Mineralogy of V-type near Earth objects

FCUP Mineralogy of V-type Near Earth Objects 1 I dedicate this thesis to all my family and to my girlfriend. FCUP Mineralogy of V-type Near Earth Objects 2 Special Thanks I firstly would like to thank my parents, grandparents, uncles and my girlfriend to all the support that they gave to me and for their patience and help. I would like to thank my supervisor Prof. Dra. Teresa Seixas, for giving me the opportunity to study and learn more about the research field of asteroids and meteorites. Without her this work would not be possible. Thankfully to Prof. Teresa Seixas, the work with meteorites started in 2011, when I have done my first extracurricular internship PEEC (Plano de Estágios Extracurriculares da FCUP), entitled “Caracterização Laboratorial de Meteoritos”, in its third edition. During my first extracurricular internship PEEC, it was possible to present two posters communications: the first was presented in the IJUP11 (Investigação Júnior da Universidade do Porto 2011) conference and was entitled “Laboratorial Characterization of Meteorites”; the second was presented in the XXI ENAA (Encontro Nacional de Astronomia e Astrofísica) conference and was entitled “Laboratorial Characterization of a Campo del Cielo meteorite”. During the work of this thesys, it was possible to present a poster communication, entitled “The early stages of the Solar System: constraints from the small icy worlds”, in the IJUP12 (Investigação Júnior da Universidade do Porto 2012) conference. In addition, it was also possible to make an oral presententation, entitled “Mineralogy of V-type near-earth objects”, in the Física2012 conference. I also would like to thank my supervisor for all the constructive comments and suggestions for this thesis. I also would like to thank all researchers who collaborated in the RELAB and The MIT-UH-IRTF Joint Campaign for NEO Spectral Reconnaissance. All (or part) of the data utilized in this publication were obtained and made available by the The MITUH-IRTF Joint Campaign for NEO Reconnaissance. The IRTF is operated by the University of Hawaii under Cooperative Agreement no. NCC 5-538 with the National Aeronautics and Space Administration, Office of Space Science, Planetary Astronomy Program. The MIT component of this work is supported by NASA grant 09-NEOO0090001, and previously by the National Science Foundation under Grant No. 0506716. FCUP Mineralogy of V-type Near Earth Objects 3 Abstract Much of our current knowledge on the formation and evolution of our Solar System is based on the study of chemical, mineralogical, and physical properties of meteorites. These are remnants of planetesimals that never formed a planet, possibly due to the gravitational influence of Jupiter, and preserve crucial information about the collisional and evolutionary processes that occurred in the early stages of the Solar System. Linking a meteorite to a specific asteroid or asteroid’s group can provide important information about petrological, thermal and collisional history of asteroids. In addition, if a linkage between a meteorite and an asteroid can be established, knowing the orbital position of the asteroid allows us to argue about the mechanism that can deliver a meteorite to Earth. Visible and near-infrared reflectance spectroscopy of meteorites, asteroids and samples of mixtures of meteorite constituent minerals provides a powerful technique to study a possible genetic linkage between meteorites and asteroids. However, reflectance spectra are not always intrinsic to surface mineralogy, as grain size effects, or surface roughness, can change the optical properties of any mineral. In the case of airless bodies like asteroids, the effects of space weathering caused by irradiation of solar wind and cosmic rays and by the bombardment of micrometeorites can redden the surface reflectance spectra and lower its albedo. In addition, the absorption features of the asteroid surface reflectance spectra can be attenuated. Hence, studying the mineralogy of asteroids is a quite arduous task. However, the V-type asteroids and the Howardites (H), Eucrites (E) and Diogenites (D) seem to share the best known meteorite-asteroid linkage. There are also possible genetic linkages that have been postulated, such as ordinary chondrites and S-type asteroids or, even more surprisingly, the recently argued possibility of the CM and CI chondrites coming from comets. In this study, it is aimed to analyze the reflectance spectra of V-type Near Earth Objects (NEOs) and 4 Vesta taken from the publically available MIT-UH-IRTF Joint Campaign for NEO Reconaissance, and reflectance spectra of HED meteorites and intimate mineral mixtures from the RELAB database. In order to fit reflectance spectra and derive the surface composition, mineralogy and grain size, the Hapke radiative FCUP Mineralogy of V-type Near Earth Objects 4 transfer model is applied. The results of the spectral analysis and modelling corroborate the known genetic linkage between V-type asteroids and the studied HEDs. A similar linkage between several NEOs and V-type asteroids is proposed. It is also hoped to constrain some processes that took place in the early Solar System. FCUP Mineralogy of V-type Near Earth Objects 5 Resumo O nosso conhecimento sobre a formação e evolução do Sistema Solar advém do estudo das propriedades químicas, mineralógicas e físicas dos meteoritos. Estes são objetos remanescentes de planetesimais que nunca chegaram a formar um planeta possivelmente devido à ação gravítica de Júpiter e preservam informação crucial sobre os processos evolutivos e colisionais que ocorreram durante os primórdios do Sistema Solar. Estabelecendo uma ligação genética entre um meteorito e um dado asteroide ou grupo de asteroides, pode providenciar informação importante sobre a história petrológica, térmica e colisional dos asteroides. Igualmente, se tal conexão poder ser estabelecida e se a posição orbital do asteroide for conhecida, é possível argumentar-se sobre quais os mecanismos responsáveis em ejetar os meteoritos para a órbita da Terra. A espectroscopia no visível e no infravermelho de meteoritos, asteroides e amostras de misturas de minerais constituintes dos meteoritos, é uma excelente técnica para estudar as ligações genéticas entre meteoritos e asteroides. Mas, os espectros de reflexão nem sempre são intrínsecos à superfície, pois o efeito do tamanho de grão e a rugosidade da superfície podem modificar as propriedades óticas de qualquer mineral. No caso dos asteroides, os efeitos da erosão espacial causados pela irradiação devida a raios cósmicos ou ventos solares ou até mesmo causada pelo bombardeamento de micrometeoritos pode envermelhecer o espectro de reflexão, atenuar as suas bandas de absorção e diminuir o albedo. De fato, tentar encontrar possíveis ligações genéticas entre meteoritos e asteroides pode ser uma tarefa bastante árdua. No entanto, os asteroides do tipo V e os meteoritos HED (Howardites, Eucrites e Diogenites) aparentam partilhar uma forte ligação genética, dadas as muitas semelhanças entre os seus espectros de reflexão. Existem também outras possíveis ligações genéticas que têm sido postuladas, como por exemplo a ligação entre os condritos ordinários e os asteroides do tipo S, ou ainda mais surpreendentemente a ligação entre os condritos carbonáceos CM e CI e os cometas. Neste estudo, pretende-se analisar espectros de reflexão de Near Earth Objects (NEOs) do tipo V e do asteroide 4 Vesta, obtidos a partir da base de dados pública MIT-UH-IRTF Joint Campaign for NEO Reconaissance, e espectros de FCUP Mineralogy of V-type Near Earth Objects 6 meteoritos HED e misturas de minerais obtidos a partir da base de dados pública RELAB. É espectável conseguir-se estabelecer uma possível ligação genética entre alguns NEOs e asteroides do tipo V e os meteoritos HED. É também espectável imporse algumas restrições em alguns processos de formação do Sistema Solar. FCUP Mineralogy of V-type Near Earth Objects 7 Indices Page: Indices of Figures…………………………………………………………………………..…12 Indices of Tables……………………………………………………………………………...25 Indices of Abbreviations……………………………………………………………………...27 Introduction…………………………………………………………………………………….29 1. Formation of the Solar System……………………………………………………..…….31 1.1. Historical Overview………………………………………………………………………31 1.1.1. Turbulent Theory………………………………………………………………………32 1.1.2. Nebular Theory……………………………………………………………...…………32 1.1.3. Tidal theory………………………………………………………...…………………..34 1.1.4. Accretion theory……………………………………………………………………….35 1.2. Formation and evolution of the Solar System: Planet Formation…………………..35 1.2.1. The protoplanetary nebula and the first solids……………………………………..36 1.2.2. The origin of planetesimals…………………………………………………………..38 1.2.3. The formation of terrestrial planets………………………………………...………..39 1.2.3.1. Origin of the Moon…………………………………………………………………..41 1.2.4. The formation of the Giant planets…………………….…………………………….42 1.2.4.1. Core-Accretion theory………………………………………………………………43 1.2.4.2. Disk Instability theory……………………………………………………………….43 1.2.4.3. Observational Constraints: theory against observations………………………..43 1.2.4.4. Formation of Uranus and Neptune…………………………………………….….45 1.2.5. The “Tidal Downsizing” hypothesis, a more plausible theory?............................45 FCUP Mineralogy of V-type Near Earth Objects 8 Page: 1.2.6. The Asteroid Belt…………………………………….………………………………..47 2. Small Bodies in the Solar System………………………………………………………..49 2.1. Asteroids: Shattered Worlds……………………………………………………………50 2.1.1. Classification of Asteroids…………………………………………………………….51 2.1.2. Orbits of Asteroids………………………………..……………………………...……53 2.1.3. 4 Vesta………………………………………………………………………………….55 2.1.4. Families of Asteroids………………………………………………………………….57 2.2. Comets: “Dirty snowballs”.……………………………….……………………………..60 2.2.1. Taxonomy of comets………………………………………………………………….60 2.2.2. Types of comets……………………………………………………………………….61 2.2.2.1. A new classification scheme for the comets……………………………………..63 2.2.3. The Oort Cloud…………………………………….…………………………………..65 2.3. Trans-Neptunian Objects: The Outer Limit……………………….…………………..65 3. Meteorites…………………………………………………………………………………..68 3.1. Classification of meteorites…………………………………………………….……….69 3.1.1. Chondrite Meteorites…………………..……………………………………….……..69 3.1.1.1. Oxidation state and bulk oxygen isotopic compositions………………………...72 3.1.1.2. Petrologic type……………………………………………………………………….72 3.1.1.3. Refractory inclusions………………………………………………………………..73 3.2.1.4. Thermal metamorphism…………………………………………………………….73 3.2.1.5. Properties of the Chondrites meteorites………………………………………….74 3.2.2. Nonchondrite Meteorites……………..……………………………………………….78 FCUP Mineralogy of V-type Near Earth Objects 9 Page: 3.2.2.1. Primitive Achondrites……………………………………………………………….79 3.2.2.2. Differentiated Achondrites…………………………………………………………80 3.2.2.3. HED’s……………………………………….………………………………………..82 3.2.2.4. Irons…………………………………….…………………………………………….84 3.2.2.5. Stony-irons………………………………………………………….………………..85 3.3. The chronology of events in the early Solar System………………………………...85 4. The parent body search…………………………………………………………………..88 4.1. The meteorite-parent body relationship…………………………………………….…89 4.1.1. HED’s-4 Vesta-Vestoids……………….……………………………………………..89 4.1.2. Ordinary Chondrites…………………………………………………………………..92 4.1.3. Iron meteorites and Enstatite chondrites…………………………………………...93 4.1.4. CM and CI carbonaceous chondrites……………………………………………….94 4.1.5. Others Meteorites Groups………………………….………………………………..95 4.2. Space weathering………………………………………………………………………..96 4.3. Spectra Modelling: The Hapke Model and the Modified Gaussian Model………...98 4.3.1. Hapke Model………………………………………………………….………………..98 4.3.1.1. Concepts and Definitions…………………………………………………………..99 4.3.1.1.1. Radiance and Irradiance……………………………………………………...….99 4.3.1.1.2. Geometric Notions……………………………………………………..……….100 4.3.1.1.3. Cross sections and efficiencies………………………………………………..101 4.3.1.1.4. Particle Scattering Albedo and Espat Function………………………………103 FCUP Mineralogy of V-type Near Earth Objects 16 move a short distance in a direction away from either end toward the center of the Lshape area. Almost any change that is made to a particle that is initially clear, smoothsurfaced and spherical decreases b and increases c. Fig.4.13 – Types of reflectances and their symbols. Fig.4.14 – Scheme of the equivalent slab model for . Fig.4.15 – External (SE) and Internal (SI) surface reflectance coefficients versus the refractive index for . The solid lines denote the exact expressions while the dashed lines denote the approximations results. Fig.4.16 – The residual error between the 900nm orthopyroxene absorption feature models, using different values of n, for a reflectance spectra of a particle with a size . Fig.4.17 – Input parameters for the Hapke Model spreadsheet. Red colored cells represent the identification of the sample, orange colored cells represent the input spectra data, green colored cells represent sample specific values and orange colored cells represent assumptions of the Hapke Model. Fig.4.18 – General overview of the demo spreadsheet developed in Microsoft Excel, to derive the mineralogical composition of laboratory mixtures. Fig.5.1 – Reflectance spectra of the seven laboratory mixtures of orthopyxorene and clinopyroxene obtained from the RELAB database, with a grain size proportion ranging between 0-45μm. Fig.5.2 – Reflectance spectra of the seven laboratory mixtures of olivine, orthopyxorene and anorthosite obtained from the RELAB database, with a grain size proportion ranging between 45-75μm. Fig.5.3 – Reflectance spectra of the seven laboratory mixtures of olivine, orthopyxorene and anorthosite obtained from the RELAB database, with a grain size proportion ranging between 70-145μm. Fig.5.4 – Plot of the single scattering albedo versus wavelength, for all the 7 laboratory mixtures with the following end-members: anorthosite, orthopyxorene and olivine. FCUP Mineralogy of V-type Near Earth Objects 17 Fig.5.5 – Plot of the single scattering albedo versus wavelength, for all the 14 laboratory mixtures with the following end-members: clinopyroxene and orthopyxorene. Fig.5.6 – Plot of the single scattering albedo of the sample XP-CMP-010 (C1XP10) and their end-members (orthopyroxene (Opx) and clinopyroxene (Cpx)) versus wavelength. The selected files for orthopyroxene and clinopyroxene were PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30), respectively. The residuals (cyan line) were translated from 0 to 1. Fig.5.7 – Plot of the residuals of the sample XP-CMP-010 (C1XP10). Fig.5.8 – Plot of clinopyroxene (Cpx) versus