between undifferentiated chondrites and fully differentiated meteorites. According to the Meteoritical Bulletin (June 2014), and not accounting for pairing, there are 368 ureilites [9.1], making them the second-largest group of stony achondrites. Only six of the ureilites are falls; the type specimen after which the group is named is Novo Urei. The most recent fall (Almahata Sitta, which fell in the Sudan in October 2008) was observed as an incoming asteroid by several telescopes as it approached the Earth, allowing its trajectory to be calculated and its fall location identified [9.2].
Mineralogy
The information in this section is drawn from review papers by Mittlefehldt, Goodrich and colleagues [9.3–9.5], plus references therein. Many ureilites are an assemblage of ∼60 vol.% olivine and ∼30 vol.% pyroxene, with minor (≤10 vol.%) matrix or interstitial material consisting of carbon (graphite and diamond), iron-nickel metal, sulphides (troilite with up to 34.0 wt.% chromium [9.6]), and finegrained silicates (Table 9.1). Olivine (mg# 76–95) and pyroxene (mainly pigeonite, but some ureilites contain orthopyroxene or augite) core compositions vary widely between ureilites but are homogeneous within each ureilite. The core compositions of olivine and pyroxene in unbrecciated ureilites are shown in Figure 9.1. Most ureilites (∼90%) are unbrecciated; the remainder are polymict breccias. (Note, we follow the suggestion of Downes et al. [9.7] and drop the term ‘monomict’ for unbrecciated meteorites.)
Olivine grains in unbrecciated ureilites are euhedral to anhedral, have inclusion-free cores (Fa5–24) and are generally rich in Cr2O3 and CaO [9.3–9.5]. Their high Cr2O3 and CaO contents distinguish ureilite olivines from olivine in other achondrites (Figure 9.2). Many olivine grains have an almost FeO-free rim (∼10–100 μm) containing small inclusions of Fe-Ni metal [9.8, 9.9]. Pigeonite grains lack exsolution lamellae, suggesting rapid cooling from high temperatures [9.10, 9.11]. Pigeonite contains fine-grained metallic inclusions (0.46–2.60 vol.%); the absence of reduction rims around these inclusions implies co-genesis with the pigeonite [9.12].
6 - Acapulcoites and lodranites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
Acapulcoites and lodranites are igneous rocks recognized as a family (or clan) of primitive achondrites that have common mineralogical, petrographic and geochemical features. In the last decade, there has been a substantial growth in the number of family members: according to the Meteoritical Bulletin, not accounting for pairing, there are currently (June 2014) 58 acapulcoites, nine meteorites transitional between the two groups, 48 lodranites and one anomalous lodranite (Y 74357) (www.lpi.usra.edu/meteor/metbull.php). There is one acapulcoite fall (Acapulco) and one lodranite fall (Lodran); these are the type specimens of the two groups within the acapulcoite–lodranite clan. Although the bulk composition of members of the acapulcoite–lodranite clan resembles that of ordinary chondrites, textural and mineralogical evidence suggest a gradation in the degree of heating of their parent body, from acapulcoites that contain relict chondrules to lodranites that have lost sulphide and feldspathic (plagioclase-bearing) melt [6.1, 6.2].
Some of the criteria originally proposed for discriminating between acapulcoites and lodranites were too limited to encompass meteorites transitional between the endmembers [6.3], leading to the proposal of other classification schemes [6.4, 6.5]. On the basis of mineral chemistry, the meteorites can be divided into one group of acapulcoites and two groups of lodranites [6.4]: magnesian lodranites (Gibson, Y 75274 and Y 8002) and ferroan lodranites (all other lodranites). An alternate classification [6.5], based on K and Se abundances (Figure 6.1), recognized subtypes related to thermal history: primitive acapulcoites with near chondritic elemental abundances; typical acapulcoites characterized by incipient melting and loss of sulphide phases; transitional acapulcoites marked by a deficit of sulphides and melting or loss of plagioclase; lodranites with subchondritic elemental abundances and depletions in troilite, metal and plagioclase; and enriched acapulcoites with a clear addition of feldspar probably resulting from trapping of partial melt lost from lodranite regions [6.5, 6.6].
Mineralogy
Much of the information in the following sections is drawn from the review by Mittlefehldt et al. [6.7], and references therein. Mineralogical characteristics of the type meteorites Acapulco and Lodran can be employed to distinguish between acapulcoites and lodranites.
16 - Lunar meteorites
- Monica M. Grady, Giovanni Pratesi, Vanni Moggi Cecchi
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Summary
Introduction
It is well nigh impossible to compile a chapter about lunar meteorites without frequent reference to information obtained from the materials returned directly to Earth from the Apollo and Luna missions. We will, however, attempt to confine ourselves as much as practicable to observations arising directly from studies of lunar meteorites. (Note that the term ‘lunaite’ is not accepted (or acceptable) for lunar meteorites.) There is a wealth of data from lunar studies, much of which has been collected in review articles, either of Apollo and Luna data [16.1–16.3] or specific to lunar meteorites [16.4, 16.5]. Other valuable resources for material about lunar meteorites can be found at specialist websites [16.6a, b]. These sources have been drawn on heavily to produce this chapter.
No lunar meteorites have, so far, been observed to fall, and all lunar meteorites have (as of June 2014) been collected from desert regions, mainly Antarctica and Northern Africa [16.7]. The first meteorite to be recognized as lunar was collected in Antarctica in 1981 [16.8]. Since then, not taking account of pairing, there are currently 181 meteorites from the Earth's Moon [16.4, 16.6a, 16.7]. Lunar meteorites account for fewer than 0.4% of all known meteorites. Kalahari 009, at 13.5 kg, is the largest so far found [16.8].
Lunar meteorites are ejected from the Moon's surface by excavation during impact. They are a more random sampling of the Moon than the Apollo and Luna specimens, which sampled material from an area that covered <5% of the Moon's surface (Figure 16.1) [16.2, 16.3]. Spectral reflectance measurements taken by the Clementine orbiter (1994) and data from gamma-ray and neutron spectrometers on the Lunar Prospector mission (1998–99) revealed that Apollo sites were not completely representative of the Moon, since most of them were collected from a radioactive ‘hot spot’ close to Mare Imbrium [16.9, 16.10]. This region (characterized by intermediate iron concentration (Figure 16.1) and high concentrations of K, Th and U) is also known as the Procellarum KREEP Terrane, or PKT.
Over the past 25 years, since lunar meteorites were first recognized [16.8], they have yielded important data that are complementary to information derived from Apollo and Luna samples, and have helped to improve our knowledge of the Moon [16.5].