orthopyroxene (Opx) nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 0 and 45μm. Fig.5.9 - Plot of orthopyroxene (Opx) versus olivine nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. Fig.5.10 - Plot of anorthosite versus olivine nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. Fig.5.11 - Plot of anorthosite versus orthopyroxene (Opx) nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. Fig.5.12 – Reflectance spectra of all three clinopyroxenes obtained from the RELAB database. Fig.5.13 – Single scattering albedo versus wavelength for all three clinopyroxenes. Fig.5.14 – Plot of clinopyroxene (Cpx) versus orthopyroxene (Opx) nominal and derived compositions for all fourteen laboratory mixtures with a grain size proportion between 0 and 45μm and 70 and 145μm. Fig.5.15 - Reflectance spectra of 4 Vesta and the four near-Earth vestoids. The reflectance spectra are incrementally shifted vertically by 0.5. FCUP Mineralogy of V-type Near Earth Objects 18 Fig.5.16 - Reflectance spectra of the selected ten diogenites (D), ten eucrites (E) and ten howardites (H). Eucrites’ reflectance spectra are shifted vertically by 1 unity. Howardites’ reflectance spectra are shifted vertically by 2 units. Fig.5.17 - Reflectance spectra of the five V-type asteroids versus respectively HED with the closest mineralogy. A vertical shift of 1 unit is introduced between different groups of spectra. Fig.5.18 – Plot of the derived single scattering albedos of the ten Diogenites (D), ten Eucrites (E) and ten Howardites (H). The single scattering albedo of Eucrites are shifted vertically by 1 unity, while the single scattering albedo of Howardites are shifted vertically by 2 units. Fig.5.19 – Plot of the derived single scattering albedo of the five selected V-type asteroids. The single scattering albedo of 3908 Nyx is shifted vertically by 0,5 units, of 4055 Magellan is shifted vertically by 1 unit, of (5604) 1992 FE is shifted vertically by 1,5 units and of (6611) 1993 VW is shifted vertically by 2 units. Fig.5.20 – Plot of the reflectance spectra of the pyroxene end-members used in this study. Reflectance spectra of Hypersthene are shifted moved vertically by 1 unity, of Augite are shifted vertically by 2 units, of Piogenite are shifted vertically by 3 units, of Enstatite are shifted vertically by 4 units, of Bronzite are shifted vertically by 5 units and of Wollastonite are shifted vertically by 6 units. Fig.5.21 – Reflectance spectra of the selected olivine and plagioclase end-members minerals, used in this study. Reflectance spectra of Anorthite are shifted vertically by 1 unity, of Olivine Fayalite are shifted vertically by 2 units and of Olivine Forsterite are shifted vertically by 3 units. Fig.5.22 – Reflectance spectra of the selected neutral phases used in this study. Fig.5.23 – Best fit model results of the selected V-type asteroids. Data points represent the single scattering albedo of the samples, while the black lines represent the best obtained fits. A vertical shift of 0.5 units is introduced between different asteroids/meteorites albedos. FCUP Mineralogy of V-type Near Earth Objects 19 Fig.5.24 – Plot of the reflectance spectrum of 4 Vesta with its appropriate endmembers minerals found in this study. All spectra are normalized by the reflectance value at 1.2μm. Fig.5.25 - Plot of the low-Ca pyroxenes content vs high-Ca pyroxenes content, of all thirty HED meteorites and all five V-type asteroids. Fig.B.1 – Pyroxene crystal structure. Yellow: SiO4 tetrahedra. Red: position occupied by smaller cations or Y. Blue: position occupied by bigger cations or X. Fig.B.2 – The nomenclature of the calcium, magnesium, iron pyroxenes. Fig.B.3 – The nomenclature of sodium pyroxenes. Fig.B.4 - The atomic scale structure of olivine looking along the a axis. Oxygen is shown in red, silicon in pink, and magnesium/iron in blue. A projection of the unit cell is shown by the black rectangle. Fig.B.5 - Compositional phase diagram of the different minerals that constitute the feldspar solid solution. Fig.E1 – Derived albedo for the sample XP-CMP-011 (C1XP11) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E2 – Plot of the residuals of the sample C1XP11. Fig.E3 – Derived albedo for the sample XP-CMP-012 (C1XP12) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E4 – Plot of the residuals of the sample C1XP12. Fig.E5 – Derived albedo for the sample XP-CMP-013 (C1XP13) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E6 – Plot of the residuals of the sample C1XP13. FCUP Mineralogy of V-type Near Earth Objects 20 Fig.E7 – Derived albedo for the sample XP-CMP-014 (C1XP14) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E8 – Plot of the residuals of the sample C1XP14. Fig.E9 – Derived albedo for the sample XP-CMP-015 (C1XP15) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E10 – Plot of the residuals of the sample C1XP15. Fig.E11 – Derived albedo for the sample XP-CMP-016 (C1XP16) and for their extreme minerals PP-CMP-021 (C1PP21) and PP-CMP-030 (C1PE30) versus wavelength. The residuals were shifted by one unity. Fig.E12 – Plot of the residuals of the sample C1XP16. Fig.E13 – Derived albedo for the sample XT-CMP-033 (C1XT33) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E14 – Plot of the residuals of the sample C1XT33. Fig.E15 – Derived albedo for the sample XT-CMP-034 (C1XT34) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E16 – Plot of the residuals of the sample C1XT34. Fig.E17 – Derived albedo for the sample XT-CMP-035 (C1XT35) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E18 – Plot of the residuals of the sample C1XT35. Fig.E19 – Derived albedo for the sample XT-CMP-036 (C1XT36) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. FCUP Mineralogy of V-type Near Earth Objects 21 Fig.E20 – Plot of the residuals of the sample C1XT36. Fig.E21 – Derived albedo for the sample XT-CMP-037 (C1XT37) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E22 – Plot of the residuals of the sample C1XT37. Fig.E23 – Derived albedo for the sample XT-CMP-038 (C1XT38) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E24 – Plot of the residuals of the sample C1XT38. Fig.E25 – Derived albedo for the sample XT-CMP-039 (C1XT39) and for their extreme minerals OL-JMS-001 (C1OL01), PE-CMP-031 (C1PE31) and PA-CMP-060-B (CBPA60) versus wavelength. The residuals were shifted by 1.1 unities. Fig.E26 – Plot of the residuals of the sample C1XT39. Fig.E27 – Derived albedo for the sample XP-CMP-001 (C1XP01) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E28 – Plot of the residuals of the sample C1XP01. Fig.E29 – Derived albedo for the sample XP-CMP-002 (C1XP02) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E30 – Plot of the residuals of the sample C1XP02. Fig.E31 – Derived albedo for the sample XP-CMP-003 (C1XP03) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E32 – Plot of the residuals of the sample C1XP03. FCUP Mineralogy of V-type Near Earth Objects 22 Fig.E33 – Derived albedo for the sample XP-CMP-004 (C1XP04) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E34 – Plot of the residuals of the sample C1XP04. Fig.E35 – Derived albedo for the sample XP-CMP-005 (C1XP05) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E36 – Plot of the residuals of the sample C1XP05. Fig.E37 – Derived albedo for the sample XP-CMP-006 (C1XP06) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E38 – Plot of the residuals of the sample C1XP06. Fig.E39 – Derived albedo for the sample XP-CMP-007 (C1XP07) and for their extreme minerals PP-CMP-023 (C1PP23) and PE-CMP-032 (C1PP32) versus wavelength. The residuals were shifted by one unity. Fig.E40 – Plot of the residuals of the sample C1XP07. Fig.E41 - Unnormalized reflectance spectrum of 4 Vesta compared with the reflectance spectra of ten diogenites (D), ten eucrites (E) and ten howardites (H). Reflectance spectra of: diogenites (D) are shifted vertically by 0.7; eucrites (E) are shifted vertically by 1.2; howardites (H) are shifted vertically by 2.2. All spectra are normalized by the reflectance value at 1.2μm. The abbreviature “gsp” means grain size proportion. Fig.E42 - Unnormalized reflectance spectrum of 3908 Nyx compared with the reflectance spectra of ten diogenites (D), ten eucrites (E) and ten howardites (H). Reflectance spectra of: diogenites (D) are shifted vertically by 0.7; eucrites (E) are shifted vertically by 1.2; howardites (H) are shifted vertically by 2.2. All spectra are normalized by the reflectance value at 1.2μm. The abbreviature “gsp” means grain size proportion. FCUP Mineralogy of V-type Near Earth Objects 23 Fig.E43 - Unnormalized reflectance spectrum of 4055 Magellan compared with the reflectance spectra of ten diogenites (D), ten eucrites (E) and ten howardites (H). Reflectance spectra of: diogenites (D) are shifted vertically by 0.7; eucrites (E) are shifted vertically by 1.2; howardites (H) are shifted vertically by 2.2. All spectra are normalized by the reflectance value at 1.2μm. The abbreviature “gsp” means grain size proportion. Fig.E44 - Unnormalized reflectance spectrum of (6611) 1993 VW compared with the reflectance spectra of ten diogenites (D), ten eucrites (E) and ten howardites (H). Reflectance spectra of: diogenites (D) are shifted vertically by 0.7; eucrites (E) are shifted vertically by 1.2; howardites (H) are shifted vertically by 2.2. All spectra are normalized by the reflectance value at 1.2μm. The abbreviature “gsp” means grain size proportion. Fig.E45 - Unnormalized reflectance spectrum of (5604) 1992 FE compared with the reflectance spectra of ten diogenites (D), ten eucrites (E) and ten howardites (H). Reflectance spectra of: diogenites (D) are shifted vertically by 0.7; eucrites (E) are shifted vertically by 1.2; howardites (H) are shifted vertically by 2.2. All spectra are normalized by the reflectance value at 1.2μm. The abbreviature “gsp” means grain size proportion. Fig.E46 – Top: Plot of the best fit obtained of a diogenite meteorite (sample MP-TXH068-A). Down: Residuals plot of the fit of the sample MP-TXH-068-A. Fig.E47 – Top: Plot of the worst fit obtained of a diogenite meteorite (sample MP-TXH071-A). Down: Residuals plot of the fit of the sample MP-TXH-071-A. Fig.E48 – Top: Plot of the best fit obtained of a eucrite meteorite (sample MB-TXH-069B). Down: Residuals plot of the fit of the sample MB-TXH-069-B. Fig.E49 – Top: Plot of the worst fit obtained of a diogenite meteorite (sample MP-TXH070-D). Down: Residuals plot of the fit of the sample MP-TXH-070-D. Fig.E50 – Top: Plot of the best fit obtained of a howardite meteorite (sample MB-TXH053-A). Down: Residuals plot of the fit of the sample MB-TXH-053-A. FCUP Mineralogy of V-type Near Earth Objects 24 Fig.E51 – Top: Plot of the worst fit obtained of a howardite meteorite (sample MB-TXH068-D). Down: Residuals plot of the fit of the sample MB-TXH-068-D. Fig.E52 – Top: 4 Vesta single scattering albedo and its modelled curve. Down: Residuals plot of the fit. Fig.E53 – Top: 3908 Nyx single scattering albedo and its modelled curve. Down: Residuals plot of the fit. Fig.E54 – Top: 4055 Magellan single scattering albedo and its modelled curve. Down: Residuals plot of the fit. Fig.E55 – Top: 5604 (1992) FE single scattering albedo and its modelled curve. Down: Residuals plot of the fit. Fig.E56 – Top: 6611 (1993) VW single scattering albedo and its modelled curve. Down: Residuals plot of the fit. FCUP Mineralogy of V-type Near Earth Objects 25 Indices of Tables Table C1 – Mineralogically characterized asteroids, which have a possible meteorite affinity. Table C2 – Detailed description about the taxonomic classes of the Tholen and SMASSII classification schemes. Table C3 – Classification of asteroids by the radii of their orbits around the Sun. Table C4 – Classification scheme for comets based on their aphelion. The Encke-type is defined by E, short period by SP, intermediate-period by I and long-period by L. Table C5 – Definition of the Tisserand parameter classification scheme. The upper limit for TP is around 2.8, limit above for which is impossible for an object to be directly injected to a single encounter. Table C6 – Classification scheme of meteorites. Table C7 – Meteorite groups and numbers of their members. Number of meteorites was taken from Grady (2000). Table C8 – Meteorite groups and their composition characteristics. Table C9 – Meteorite groups and their postulated parent bodies or source bodies. Table C10 – Space weathering and asteroid types. Table D1 – Selected Laboratory Mixtures with their grain size proportions and the mineralogical composition derived by Duffard (2005) and the mineralogical compositions derived in this work using the Solver software of Excel, with the residuals of their derivation. Table D2 – Selected Laboratory minerals with their grain size proportions, to use in the Solver software of Excel to derive the mineralogical compositions of all Laboratory Mixtures presented in Table D1. FCUP Mineralogy of V-type Near Earth Objects 32 proposing that the planets orbit around the Sun – heliocentric model. Because of his idea, he was charged with impiety by Cleanto. Only several centuries later, Nicolau Copernic (1473-1543) argued in his favor. Copernic proposed that: (i) the planets orbit the Sun and rotate around its rotation axis; (ii) the dimensions of the Solar System can be neglected when compared to the dimensions of the Universe. Some years later, Galileu Galilei (1564-1642) observed some planets and stars with a telescope, for the first time in the humankind history. He confirmed the Copernic theory’s and then introduced the term dynamics. After this, several theories about the Solar System formation were proposed: turbulent theory (proposed by René Descartes, 1596-1650); the nebular theory (proposed by Immanuel Kant (1724-1804) and Marquis de Laplace (1749-1827)); tidal theory (proposed by Buffon (1707-1788) and later on modified by Sir James Jeans); accretion theory (proposed Alfvén and Arrhenius in the last century). 1.1.1. Turbulent Theory The turbulent theory was introduced by René Descartes (1596-1650), in which the Universe was made by ether and matter. In his theory the Universe was full of swirling movements. However, this theory was discarded as it did not explain the presence of the ecliptic, later confirmed by the Newton’s laws. In its theory, René Descartes introduced the term evolution. 1.1.2. Nebular Theory The existence of a nebula from which the planets formed was firstly proposed by Immanuel Kant. Several years later Marquis de Laplace proposed a similar theory. He predicted that: (i) the orbits of all planets are roughly in the same plane; (ii) all planets orbit around the Sun in the same direction; (iii) all orbits have a low eccentricity, meaning that they are close to circles; (iv) all planets spin on their axis in the same direction. In his theory, the beginning point was a slowly rotating spherical nebula made of gas and dust, which were slowly collapsing due to the gravitational force (step a, in Fig. 1.1). When the nebula collapsed it spun faster and flattened along the spin axis, in order to conserve the angular momentum (step b, in Fig. 1.1). After this, maybe the nebula acquired a lenticular shape and the material found in the broadside regions was free to orbit the central mass (step c, in Fig. 1.1). As the collapse was undergoing, material was left behind in the equatorial plane forming a set of annular rings (step d, in FCUP Mineralogy of V-type Near Earth Objects 33 Fig. 1.1). In the end, each ring originated a planet. Eventually, the Sun was formed in the central region (step e, in Fig. 1.1) [11]. This theory was capable of explaining the observations made during the eighteenth century. In addition, it is capable to explain why all planets orbit around the Sun in the same direction and why their orbits are almost circles lying in the same plane, called ecliptic. However, Laplace’s theory did not explain how a planet formed from a ring of material. Actually, it also did not provide an explanation for one intriguing fact: why 99.9% of the Solar System’s mass is in the Sun and 99% of the Solar System’s angular moments is in planets. Fig.1.1 – A schematic overview of the Laplace’s nebular theory. (a) A slowly rotating and collapsing gas-and-dust sphere. (b) An oblate spheroid formed as the spin axis increased. (c) The critical lenticular form. (d) Rings were left behind in the equatorial plane. (e) Each ring originated a planet [11]. Several modifications to Laplace’s theory were attempted in order to bridge the latter evidence. Berlage introduced a viscosity effect, while Hoyle and Schatzman introduced the magnetic field effect and the solar wind, respectively, attempting to give one plausible explanation which could bridge the problem of the angular momentum. FCUP Mineralogy of V-type Near Earth Objects 34 1.1.3. Tidal theory Buffon (1707-1788) proposed that a comet shocked with the Sun, 70Myr ago. This hypothesis has no scientific fundament. But, in the beginning of the twentieth century, Sir James Jeans reformulated his theory by substituting the comet by a massive star. Jeans separated the formation of the Sun from the formation of the planets. The mechanism proposed by Sir James Jeans is illustrated in Fig. 1.2, where a massive star passing close to the Sun, raised a tide on the Sun pulling out material, which formed a filamentary structure. This filament was gravitationally unstable and as its mass was greater than the Jeans critical mass, it condensed to form planets. In the end, the planets were left in orbits around the Sun. Fig.1.2 – A schematic overview of the tidal theory, modified by Sir James Jeans [11]. If this theory is correct, the presence of planetary systems around stars should be rare, which contrasts with growing up number of discovered exoplanets in the last few years. FCUP Mineralogy of V-type Near Earth Objects 35 1.1.4. Accretion theory Alfven and Arrhenius argued that the Sun had encountered two nebulae: one composed by volatile elements and another composed by hydrogen. They proposed that the Sun accreted matter from these nebulae and several collisional processes between particles led to the formation of flows from which the planets could have formed. 1.2. Formation and evolution of the Solar System: Planet Formation Modern theories about the formation and evolution of planetary systems are based on the observations of star-forming regions and numerical simulations. In the case of the Solar System, there are several observational constraints supporting the most accepted theory, which will be described in this chapter. Some key observations are:  The Solar System contains eight large planets, with almost circular and coplanar orbits, lying in the region between 0.4-30 AU. There are also a few locations between the planets where additional large objects could exist on stable orbits.  All planets are grouped in two distinct regions: in the inner region of the Solar System there are the volatile-poor planets while the volatile-rich planets are found further out. The Main Asteroid Belt is found between these two regions (2-4 AU). The Main Asteroid Belt is substantially depleted in mass compared to the other two regions.  The degree of fractionation decreases with distance from the Sun: the terrestrial planets and inner-belt asteroids are highly depleted in volatiles while the outerbelt asteroids are less so and many satellites in the outer Solar System are ice rich.  Ancient planetary surfaces, such as those of Mercury, Mars or Callisto, are covered by impact craters.  The terrestrial planets and many asteroids have undergone differentiation.  There is strong evidence that Saturn is highly centrally condensed, containing a core of mass of 10MEarth and weaker evidence that Jupiter has a core of similar FCUP Mineralogy of V-type Near Earth Objects 36 mass. These cores may have masses comparable to the ones of Uranus or Neptune. The observational evidences led the development of the Solar System formation theory postulated by Laplace, described in the last section. In this theory, planets formed from a protoplanetary nebula by pairwise accretion of small objects. A gravitational instability in some regions of this disk is added to this theory to explain the formation of the giant planets. However, before going deeper in the description of this theory it is important to quote that the standard model only attempt to explain the planets observed in the Solar System. 1.2.1. The protoplanetary nebula and the first solids The Solar System probably formed from the collapse of a fragment of a molecular cloud. When the Jeans’ criterion was satisfied, the fragment of the molecular cloud collapses. These collapse may have occurred spontaneously or been triggered by an external factor, such as a supernova [12]. As the cloud fragment collapses the bulk of its mass fell to the center to form the protosun, while the remaining material formed a rotationally supported disk. The presence of this disk is constrained by the observations of young “T-Tauri” stars, with ages less than 10Myr, which have optically thick disks of gas and dust with masses of 0.001-1MSun [13]. However, the question concerning to initial mass of the disk is still under debate. These disks have spectra containing absorption features caused by the presence of water ice and silicates. According to the observations, disks are not observed around stars older than 10Myr, providing an upper limit of the solar protoplanetary nebula lifetime. But, as young stars are observed to be accreting material and due to the limited lifetime of circumstellar disks, the solar protoplanetary nebula is viewed by this model as a viscous accretion disk in which the material is transported radially inwards, ultimately falling onto the Sun [14]. However, the source of viscosity in the disk is still under debate. Nowadays, the maximum temperature reached in the Sun’s protoplanetary disk is still under current discussion. As it is constrained by the approximate isotropic FCUP Mineralogy of V-type Near Earth Objects 37 homogeneity of planetary material in the inner Solar System, most of the material in the inner regions were vaporized and mixed. As the nebula cooled and refractory material began to condense, chemistry becomes important. The sequence of the condensation of different elements can be seen in Fig. 1.3. The list along the top of Fig. 1.3 shows the 15 most abundant elements in the protoplanetary nebula. Below, there is a list of compounds in which these elements were mostly found when the nebula temperature’s was around 2000K. As an illustrative example, oxygen atoms were enclosed in water molecules. The staircase found across Fig. 1.3 separates the solid phase from the gas phase. By looking down in the staircase from the top right to the bottom left we can follow the sequence in which the elements condensed as the nebula cooled [15]. Fig.1.3 – Predictions of the sequence in which chemical species condensed in the solar nebula [15]. The first elements which condensed were metals with high boiling points as there are the cases of the calcium, aluminum and nickel. The first two cases condensed as oxides CoO and Al2O3, while nickel condensed as solid nickel. The arrows below the staircases denote how these solids reacted with other elements, which were still in the gas phase. As it can be observed by this diagram in Fig.1.3, a major quantity of aluminum ended up in the silicate mineral feldspar. The presence of FCUP Mineralogy of V-type Near Earth Objects 38 minerals like olivine and feldspar is reassuring as Earth is made by them. The others compounds with lower boiling points started to condense as the nebula cooled, even some compounds as water, ammonia and methane. The complexity of this process as a whole is evidenced by the possible routes for water molecules: it started to enter the solid phase when the temperature was still high enough by becoming incorporated in the hydrated minerals tremolite and serpentine. Furthermore, water molecules only formed pure water ice when the nebula cooled. At this stage, in the inner region of the nebula, around 0.5% of the mass would have formed solids following any initial hot phase of the nebula evolution. These solids were mainly composed by silicates, metal and sulfides. In the outer regions, where temperatures were colder, icy materials also condensed. In this region, water ice was the most abundant. The boundary between these two regions is established by a discontinuity in the surface density of solid material called the “ice line” or “snow line”. This line would have moved inwards over time as the nebula cooled. 1.2.2. The origin of planetesimals This earliest stage of planetary accretion is the most poorly understood at the present. Nowadays it is argued that planetesimals could be originated by one of the following processes: via gravitational instability or via dust grain sticking. This stage of the nebula evolution is marked by the lacking of gas pressure support, settling dust grains and aggregates towards the midplane of the nebula, with larger objects falling faster than small ones. During this process, dust grains would have coagulated to form larger objects, increasing the rate of sedimentation. As a consequence, the solid to gas ratio in this region was increased. If the inner layer of the nebula was very thin, portions of it would become gravitationally unstable and collapsed to form solid bodies of around 1 km in diameter. These bodies are called “planetesimals”. The formation planetesimals via gravitational instability are not so simple and are more probable to occur when the volume density of solid particles is high and their relative velocities are low. In one hand, if the disk was nonturbulent, particles would have migrated at different rates in different regions, leading to a pile up of solid material in certain regions [16]. On the other hand, particles would migrate more slowly in FCUP Mineralogy of V-type Near Earth Objects 39 regions where the volume density of particles is high, probably leading to a local increase of the solid-to-gas ratio [17]. These two factors favor the formation of planetesimals via gravitational instability. If planetesimals were not formed through gravitational instabilities, such bodies are probably a result of a sticking process of dust grains and aggregates during collisions. Experimental studies suggest that small dust grains will stick together if they collide at velocities lower than a few meters per second [18]. This sticking process is mainly due to van de Waals forces, which are weak. The grain growth process may have occurred in a short timescale for the objects to survive against gas drag due to the differential velocity of solid objects relative to the gas nebula. Nowadays it is still unclear how aggregation continued for boulder-sized and larger objects. But, the presence of gas may have helped. 1.2.3. The formation of terrestrial planets Once solid bodies reached a size of around 1km, gravitational interactions between them become significant. Close passages between planetesimals tended to increase their random velocities ν, as a result of mutual gravitational attraction. The frequency of collisions is strongly dependent on their random velocities. In one hand, if ν was high, close encounters were briefly reducing the probability of collision and thus the growth rate of planetesimals. On the other hand, if ν was small during a close encounter, planetesimals remained close to each other, increasing the probability of collision. In this scenario, three growth modes may have occurred during this stage, each at different times in the solar nebula. The early stages of accretion were marked by a “runaway growth”. Actually, if the disk was dynamically too cold, the velocity dispersion of the planetesimals, , may be inferior to their escape velocity, . In this case a “runaway growth” begins, which means that relative mass growth of each individual object is an increasing function of its own mass M [19]: FCUP Mineralogy of V-type Near Earth Objects 40 In the “runaway growth” mode, the largest planetesimals tended acquire the smallest random velocities, due to the dynamical friction process. Hence, the largest bodies experienced strongest gravitational focusing of their trajectories and they grew more rapidly, during this stage. Once the largest objects, become around 100 times more massive than a typical planetesimal (i.e. dubbed “planetary embryos”), the random velocities of planetesimals were mainly determined by gravitational perturbations from the embryos. So, when the “runaway growth” slowed down and accretion continued in an “oligarchic growth” mode (in which the relative mass growth for the largest objects is proportional to ) [20]. During this growth mode, the larger a planetary embryo was, the more it stirred up the velocities of nearby planetesimals and the slower it grew up. Gravitational interactions between planetary embryos tended to keep them in their “feeding zones” in the protoplanetary disk, where they accreted most of their mass. Hence, depending on where planetary embryos formed in the protoplanetary disk, they may have had different compositions. The final stage of accretion is the most violent phase, marked by catastrophic collisions between planetary embryos [21]. This last stage began when the remaining planetesimals had too little mass to sustain the random velocities of planetary embryos. While the random velocities of the planetary embryos increased, their growth slowed drastically and their orbits began to cross those of their neighbors. A new growth mode began, when almost half of the total solid mass was contained in planetary embryos [22]. This growth mode is known as “orderly growth”. In this last stage, planetary embryos started to collide with each other and despite the unknown efficiency of these collisions, growth continued. However, this process is supported by numerical simulations [23]. The remaining planetesimals probably fell into the Sun or were ejected from the infant Solar System. Several impacts may have been so energetic that had completely molten each of the inner planets, originating a magma ocean which homogenized existing material and erased any chemical signatures of the earlier stages of accretion. One example of a catastrophic collision may have occurred with Mercury after it had differentiated, constraining its observed high density [24]. Due to the heat from the decay of radioactive isotopes allied with energy released from impacts, planetary embryos were differentiated once these objects FCUP Mineralogy of V-type Near Earth Objects 41 became partially molten [25]. Iron and siderophile elements (such as, platinum, palladium and gold) preferentially sank to the center to form a core, while the lighter silicates and lithophile elements formed a mantle. As it is evidenced by the presence of siderophile elements in Earth’s mantle and crust, some material was accreted as a “late veneer” when Earth’s differentiation was largely complete. However, the presence of volatile elements on Earth poses a major problem on planetary formation. As it was already described above, the inner region of the protoplanetary nebula may be too hot for the volatiles elements to condense at the same time as planetesimals were being formed. Hence, it seems likely that Earth acquired its volatile elements by accreting material which formed in regions beyond 2.5 AU, i.e. in regions where the nebula was cold enough for ices to condense. Another important aspect which has to be stated out is how Earth probably acquired water. As it constrained by the difference between the measured D/H ratio of seawater on Earth and the one’s of comets, an asteroidal source is more promising to explain the origin of water on Earth. Actually, according to numerical simulations, Earth could have accreted several oceans worth of water from the asteroid belt, in particular if lunar-to-Mars size planetary embryos had formed in this region [26]. 1.2.3.1. Origin of the Moon One of the most intriguing properties of the Solar System is the origin of the Moon itself. Between the more than 100 moons in the Solar System, our Moon is an object of peculiar interest as the ratio between Earth’s and its masses is quite large. Most of the moons in the Solar System are captured objects: see for example the case of Phobos and Deimos (moons of Mars), which are probably captured asteroids, or also the case of Triton (moon of Neptune). In addition, we can also count lots of small icy worlds orbiting Jupiter, Saturn and Uranus which are probably captured objects by the gravitational influence of these planets. Historically, there are three theories for the origin of the Moon: the capture theory, the fission theory and the cocreation theory. In the capture theory, the Moon formed elsewhere in the protoplanetary disk and was consequently captured by the Earth’s gravitational field. However, the major contradiction against this theory is the FCUP Mineralogy of V-type Near Earth Objects 48 consequence of the latter process, a silicate crust in these bodies was formed. With the ending of this internal activity, only several collision processes in the Main Asteroid Belt could modify their characteristics: as an example, the objects with lower dimensions shocked with the ones of bigger dimensions originating a regolith in their surfaces [49]. FCUP Mineralogy of V-type Near Earth Objects 49 2. Small Bodies in the Solar System “The clear and very uplifting lesson of all this is that we remain only dimly aware even of our immediate astronomical backyard. The Kuiper belt with its unexpected dynamical subdivisions and structures, the main-belt comets, the (perhaps) densely populated inner Oort cloud, were all completely unknown when, for example, most of the readers of this article were born. The obvious next question is “What else is out there?”.” in Jewitt, D. (2009), Icy bodies in the New Solar System, Proceedings of the International Astronomical Union, IAU Symposium, 263, 14 he last 20 years have been particularly fruitful in the study of these icy objects. Due to several advances in technology it was possible to reveal the presence of such objects in regions thought to be empty. Nowadays, it is known that they carry an unimaginable fraction of scientifically useful information about the origin and the evolution of the Solar System. Firstly, the majority of these objects escaped substantial thermal alteration since its formation. Their chemical and molecular compositions may be closely to the initial conditions of the Solar System. Secondly, their number is so high and their orbits so accurate that they can be used with the purpose of mapping dynamical space parameters or to trace processes occurring in the protoplanetary disk [50]. In our Solar System, there are a countless number of small objects, ranging from planetesimals which grew up until modest sizes of a few hundreds of kms in radius to post-collisional debris. These objects are known as asteroids, comets or Trans-Neptunian Objects. They are located in the three ice reservoirs: the Oort cloud; the Kuiper Belt; and in the Main Asteroid Belt. The presence of such small bodies in our Solar System seem to be quite common in other planetary systems, as the number of known exo-Kuiper-Belts around nearby stars has recently increased from 29 to 34 [51]. As an illustrative example, the HR8799 planetary system resembles to a scaledup version of the Solar System [52]: the presence of an asteroid belt interior to and a Kuiper Belt exterior to the recently discovered planets was inferred [53]. In this scenario, the obvious question is “What else is out there?”. T FCUP Mineralogy of V-type Near Earth Objects 50 In this chapter, the small icy objects of the Solar System will be described. This chapter is divided into three subsections. In the first subsection, asteroids will be presented. Then, in the second subsection, the “dirty snowballs”, also known as comets and what is their role in our Solar System will be pointed out. Finally, the last subsection will be dedicated to the Trans-Neptunian Objects. Keywords: asteroids; Main Asteroid Belt; taxonomy; Near Earth Asteroids; Kirkwood gaps; Trojan Asteroids; Centaurus Asteroids; 4 Vesta; Vestoids; families of asteroids; secular resonances; comets; long-period comets; short-period comets; Tisserand parameter; Oort Cloud; Trans Neptunian Objects; 2.1. Asteroids: Shattered Worlds The term asteroid is applied to a peculiar class of objects, which are sufficiently non volatile not to be classified as a comet and small enough not to be classified as a planet. In addition, they were called asteroids as they seemed to be star-like points of light. Nowadays, asteroids provide our only in situ record of the conditions and processes that the inner portions ( 1.8-3.5 AU) of the late solar nebula and the infant solar system have experienced [54]. Asteroids have irregular-shaped bodies, which in some cases they look like potatoes. Thousands of these objects orbit the Sun in the so known Main Asteroid Belt (hereafter, also reported as MAB) (Fig. 2.1), between the orbits of Mars and Jupiter [55]. Due to the several advances in technology, we are now able to get measurements of asteroid light-curves, as well as other important information such as body shapes, polar orientations or rotational periods. Ceres was the first asteroid to be discovered in 1801 by the Italian astronomer Giuseppe Piazzi. As it was predicted by the Titius-Bode rule 1 , the semi-major axis of Ceres’ orbit is 2.8 AU. Ceres is the biggest known rock in the MAB, with a diameter of 950 km (around 560 miles). In 1802, a second asteroid named Pallas was discovered. Ceres and Pallas were found really close to each other suggesting that they were a result of a fragmentation process of a larger body that could be split by several collisional processes. Hence, this evidence has encouraged people to look for more possible fragments, thus leading to the discovery of Juno (1804) and 4 Vesta (1807). 1 Titius Bode Rule: gives the mean orbital distance of the planets and which can be written in the Blagg-Richardson formulation as , where is the distance of the nth planet from the Sun (in AU), A is the mean ratio between two consecutive planetary distances and . The MAB distance is considered as a regular planetary distance AU. FCUP Mineralogy of V-type Near Earth Objects 51 Nowadays, many thousands of asteroids are catalogued [56]. In Table C1 some mineralogically characterized asteroids with some meteorite affinity associated can be found. Fig.2.1 – Main Asteroid Belt and Trojan Asteroids. Source: http://www.aerospaceweb.org/question/astronomy/q0270.shtml 2.1.1. Classification of Asteroids Asteroids can be classified by their location in the Solar System (which will be discussed in the next section of this chapter) or by their reflectance spectra. There are two taxonomic classifications, which are the most used to classify asteroids taking into account their reflectance spectra: Classification of Tholen; Classification of SMASSII. A brief description of these taxonomic classifications can be also found in Table C2. The taxonomic classification proposed by David J. Tholen [57], in 1984, is the most commonly used to classify asteroids by their spectral type. He studied the spectra of approximately 600 asteroids, in the range between 300 nm and 1100 nm, obtained by ECAS (Eight-Color Asteroid Survey). In this classification scheme, spectral types of asteroids are nominated by capital letters E, S, M, V, A, C, P, D and Z. Each spectral type has several linkages with a reflectance spectrum of a certain meteorite. This classification scheme tends to identify the asteroids from which different irons and stony-irons meteorites could have FCUP Mineralogy of V-type Near Earth Objects 52 been originated. This possible relationship between asteroids and irons and stony-irons meteorites will be discussed in chapter 4. The distribution of different spectral types of asteroids is not uniform in the MAB. As an example, the abundance of the E-type asteroids decreases from 60% of all known asteroids at 1.8 AU to only 2% at 2.15 AU (Fig. 2.2). Most E-type asteroids tend to have high reflectance and relatively featureless spectra [50]. S-type asteroids are more abundant at 2.3 AU declining to 20% at 3.0 AU and reaching 0% at 3.25 AU (Fig. 2.2). The S-type asteroids have moderate albedo and are mainly composed by silicates (Fe and Mg) plus metal (Fe-Ni). This type of asteroids can represent the fusion crust of a differentiated body, in which they had experienced at least partial melting in their parent body [58]. By the analysis of S-type asteroids survey spectra, seven mineralogical subtypes of the S-taxonomy were reported: the silicate assemblages ranged from nearly monominerallic olivine (subtype S(I)) through basaltic silicates (subtype S(VII)). Only the subtype S(IV) did not require igneous processing to produce the observed assemblage [58]. C-types asteroids as well as P and D-types are the most abundant in the region between 3.0 and 5.0 AU and contain organic material such as kerogen [56]. They also contain chemically bound water, clay minerals (montmorillonite), organic polymers, magnetite and sulfides. C-type asteroids have low albedo and are mainly composed of hydrated silicates and carbon. This type of asteroids possibly has a similar composition to the solar nebula. P and D-types may also contain mixtures of various kinds of ices and other organic compounds. Only one meteorite was catalogued as being possibly linked to the D-type asteroids: Tagashi Lake meteorite, which fell on 16 of January of 2000 in British Columbia, Canada [59]. A-type asteroids are usually reported as having olivine-rich assemblages, as they only have a single absorption feature near 1.05 µm [60]. The SMASSII taxonomic classification was recently introduced by Schelte L. Bus and Richard P. Binzel, in 2002, based in SMASS (Small Main-Belt Asteroid Spectroscopic Survey) with 1447 asteroids. In their work, they obtained higher resolution spectra of asteroids comparing to ECAS. Despite their observations were only in the range between 440 nm and 920 nm, they were capable of distinguishing several particular characteristics. In order to conserve the taxonomic classification FCUP Mineralogy of V-type Near Earth Objects 53 proposed by Tholen, they catalogued 26 types of different asteroids [61]. For a comparison between these two classifications schemes, please see Table C2. Fig.2.2 –The fraction of the total population of the different spectral types of asteroids in the MAB, varying with the distance from the Sun [62]. 2.1.2. Orbits of Asteroids Asteroids can also be classified based on their location within the Solar System and thus by their orbits (for further details, please check Table C3). Most asteroids orbit the Sun between 2.2 and 3.3 AU, which from the third Kepler’s law corresponds to a revolution period between 3.3 to 6 years, in the so known MAB. The belt is divided in three regions taking into account the average distance (a) of asteroidal orbits: Inner Belt: a < 2.5 AU; Middle Belt: 2.5 < a <3.1 AU; Outer belt: 3.1 < a <4.1 AU. However, asteroids can also be found close to the orbit of Jupiter (which are called Trojans), between Saturn and Uranus (which are called Centaurus) and close to Earth’s orbit (also known by Near Earth Objects, also reported by the acronym NEOs). In some regions of the MAB, some kind of gaps can be found (Fig.2.3). In 1866, Daniel Kirkwood explained these gaps for the first time to our knowledge as being a result of a resonance effect between the revolution periods of asteroids that occupied those regions and that of Jupiter. These gaps are also known as the Kirkwood gaps (Fig. 2.3). Several examples of these gaps can be counted in the MAB: the gap at 3.28 AU corresponds to a revolution period of 5.93 years, which is half of the Jupiter’s period; the gap at 2.5 AU corresponds revolution period of 3.95 years, which is one FCUP Mineralogy of V-type Near Earth Objects 54 third of Jupiter’s period. There are also two gaps corresponding to two-fifths and threefifths of Jupiter’s orbital period (resonance 3:1). These gaps can be explained by a simple fact: synchronous tugs from a large body on smaller particles gradually move those particles out of their orbits, like if it was “vanishing” and “cleaning” those regions. Actually a similar example to these gaps in the MAB is the Cassini division in Saturn’s rings. Fig.2.3 – Kirkwood gaps in the MAB [62]. The NEA’s orbit the Sun in an orbit close to the Earth’s one. This fact can be easily explained by the existing gaps (Fig. 2.3) in the MAB, where due to the resonance with Jupiter, these asteroids were completely pushed away from their initial orbits. There are three known populations of NEA’s: Aten, Apollo and Amores. The orbits of the asteroids from the Aten group lie inside Earth’s orbits as their perihelion distance is less than 1 AU. Their diameters range from 0.2 up to 3.4 km and their reflectance spectra (C and S type) resembles to those of the stony meteorites [56]. 5604 1992 FE is one exemplar of an asteroid belonging to the Aten group, classified as a V-type asteroid. Asteroids from the Apollo group cross the Earth’s and Mars orbit as their perihelion distance is less than 1 AU and their aphelion distance is greater than 1.52 AU. Their diameters range from 2 to 10 km and their reflectance spectra (S-type) link their compositions to the stony meteorites. Several asteroids, such as 1866 FCUP Mineralogy of V-type Near Earth Objects 55 Sisyphus, 1685 Toro, 1864 Daedalus and 6611 (1993 VW) are exemplars of the Apollo group. The asteroids from the Amor group all have perihelion distances greater than 1 AU and aphelion distances exceeding the radius of Mars’s orbit. Their diameters range from 6km up to 30km and their reflectance spectra (C and S-type) are similar to those of stony meteorites. 433 Eros, 1036 Ganymede, 3908 Nyx and 4055 Magellan are four exemplars of this group. The Trojan asteroids are found near Jupiter’s orbit, close to two points of Lagrange, oscillating around them. This group is named in allusion to the wooden horse which concealed Greek fighters. The association between the configuration of the magnetic field and the Lagrange points restricts the Trojan asteroids loosely to a comparatively small volume of space partly within and partly outside of the orbit of Jupiter. These types of asteroids contain organic matter which gives them a reddish color. The five largest Trojan asteroids are: 684 Hektor (150x300km), 911 Agamemnon (with a diameter maybe of 138km), 617 Patroclus (140km), 1437 Diomedes (possibly having a diameter of 130km) and 1172 Aeneas (125km). However, all the ten largest asteroids have diameters greater than 100km and are usually linked to the C, D and P spectral types [62]. The Centaurus asteroids are usually associated as the transition point between asteroids and the icy worlds of the Edgeworth-Kuiper belt. Their chemical compositions (spectral types C and D) are similar to those of carbonaceous chondrites but may also contain ices of volatile compounds [62]. 2.1.3. 4 Vesta The asteroid 4 Vesta was discovered by the German astronomer Heinrich Wilhelm Olbers in 1807 and was named 4 Vesta by the mathematician Carl Friedrich Gauss. It has an ellipsoidal shape (Fig.2.4) with an average radius of km [63] and a volume of . Its orbit has a semi-major axis of 2.362 AU, with an eccentricity of 0.089, and an inclination of 7.133o. 4 Vesta has a rotation period of 5.3h and, as it turns, its reflectance spectrum appears to change, showing a heterogeneous surface. Its reflectance spectrum shows a strong band near 1000 nm, due to the presence of pyroxenes. However, its reflectance spectrum also shows FCUP Mineralogy of V-type Near Earth Objects 56 several diogenetic or even olivinic features, which are interpreted as 4 Vesta being a differentiated asteroid and also having several impact craters. The surface of 4 Vesta shows regions with lava composed by basalt eucritic material. Its impact craters have excavated across the extrusive eucritic crust an intrusive plutonic layer or maybe a plutonic body of diogenitic composition. 4 Vesta also shows a dichotomous surface as the east hemisphere is mainly composed by pyroxene-rich in magnesium and pyroxene-poor in calcium (similar to diogenites 2 ), while the west hemisphere is dominated by a pyroxene-rich in iron and calcium (similar to eucrites2). Also another feature which reinforces this dichotomy is: the east hemisphere is probably composed by excavated plutonic rocks (which were possibly originated due to a major impact), while the west hemisphere is possibly composed by some lava flows. Fig.2.4 – Mapping 4 Vesta using several images captured by the Hubble Space Telescope (also known as HST). Source: Ben Zellner (Georgia Southern University), Peter Thomas (Cornell University) and NASA. 2 A description about Howardites, Eucrites and Diogenites will be provided in the next chapter. FCUP Mineralogy of V-type Near Earth Objects 57 Images obtained by the HST showed a “huge” crater (Fig. 2.4) with a diameter of 460 km. This crater has a depth of around 8 km and a central peak with approximately 13 km of height. The color change rate of the crater indicates a possible transition from pyroxenes to olivines to pyroxene-rich in calcium with the topographic increasing depth of the crater [64]. According to the spectrum of 4 Vesta in the visible and IR, its surface exhibits several absorption features of Ca-low pyroxenes, which have a similar composition to eucrites [65]. It was also observed that its surface is composed by howardites. Another peculiar absorption feature centered in 506.5 nm, was observed in the 4 Vesta spectrum evidencing the presence of Ca-rich augite. Recently, the Dawn Mission flew-by the surface of 4 Vesta. This mission belongs to NASA’s Discovery Program. Dawn’s goal is to characterize the conditions and processes of the infant Solar System, by investigating two of the largest asteroids in the MAB: 1 Ceres and 4 Vesta. In the latter case, it is aimed to verify the basaltic nature both inferred from its reflectance spectrum and from the howardites, eucrites and diogenites meteorites (which are believed to be originated from 4 Vesta). In a close future, several results about the data analysis of this mission may be found among the literature. 2.1.4. Families of Asteroids Regarding the problem of the missing planet somewhere between Mars and Jupiter, a partial answer may have been finally found: during the formation of the Solar System, one or a few more planetary embryos suffered several collisions between them or bombardments that led to their fragmentation. Shock after shock, these bodies were probably reduced to a smaller size and volume. The collisional processes seem to play a very important role in the life of these small objects [66]. Actually, the socalled families of asteroids are the direct proof that these collisions really have occurred. It was first noted by Hirayama (1918) that asteroids can form prominent groupings in space of orbital elements. Asteroid families are dynamical associations which are the individual fragments of the full or partial collisional breakup of a parent body. Such associations are properly designated by genetic families when the family members have compositions which are FCUP Mineralogy of V-type Near Earth Objects 64 perihelion q, aphelion Q and Tisserand parameter TP, this new taxonomy provides an instantaneous classification for any comet-like body [84]. One way to think about comets is in terms of their Tisserand parameter, measured in relation to a certain planet. The Tisserand parameter is an approximation to the Jacobi constant (which is an exact integral of motion in the circular restricted three-body problem). The parameter is defined as: √ (2.1) where aj and a are the semi-major axis of the planet and of the comet, respectively. Also e and i are the eccentricity and inclination of the comet’s orbit. It is assumed that the inclination of Jupiter is . This new classification scheme that was proposed divides comets with perihelion distances less than AU in order to include the Encke (E), short-period (SP), intermediate-period (I) and long-period (L) types [84]. The division is based on the aphelion distance (please see Table C4, for further details) and the sub-division takes into account the Tisserand parameter classes (please see Table C5, for further details), denoted by Roman-type subscripts. SP-type comets are defined as having an aphelion distance Q ranging between 4 and 35 AU. This class includes Halley-type comets and several comets belonging to the JFCs. In order to distinguish them in the class it is used the Tisserand parameter. For example, a SPI comet is a short-period comet corresponding to a Halley-type, according to this convention. E-type comets lie in an orbit which does not approach closer than a few Hill radii to Jupiter. The members of this class have orbital features close to the Encke’s comet (Q 4.1 AU). I-type comets have aphelia distances between 35 and 1000 AU. One exemplar of this class is the Hale-Bopp (Fig.2.7, Left), which is categorized as II because its Tisserand parameter is 0.879. FCUP Mineralogy of V-type Near Earth Objects 65 The fourth and last class is the L-type comets. These comets have aphelion distances greater than 1000 AU. It is thought that the source of these type of comets if the Oort Cloud. 2.2.3. The Oort Cloud In 1950, Jan Oort (1900-1992) noticed that the aphelion distances of several long-period comets extended to distances between 50000 and 100000 AU from the Sun [85]. He postulated that the Solar System is surrounded by a very large region of space containing a high number of objects that are bound to the Sun. The objects that belongs the Oort Cloud are composed by ices, such as water, ammonia and methane. Based on the isotropic distribution of the arrival of long-period comets, he hypnotized that this region is spherical. In honor to Jan Oort, this structure was named by “The Oort Cloud”. However, this region does not resemble to a cloud at all, as it is constrained by the low number of objects per unit of volume. It is thought that the objects of the Oort Cloud formed close to the Sun and then were ejected to their present location due to several gravitational interactions with the gaseous giants protoplanets. It is estimated that the cloud contains 1011-1012 objects [85-86]. This estimative is calculated by comparing the rate of arrival of new long period comets and the rate of erosion that is expected from external perturbations due to the combined action of passing stars allied with the torque from our own galaxy [87]. The Oort Cloud may be disturbed by the gravity of passing stars in the vicinity of the Sun. Also, it is speculated that the Sun has a dark massive planet or companion called Nemesis [88], which can periodically disturb the orbits of several objects in the Oort Cloud. 2.3. Trans-Neptunian Objects: The Outer Limit The existence of a large number of icy objects orbiting beyond Neptune was firstly postulated by the Irish astronomer Kenneth Edgeworth and by the American astronomer Gerard Kuiper (1951). Their predictions were only confirmed in 1992, when the first Trans-Neptunian Object (1992 QB1) was discovered at a distance of 44 AU from the Sun. Until today, the number of known Trans-Neptunian Objects has not FCUP Mineralogy of V-type Near Earth Objects 66 stopped to increase. Trans-Neptunian Objects (hereafter TNOs), also known as Edgeworth Kuiper Belt (hereafter, EKB) objects, represent a population of small icy bodies orbiting beyond Neptune [89]. It is believed that these objects represent the most pristine and thermally unprocessed objects of the Solar System [90]. Their diameter varies from less than 1 km to more than 1000 kms. However, their sizes and shapes may be continuously being modified by collisional processes. It is quoted that the TNOs can provide a significant portion of NEAs [91-92]. TNOs reside in the so known Edgeworth Kuiper Belt (E-K Belt), as well as some of the Solar System dwarf planets: Pluto, Eris, Sedna, Quaoar, Ixion and Varuna. TNOs can be divided dynamically in several classes: Classical, Resonants, Scattered and Extended scattering disk objects (Fig.2.10) [90]. Classical objects are characterized by having orbits with low inclination and eccentricities (e≲0.2). Also its semi-major axis lies between 42 and 48 AU, so they are dynamically stable in timescales of Gyr as they approach to Neptune. Resonant objects are trapped in resonances with Neptune. They are usually found near the 3:2 mean motion resonance. The most famous resonant TNO is 134430 Pluto. Scattered objects are characterized by having a high inclination as well as high eccentricities (e 0.5, please see Fig. 2.12). Their perihelion distance is found near 35 AU and numerical simulations showed that their orbits are unstable on Gyr timescales due to perihelion interactions with Neptune. The Extended scattering disk objects are located out of interacting gravitational encounters with Neptune. These are the most intriguing TNOs objects: their orbits suggest the existence of another force or perturbation by an unknown planet (Nemesis?), by a passing star or by several resonant effects, which lifted their perihelion to the outer planetary domain. In Fig. 2.12 the Centaurus population is also represented. As it was explained above in the first section of this chapter, Centaurus are found in orbits between Jupiter and Neptune and they are probably ejected objects from the E-K belt or the MAB. As they strongly interact with the giant planets, their lifetime is short: in average 10 Myr [93]. The total mass of the TNOs can be estimated based on their observed size distribution and by assuming that their bulk composition is 1.0 g.cm-3. Hence, the resulting total mass of TNOs is around 0.30 MEarth [89]. FCUP Mineralogy of V-type Near Earth Objects 67 Fig.2.10 – Semi-major axis vs orbital eccentricity for the Outer Solar System. The dynamical subtypes of TNOs are marked on the plot. Only three dynamical resonances with Neptune are indicated. The curve NN represents the semi-major axis of the orbit of Neptune, while the lower curve found at 40 AU represents the locus of orbits with q=40 AU. Centaurus population are also represented in this plot [50]. The colors of TNOs vary from grey to different tonalities of brown or red, depending on the deposits that cover their surface. TNOs are mainly composed by mixtures of water ice and refractory particles with admixtures of frozen methane, ammonia and oxides of carbon [89]. The state of knowledge about TNOs is nowadays in a lower level respecting to the MAB. However, it is hoped that in the future more and more questions will come: “this is the main rule of the game of knowledge.” [94]. FCUP Mineralogy of V-type Near Earth Objects 68 3. Meteorites “All known meteorites are fragments of either asteroids, the Moon, or Mars, with the former dominating. They display a large diversity of texture and mineralogy. Most are ancient, dating from the first 10 million years of the solar system formation. They trace key stages of its evolution, such as the composition of the initial interstellar medium, the early stages of the Sun’s protostellar collapse, disk formation, dust condensation and coagulation, thermal processing, and planetesimal and planet formation.” in Perryman, M. (2011), The Origin of the Solar System, arXiv:1111.1286 eteorites provide us important tools to understand the evolutive processes that occurred during the infant solar system. In particular, they enable us to understand the chemical, physical and mineralogical properties of the nearearth objects (NEOs) as well as objects in the asteroid belt [95-97]. Solid bodies with extraterrestrial origin that penetrate and reach Earth’s surface are called meteorites. It is important to clearly distinguish between a fall and a find of a meteorite: meteorites are considered falls if they can be associated with an observed fall event and find if it cannot be associated with a recorded fall event. A meteorite is a fragment of an asteroid or a comet that can have different sizes and shapes. . Several perturbations in the orbit of an asteroid, due to the gravitational influence of Jupiter, can make the asteroid go through the inner regions of the Solar System and cross Earth’s orbit. As an evidence of this fact are the cases of the 4 Vesta fragments. Meteorites are named taking into account the location where they were recovered. As an example, the Allende meteorite fell in Allende (Mexico). The number of known meteorites is currently growing up through the discovery of new samples in the hot [98] and cold deserts [99]. As numerous samples are found in the same location, these meteorites are given names and numbers. The increasing number of meteorite finds, allied with the beginning of the use of charge-coupled device (CCD) detectors to obtain the visible and NIR reflectance spectra of small objects, enables us to understand the diversity of asteroid mineralogies. M FCUP Mineralogy of V-type Near Earth Objects 69 This chapter will mainly provide a description about the classification of meteorites. In addition, a final section describing the chronology of events in the early Solar System and their relation with meteorites will be presented. Keywords: meteorites; chondrites; nonchondrites; textures; primitive achondrites; differentiated achondrites; carbonaceous chondites; ordinary chondrites; enstatite chondrites; stony-irons; irons; clan; grouplet; chondrules; oxidation state; breccias; refractory inclusions; CAIs; AOAs; thermal metamorphism; HED’s; Howardites; Diogenites; Eucrites; 3.1. Classification of meteorites In the 1860s, G. Rose drew the first lines of the picture of the current classification scheme of meteorites. Rose was the first to split stones into chondrites and nonchondrites. Nowadays, based on the bulk composition and textures, meteorites can be separated into two major categories: chodrintes and nonchondrites meteorites. The latter contain the primitive achondrites and the differentiated meteorites. The main goal of a scheme classification is to divide meteorites into classes that seem to share the same origin or formation history. Indeed, it can reveal possible genetic linkages between meteorites of the same class. The chondrites and nonchondrites meteorites are classified in classes, such as: carbonaceous, ordinary, enstatite, R, achondrites, stony-iron and irons. These classes of meteorites are also divided in groups, as for example the enstatite chondrites are divided into two groups: EH and EL. For a more detailed review of the basis of this classification scheme, please see Table C6. Also for a detailed review of the meteorite groups and the number of their registered falls or finds, please see Tables C7 and C9. 3.1.1. Chondrite Meteorites Chondrites are the oldest rocks known in the Solar System, providing the best clues about the physical and chemical processes which have occurred during the infant solar nebula. The term chondrite comes from the greek word for seed, referencing the ameboid olivine aggregates (also reported as the acronym AOAs) and Ca-Al rich inclusions (also reported as the acronym CAIs), which will be explained later on. Chondrites have formed about 4.6 Myr ago and their abundances of nonvolatile elements are close to those in the solar photosphere. FCUP Mineralogy of V-type Near Earth Objects 70 Chondrites are mainly composed of four major components: chondrules, FeNimetal, refractory inclusions (which can be CAIs or AOAs) and some fine grained size material. It is commonly accepted that the refractory inclusions, chondrules as well as the FeNi-metal were formed in the solar nebula due to condensation and evaporation processes that have occurred at high temperatures. Several CAIs, FeNi-metal and chondrules were subsequently melted during some heating processes. However, some CAIs and metals in some chondrites appear to have escaped these high temperatures in the solar nebula. Fig.3.1 - Scanning electron micrographs of chondrules and CAIs. A. Porphyritic Olivine type I chondrule. B. Porphyritic Olivine Pyroxene type I chondrule. C. Porphyritic Pyroxene type I chondrule. D. Barred Olivine type I chondrule. E. Porphyritic Olivine type II chondrule. F. Amoeboid Olivine Aggregate. G. Fine-grained CAI with nodular texture shown in the insert. H. Hibonite-glass CAI spherule. I. Coarse-grained ultra-refractory igneous CAI E101,1, All objects are from the Efremovka carbonaceous chondrite, except H from the Murchison carbonaceous chondrite [107]. Chondrules are 100 µm to 10 mm spherules of once melted silicates. They are mainly composed by ferromagnesian olivine and pyroxene minerals embedded in a CaAl-rich glassy mesostatis of close-to-feldspathic compositions. Based on their textural FCUP Mineralogy of V-type Near Earth Objects 71 and chemical properties, different groups of chondrules are pointed out (Fig.3.1). There are two major groups of chondrules: type-I chondrules showing the presence of reduced iron as sulfides, magnesian olivine or pyroxene; type II chondrules containing oxidized Fe entering the structure with silicates. Chondrites are divided into five classes that share similar mineralogy, petrography, isotopic properties and chemical composition; carbonaceous chondrites, ordinary chondrites and enstatite chondrites. Some chondrites are also classified into two additional groups: the R and K groups. The R group (also known as Rumuruti-like) and K group (also known as Kakangari-like) are quite different from the other chondrites, suggesting that they represent two additional groups. The K chondrites are a small group as only two meteorites are catalogued. This group is also designated as a grouplet. Indeed, the three major classes of chondrites are the carbonaceous class (hereafter, also reported as C-class), the ordinary class (hereafter, also reported as Oclass) and the enstatite class (hereafter, also reported as E-class). These major classes also contain several groups. In the C-class there are eight groups: CI (also known as Ivuna-like), CM (also known as Mighei-like), CO (also known as Ornans-like), CR (also known as Renazzo-like), CH (also known as ALH85085-like), CB (also known as Bencubbin-like), CV (also known as Vigarano-like) and CK (also known as Karoonda-like). The letters designating the group are usually associated with a particular chondrite. The only exception is the CH group, where the “H” refers to high metal abundance and high iron concentration [101]. The O-class is divided in three groups: H, L and LL. The letters designating these groups refer the bulk content of iron: the H group possesses a high content of iron; the L group possesses a low content of iron; the LL group has a low content of metallic iron compared to the total iron bulk content, as well as a total low bulk content of iron. The E-class is divided into two groups: EH and EL. These two groups have different metallic iron bulk content. But, there is a meteorite that is classified as an enstatite which is not catalogued in one of these groups: LEW87223 [102]. Several groups of the chondritic meteorites have similar chemical, mineralogical and isotopical properties, which constitute “clans”. As an example, the H, L and LL groups are chemically and isotopically similar, so these three groups form a clan. Four other clans are also reported: EH-EL clan, CM-CO clan, CV-CK clan and CR-CH-CB clan. FCUP Mineralogy of V-type Near Earth Objects 72 3.1.1.1. Oxidation state and bulk oxygen isotopic compositions Chondrites show a wide range of oxidation states that were almost certainly due to both combined asteroidal and nebular processes [103-104]. The oxidation state is characterized taking into account the distribution of iron along with its three usual oxidation states: 0 (FeNi-metal and Fe-sulfides), 2 (silicates) and 3 (oxides). The oxidation state of the chondrites increases as follows: E-class, O-class, C-class+R. The K group is in the intermediate position between the E and O classes. In the C-class, the oxidation state increases as follows: CB-CH-CR-CO-CV-CKCM-CI. In the E-class the oxidation state increases from H to L to LL as denoted in the ratio of their olivine and pyroxene. Oxygen has three stable isotopes synthesized in stars: 16O, 17O and 18O. 16O is formed during the fusion of four He nuclei during its burning in supernovae. But, 17O and 18O are secondary isotopes which are formed from 16O and 14N, respectively. Oxygen isotope abundances in meteorites can help us to understand several physical and chemical processes occurred during the formation of our Solar System. Average values for δ17O 3 can be found in Table C8. 3.1.1.2. Petrologic type Chondritic meteorites are also classified in petrologic types [105]. This classification scheme yields a guide to the degree of thermal and aqueous alteration suffered by the chondrite. They are divided into six petrologic or petrographic types: type 1 represents the higher degree of aqueous alteration in relation to the type 2; type 3 is the less modified by secondary processes; the sequence from type 3 to 6 denotes an increasing degree of chemical equilibrium. In the C-class: CI are classified as being of petrologic type 1; CM are classified as being of petrologic type 1 or 2, however most of CM chondrites are petrologic type 2 with fully hydrated matrices and variably altered chondrules, CAIs, etc; both CO and CV are classified as petrologic types 3 or 4; both CB and CH are classified as petrologic type 3; CR are classified as being of petrologic type 1 or 2, but the majority of them are classified as type 2. All the others (meaning 3 The conventional way to express the variations in abundances of isotopes in terms of isotopic ratios, in relation to an arbitrary standard called SMOW (standard mean ocean water), is as follows: [( ⁄) ( ⁄) ] . FCUP Mineralogy of V-type Near Earth Objects 73 the CK, L, LL, H, EH, EL and R), exception made of the K group which is classified as petrologic type 3, are classified as being of petrologic types between 3 to 6. 3.1.1.3. Refractory inclusions Refractory inclusions are the oldest objects in our Solar System, with a Pb isotope age of 4.567 Gyr [106]. There are three roughly families of refractory inclusions: coarse-grained Ca-Al-rich inclusions; fine-grained Ca-Al-rich inclusions; amoeboid olivine aggregates. Calcium-aluminum rich inclusions (also reported in the literature as CAIs) are submillimiter to the centimeter sized clasts presented in chondritic meteorites (Fig.3.1). Their mineralogy is usually described as a “primary” and “secondary”. The term “primary” reports to a phase when the inclusions were formed by direct condensation melt solidification or solid-state recrystallization. The term “secondary” is used for any phase that texturally seems to be replacing another phase. Several minerals are present in the CAIs: corundum, hibonite, grossite, perovskite, melilite, spinel, Al-Tidiopside, anorthite and forsterite. The coarse-grained family of CAIs has magmatic textures and crystal sizes ranging from roughly 10 µm to 1 mm, whereas the finegrained family of CAIs consists of aggregates of 20-50 µm nodules [107]. The magmatic texture presented in the coarse-grained CAIs is indicative of crystallization from a melt process. AOAs are irregular shaped objects, which have fine grain sizes (tipically 5-20 µm) [107] that reside in a few percent of type 2-3 carbonaceous chondrites. They do not appear to show mineralogical and isotopic differences among groups. 3.2.1.4. Thermal metamorphism Thermal metamorphism affected most of the chondritic classes and plays an important role in the classification scheme developed for chondrites. It is known that thermal metamorphism modified substantially the characteristics of the chondrite meteorites and in particular it changed textural integration and recrystallization, mineral equilibrium, destruction of primary minerals, and growth of secondary minerals [108]. Several heat sources have been proposed to be the cause of thermal metamorphism, FCUP Mineralogy of V-type Near Earth Objects 80 Acapulcoites and londranites (Fig.3.6) have approximately chondritic bulk compositions, recrystallized textures, chondritic mineralogy and planetary-type rare gases. They are mainly composed of orthopyroxene, olivine, Cr-diopside, Naplagioclase, FeNi-metal, schreibersite, troilite, whitlockite, Cl-apatite, chromite and graphite [131]. Acapulcoites differ from londranites in several characteristics, such as: small grain sizes, abundant Fe, Ni-FeS as µm-sized veins and plagioclase that have escaped from partial melting. Usually, acapulcoites are fine-grained size (150-230 µm) with more or less chondritic abundances of olivine, pyroxene, plagioclase, metal and troilite. On the other hand, londranites are coarse grained (540-700 µm) and are also characterized by the depletion in troilite and plagioclase [132]. Winonaites have approximately chondritic mineralogy and chemical composition. They can contain mm-to-cm sized regions with different textures and/or mineralogies. These regions include plagioclase, diopside-rich, coarse grained olivinerich areas and/or calcic pyroxenes [133]. On the other hand, they also exhibit recrystallized textures and reduced mineral chemistries. Winonaites have experienced partial melting and show evidence for brecciation and metamorphism on the precursor body. It is known that winonaites are linked to the IAB irons (Fig.3.6) by the oxygen isotopic and mineral composition, as it was reported above about the winonaite-IAB iron silicate inclusion clan [134]. They may also be related to the IIICD irons [135]. 3.2.2.2. Differentiated Achondrites Differentiated achondrites are divided into seven subgroups, where the HED seem to be the most interesting group as there is strong evidence that they originated in the same parent body (a detailed discussion about this topic can be found in the next chapter). The others six reported groups are: angrites, aubites, brachinites, ureilites, Martian (SNC) and Lunar meteorites. Angrites are in between the oldest differentiated achondrites and are crucial references for Solar System chronology [134]. They are medium to coarse-grained (up to 2-3 mm) igneous rocks, usually associated with a basaltic composition. Their composition mainly consists of Ca-Al-Ti-rich, pyroxene, calcium-rich olivine, anorthite and plagioclase [132]. Oxygen isotopic compositions of angrites are similar to those of HED meteorites, brachinites and mesosiderites. FCUP Mineralogy of V-type Near Earth Objects 81 Aubites typically are highly reduced enstatite pyroxenites. Almost all aubites are fragmental, except in some cases that are regolith breccias as it is evidenced by their high contents of noble gases implanted by the solar wind [136]. Their compositions mainly consist of FeO-free enstatite, but it also contains plagioclase, FeO-free diopside and forsterite. Aubites have close affinities with the E-class of chondrites and are sometimes referred to as enstatite achondrites. Brachinites (Fig.3.7) are essentially made of olivine (typically 74-98%). However, others minerals can also be found: augite (4-15%), plagioclase (0-10%), some traces of orthopyroxene, chromite (0.5-2%), Fe-sulfides (3-7%), few amount of phosphate and FeNi-metal. The differences in petrology and geochemical compositions among the different brachinites propose that they possibly were not formed by the same process [132]. Fig.3.7 – Photomicrograph of a brachinite in transmitted light with partially crossed nicols: Brachina [121]. Ureilites are mainly composed by olivine, pyroxene and <10% dark interstitial material such as carbon, metal sulfides and minor fine-grained silicates [132]. Several characteristics differentiate ureilites from the other differentiated achondrites: high content of CaO in olivine and pigeonite; high Cr2O3 contents in olivine; high carbon contents; and a large range in oxygen isotopic compositions. The martian meteorites is also know by the acronym SNC, which means S for Shergottites, N for Nakhlites and C for Chassigny. They are volcanic or subvolcanic rocks. Due to their highly fractionated compositions and young crystallization ages FCUP Mineralogy of V-type Near Earth Objects 82 (possibly 180 Myr), it seem that they are derived from a planet-size body. They were linked to Mars, as several similarities between the isotopic composition of nitrogen and noble gases of the martian atmosphere and the one’s found in some shegorttites were reported. 3.2.2.3. HED's The letters from the acronym HED refers to the Howardites, Eucrites and Diogenites, respectively (Fig.3.8). Fig.3.8 – Thin section photomicrographs in transmitted light with crossed polars of typical HED meteorite types. Top Left: Howardite Kapoeta consisting of mineral fragments of highly variable grain sizes. Top Right: Unequilibrated noncumulate eucrite Pasamonte showing the typical basaltic texture of plagioclase (light) and pyroxene (colored). Down: Diogenite Johnstown illustrating the highly brecciated nature of the rock that consists essentially entirely of orthopyroxene [121]. Eucrites (Fig.3.8) are the most common type of achondrites. They have a dark fusion crust and their interior is calcium-rich presenting a brown tonality. They are named after a Greek word meaning “easily distinguished”, as they are easily distinguish from chondrites. Eucrites are also reported as the “extraterrestrial basalts” FCUP Mineralogy of V-type Near Earth Objects 83 because they are volcanic rocks of magmatic origin, probably representing the surface of their parent body (Fig.3.9). They are primarly composed by Calcium-poor pyroxenes, piogenite and anorthite. However, other accessory minerals can also be present, such as silica, ilminite, chromite, troilite or fayalitic olivine. According to their mineralogical composition they can be divided into three subclasses: noncumulate eucrites, cumulate eucrites and polymict eucrites. Noncumulate eucrites formed due to a quickly cooling of lava flows and most were then metamorphosed (Fig.3.9). Cumulate eucrites are coarse-grained gabbros and many of them are unbrecciated. They probably formed in a region of around 8km depth (Fig.3.9). Polymict eucrites are polymict breccias mainly composed by eucritic material, but they also contain less than 10vol.% of a diogenitic component in the form of orthopyroxene. Fig.3.9 – Schematic model of four hypothetical impacts into the layered crust, producing different types of Eucrites. The crustal layers from bottom to top are arranged as follows: (1) Mg-rich diogenite layer with mainly orthopyroxene; (2) cumulate eucrites with Mg-rich inverted with piogenite and calcic plagioclase; (3) ordinary eucrites; (4) surface lava flows with chemically zoned pyroxene [137]. FCUP Mineralogy of V-type Near Earth Objects 84 Diogenites were maybe formed in an interior layer of a differentiated asteroid, probably in a Fe-rich mantle (Fig.3.9). They are mainly composed by Magnesium-rich orthopyroxenes. However, other minerals such as olivine, cyclopyroxene, chromite, plagioclase, FeNi-metal, troilite and a silica phase are also present in lesser quantities [138]. Olivine is a minor phase while troilite, metal and silica are accessory phases in many diogenites. Diopside occurs as exsolution products. Plagioclase can be present in some members while phosphates are rare accessory members. Howardites are named after Edward Howard, a renowned British chemist of the 18th century and one of the pioneers in the study of meteorites. They have broken faces revealing a grey interior. Actually, they are regolith polymitic breccias mainly composed by eucritic and diogenitic debris components. In addition, they also contain abundant impact-melt clasts and breccias fragments [139]. It seems that howardites were material exposed in the surface of their parent body, as it is evidenced by the presence of sputtering gases by solar wind. 3.2.2.4. Irons Irons and stony-irons seem to be originated in the same asteroid, which have experienced high temperature events to completely differentiate. Iron meteorites are usually classified based on their chemical and structural properties. Some determinations of the most volatile siderophile elements as gallium and germanium, resolved only four groups [140], designated by Roman numerals I-IV (ordered by the decreasing content of these two elements). Letters were also added in order to distinguish additional groups. However, not all groups seem to be independent, leading to this (and sometimes confusing) nomenclature. In some cases, some groups have a combination of two letters after their Roman numeral, as for example the IAB or IIICD group. Several chemical characteristics indicate that magmatic iron meteorite groups formed by fractional crystallization, maybe in the cores of differentiated asteroids. Of the 12 groups of iron meteorites, the most enigmatic groups are the IAB and the IIICD groups. Only these two groups have broad ranges in nickel and some trace elements, which are difficult to explain by a simple fractional crystallization process or by a metallic core. For a summarized review of irons compositions and some genetic linkages between groups, please see Table C8. FCUP Mineralogy of V-type Near Earth Objects 85 3.2.2.5. Stony-irons Stony-iron meteorites are divided into two groups: mesosiderites and pallasites. Mesosiderites are breccias mainly composed of roughly proportions of silicates and FeNi-metal plus troilite. They can be divided into three petrologic classes, taking into account the orthopyroxene abundance: the content of orthopyroxene increases from A (basaltic) to C (orthopiroxinite), with an intermediate state B (more ultramafic). Pallasites are essentially composed of an amount of silicate (olivine), metal and troilite. Based on the differences of silicate mineralogy, composition, metal and isotopic composition there are three types of pallasites that can be recognized: the main group; the Eagle Station grouplet (with high Ni, Ge, and Ir and siderophiles in the metal and high content of fayalitic olivine [141]); the pyroxene-pallasite grouplet [132]. 3.3. The chronology of events in the early Solar System The chronology of the early Solar System has been extensively studied in the last decade, using the isotopic composition of several radiogenic elements in meteorites [142]. For example, several decays have been used in studying the origin and evolution of the terrestrial planets. In order to date the CAIs formation, we can point out the following decays presented in the literature: 7Be to 7Li; 10Be to 10B; 26Al to 26Mg; 41Ca to 41K; 129I to 129Xe; 244Pu to 136Xe. The latter two decays can also be used to date the chondrules chronology. In the middle 80’s, the following chronology of events in the early Solar System was proposed [143]: 4.568 Gyrs ago: Beginning of the collapse of the presolar cloud, formation of “anomalous” CAIs before incorporation of radionuclides of stellar origin; 4.567 Gyrs ago: CAI formation during a few 105 years; 4.566 Gyrs ago: Chondrule formation stars and lasts for 1-2 Myrs; 4.565 – 4.564 Gyrs ago: Accretion of chondritic planetesimals, large scale melting in the parent body of differentiated HED meteorites (postulated as being 4 Vesta, which is discussed in the next chapter). FCUP Mineralogy of V-type Near Earth Objects 86 However, during the last few years several modifications and additions were made in the latter chronology (Fig.3.10). Using bulk high precision Al-Mg dating, it was possible to establish a relative chronology for CAI and chondrules formation [144]. Also, using the Mn-Cr chronometer it was possible to date the formation of CAIs 4571 Myr ago [145]. Nowadays it is believed that CAIs formed probably during an embedded phase and during the initial phase of cloud collapse [146] and the formation of the first small planetesimals. However, an early and swift formation of planetesimals is necessary in order to prevent the rapid loss of CAIs into the Sun. Fig.3.10 – Mn-Cr and Al-Mg early Solar System timelines. Both chronometers are anchored to the time of 4571Myr, when most CAIs are believed to have been formed. Small planetesimals formed almost at the same time in order to preserve CAIs. The interiors of some intermediate-sized planetesimals were melted through the 26Al decay. The formation of the primary chondrules occurred within 2Myr. Chemical differentiation, formation of a stratified mantle and a core in these larger bodies ended 4565Myr ago (according to the case of 4 Vesta). Metamorphism occurred for many million years [145]. Al-Mg and Mn-Cr chronometers are consistent in dating the primary chondrules formation within the first 2 Myr (Fig.3.10) of the Solar System history [145]. Their formation may be concluded around 4569 Myr ago. It seems that this time was an era of intense thermal processing in the inner Solar System materials due to the decay of 26Al. Later collisional disruption distributed chondrules to their meteorite parent bodies. FCUP Mineralogy of V-type Near Earth Objects 87 Proposed mechanisms for chondrules formation include the interaction with the infant Sun, through jets or solar flares at distances less than 0.1 AU (the so-called Xwind model, which is also reported to explain the formation of CAIs) [146]. However, the most accepted theory for chondrules formation is the so-called “shock model”, in which chondrules formed in solar nebula shocks [146-147]. Several mechanisms were proposed for producing shock fronts inside the solar nebula [148] however their origin is still under debate. Identifying the mechanism responsible for the production of shock fronts can help us to understand the way that chondrules formed. Chemical differentiation, formation the core and a stratified mantle was concluded no later than 4565 Myr ago [145]. Also during this time, the H chondrite and enstatite chondrites parent bodies cooled below the temperature where isotopic closure occurs. The formation of pallasites occurred 4558 Myr ago and their younger age mark the time when either the interiors of their parent bodies cooled or the parent bodies were disrupted to allow rapid cooling of the exposed interior [145]. It is clear that accretion, accumulation of heat and chemical differentiation on larger asteroids occurred during a period of 6 Myr. Understanding the chronology of the events occurred in the early stages of the Solar System can help us to accomplish the puzzling task of understanding the steps beyond the planetary formation. But, it is important to note that the various chronometers and measurements presented in the literature are not always consistent with each other. FCUP Mineralogy of V-type Near Earth Objects 88 4. The parent body search “Asteroids investigations can provide a “map” to complement the meteorite “clock”.” in Graffey, M. J., Burbine, T. H. and Binzel, R. P. (1993), Asteroids spectroscopy: Progress and perspectives, Meteoritics, 28, 2, 161-187 s the number of known meteorites is rapidly growing up through the discovery of large concentrations of meteorites in the hot and cold deserts, this reports us to an intriguing question: “Where do meteorites come from?”. Several evidences indicate that the major source of meteorite is the inner asteroid belt. However, linking a meteorite to a specific body has been a challenging task through the past years. With the advances in the use of the charge-coupled device (CCD) detectors, it is possible to obtain visible and near-infrared reflectance spectra of smaller and smaller objects [96]. The main advantage is that it may be possible to determine parent or source bodies for particular meteorites. A parent body is the body from which the meteorite acquired its current chemical composition and mineralogical characteristics, while a source body is a fragment of the parent body by which the meteorite was completely shielded from cosmic rays for most of the Solar System’s history [96]. The postulated parent or source bodies for almost all groups of meteorites can be found in Table C9. Establishing linkages between asteroids and meteorites are of crucial importance as they allow us to understand the compositional and thermal gradients in the solar nebula. A key of interest is the very earliest history of nebular solids, which provide important constraints on nebular conditions and processes that are the ground truth for astrophysical models of nebular evolution. In the past decade, the strong relationship between the HED meteorites and the asteroid 4 Vesta was reinforced [149-150]. Linkages between ordinary chondrites and S-type asteroids, CM chondrites and C-type asteroids and irons as well as enstatite chondrites and M-type asteroids have been subject of discussion. In this chapter we will discuss the linkages that asteroids and comets may share with meteorites as well as methods to identify and study their relationships. In addition, we will describe how space weathering can play an important role in the identification A FCUP Mineralogy of V-type Near Earth Objects 89 of possible linkages between asteroids and meteorites. Finally, the model for deconvulting reflectance spectra, which will be used in this study, will be presented: the Hapke Radiative Transfer Model. Keywords: parent bodies; source bodies; 4 Vesta; Vestoids; HEDs; comets; space weathering; albedo; reflectance; absorption band; valence band; conduction band; Hapke Model; bidirectional; 4.1. The meteorite-parent body relationship Historically, it has been an arduous task linking telescopic measurements of asteroids with laboratorial measurements of meteorites. Such linkages provide physical constraints in order to develop and/or test sophisticated models of the origin and evolution of our Solar System. Long-running modeling studies of the dynamical evolution of the MAB as well as long-running spectral investigations of asteroid compositions, finally converged to yield a clear view of the asteroid-meteorite linkages. However it is interesting to note the enormous disparity between the 100 to 150 postulated parent bodies or source bodies and the approximately number of known asteroids [151]: 1000000! In the following subsections it will be pointed out a number of postulated asteroid-meteorite linkages for several groups of meteorites. As it was referred above, the postulated parent or source bodies for almost all groups of meteorites can be found in Table C9. 4.1.1. HED's-4 Vesta-Vestoids Asteroid 4 Vesta is the only known asteroid with a diameter larger than 400 km which has a basaltic crust. It is believed that the geological history of 4 Vesta occurred in the first 10 Myr of the Solar System [152]. It should be noted that basalts are rocks formed in the interior of a planet, during a stage of high temperatures and pressures and later brought to the surface. Hence, the study of basalts provides information about: the formation and differentiation of a body; the mineralogy and chemistry of its formation region in the interior of the body; thermal history and characteristics of the parent body aside from the process of the formation of magma and its evolution. FCUP Mineralogy of V-type Near Earth Objects 96 are cited as being linked to the Aor S-type asteroids. These genetic linkages between asteroids and meteorites can be found in Table C9. 4.2. Space weathering Since the discovery of the possible linkage between 4 Vesta and the HED’s meteorites, several studies aimed to link other meteoritic classes to asteroids as it was discussed in the last section. But, in some cases, the spectra of some types of asteroids do not fit well the spectra of their meteoritic analogs (Fig.4.7). This difference between some asteroid-types and their meteoritic analogs has long been cited to be caused by “space weathering” [162]. Space weathering is the amount of physical and chemical alterations which occur in the material exposed in the surface of an airless body. Nowadays it is currently accepted that space weathering can be a key modifier of visible and near infrared reflectance spectra of an airless body. It is believed that space weathering is caused by vapor recondensation after either ion sputtering or impact by micrometeorites [163]. One result of space weathering is the production of submicroscopic metallic iron (SMFe) in the regolith of airless bodies (Fig.4.6). Also, in the literature some authors use the phrase “nanophase metallic iron” (abbreviated npFe0) to refer to the carrier of the space weathering [164]. Space weathering was firstly proposed to occur on the Moon, as it was not understood why the lunar soil becomes darker and spectrally redder with time [165]. It was argued that space weathering was responsible for redden the spectral slope, obscuring absorption bands and lowering the albedo of the Moon [164]. In addition, the characteristic magnetic spin resonance of the lunar regolith is thought to be due to space weathering. The latter facts were later confirmed when astronauts of the Apollo’s mission brought back some lunar soil, which through laboratorial studies provided crucial information on the space weathering process on the Moon. However, space weathering effects on asteroids are still not well understood because we do not have soil samples of asteroids yet available for laboratorial studies. Asteroid types which show evidences of space weathering include the S, A, Q and V-types. The effects of space weathering on these types of asteroids are summarized on Table C12. It is also presented their probable meteoritic analogs and FCUP Mineralogy of V-type Near Earth Objects 97 inferred space weathering effects. In Fig.4.7, the space weathering effects are evidenced for Sand V-type asteroids [162]. Fig.4.6 - TEM image of an anorthosite (An) grain from a mature lunar soil that exhibits a rim of Fe metal particles (SMFe) [162]. Fig.4.7 – Two examples of reflectance spectra of asteroids paired with their meteoritic analogs. One possible explanation for this mismatch is the space weathering process which is thought to be responsible for the alteration of several spectral features of asteroids. Asteroids are displayed as closed triangles and meteorites as open diamonds. The top pair is offset vertically by 0.7 for clarity and show the S (IV) asteroid 7 Iris with the L6 ordinary chondrite Drake Creek. Although both objects have approximately the same mineralogy as shown by their band minima, the spectra of 7 Iris is reddened by 30% relative to Drake Creek. Also, Iris's major absorption bands are significantly reduced. The pair of 4 4 Vesta and the eucrite Jonzac also show similar mineralogy, but the band depth of 4 Vesta is reduced by 40% relative to its analog meteorite [162]. FCUP Mineralogy of V-type Near Earth Objects 98 The Q-type asteroids, for which ordinary chondrites are postulated as being their meteoritic analogs, denote on their spectra evidences of band suppression and albedo reduction. The red continuum slope and modest band suppression evidenced on the spectra of A-type asteroids are consistent of olivine weathering. In the case of V-type asteroids (Fig.4.7), the largest asteroid 4 Vesta shows a reduction in band depth evidencing the presence of space weathering, despite the band centers and shapes of the 4 Vesta and Jonzac eucrite spectra are similar. But, the strongest evidence of space weathering in asteroids is found among the S-type asteroids. Despite the spectral band center of 7 Iris (Fig.4.7) shows a similar mineralogy to the Drake Creek L6 ordinary chondrite, the weathering process seems to produce differences in band depth, red slope and albedo [162]. 4.3. Spectra Modelling: The Hapke Radiative Transfer Model In the visible and near-infrared regions, the reflectance spectra of meteorites show several absorption features, which are function of the composition and crystalline structure of the absorbing materials. Aiming to establish a possible linkage between parent body and source body, studying the reflectance spectra of meteorites can be an important tool to accomplish this hard task. However, how can we study the reflectance spectra of meteorites? In fact, this is possible through remote sensing of reflectance spectra, by inferring the surface mineralogy of asteroids and meteorites. But, to describe the absorption features it is necessary to understand the physics which controls them. Hence, the Hapke Radiative Transfer Model (hereafter also reported as the acronym HM) [166] will be described in this section. 4.3.1. Hapke Radiative Transfer Model The Hapke radiative transfer model provides a method to derive the composition and mineralogy of a surface as well as the size of its particles, through its reflectance spectra. If the surface optical constants are available, the analysis of the reflectance spectra can be made with high accuracy. The optical constants are fundamental properties of a mineral and are independent of the size and shape of the particle. The refraction index m, is defined as follows: FCUP Mineralogy of V-type Near Earth Objects 99 where n denotes the real refraction index and k denotes the complex refraction index. The spectral signatures found in absorption spectra or in reflectance spectra of mafic minerals, in the near-infrared regions are dominated by variations of k [1]. 4.3.1.1. Concepts and Definitions In this section several concepts and definitions will be introduced, such as: the difference between radiance and irradiance; the geometrical notation for the incident, emergent and phase angles of a light beam; the definitions of cross sections, efficiencies, particle scattering albedo and espat function. These definitions and concepts will be useful later to understand the basis of the Hapke Model. 4.3.1.1.1. Radiance and Irradiance In the radiation field, where the light is uncollimated, it is important to distinguish between radiance and irradiance. The amount of power at position r crossing unit area (dS) perpendicular to the direction of propagation 𝛺, travelling into unit solid angle about 𝛺 (denoted by d 𝛺 ), is called the radiance (Fig.4.8). The radiance will be denoted by . The irradiance term refers to the radiative power per unit of area of a collimated light beam (Fig.4.8). It is usually denoted by J. Fig.4.8 – Irradiance and Radiance [166]. FCUP Mineralogy of V-type Near Earth Objects 100 4.3.1.1.2. Geometric Notions Let’s assume that a light beam is focused on a semi-infinite medium. There are two crucial geometrical conventions, which are related to the latter assumption: the incident angle and the emergent angle (Fig.4.9). Let 󰇍 be the unitary vector normal to the surface ΔA, in Fig.4.9. The incident direction of the light beam can be decomposed in two angles (Fig.4.9): 1. The incident angle i: is defined as the angle between the 󰇍 and the incident direction; 2. The azimuth angle : is defined as the angle between the North direction and the projection of on the surface; After scattered by the surface, several rays of light will emerge through the surface in a certain direction , called the emergent direction. As in the case of the incident direction, can be decomposed in two angles (Fig.4.9): 1. The emergence angle e: is defined as the angle between 󰇍 and the emergent beam; 2. The azimuth angle : is defined as the angle between the North direction and the projection of on the surface; There are some aditional angles which can be defined from the ones described above, such as: 1. The phase angle: it is the angle between the incident and emergent beams; 2. The diffusion angle θ: it is the complementary angle of g (θ ); 3. The azimuth angle relative to : it is the angle between the projections of and ( | |); A relation between the phase and the azimuth angles can be derived through spherical trigonometry: FCUP Mineralogy of V-type Near Earth Objects 101 Fig.4.9 – Geometrical conventions [166]. A commonly used notation in the literature when referring to the cosine of the incident and emergence angles is the following: The plane containing 󰇍 and is called the incident plane, and that containing 󰇍 and is called the emergence plane. 4.3.1.1.3. Cross sections and efficiencies Let PE be the total power of the incident irradiance J, which is affected by the particle. The extinction cross section is defined as the ratio between PE and J: FCUP Mineralogy of V-type Near Earth Objects 102 However, a certain portion of the total power is scattered into all directions and the remaining is absorbed by the particle (Fig.4.10). Let PS be the scattered power and PA be the absorbed power. Hence: The scattering cross section, is defined as: In addition, the absorption cross section is defined as: Rewriting equation 4.6, using 4.5, 4.7 and 4.8 in leads to: Considering, now, a particle with radius a, its geometrical cross-sectional area σ, is defined as: Hence, the extinction, scattering and absorption efficiencies are, respectively defined as: Where: FCUP Mineralogy of V-type Near Earth Objects 103 Fig.4.10 – Scattering by a single particle. The plane containing J and I is called the scattering plane [166]. 4.3.1.1.4. Particle Scattering Albedo and Espat Function The particle single scattering albedo w, is defined as the ratio between the total scattered power and the total power that is removed from the wave: In addition, a parameter related to the single scattering albedo is the so-called particle espat function: In general, the efficiencies and the single-scattering albedo are functions of the wavelength. 4.3.1.1.5. Physics of absorption mechanisms When a light beam is focused on a material, where the refraction index changes, part of the light is reflected while the other is absorbed by the medium. When the photons collide with an absorber medium, they are absorbed according to the BeerLambert Law. FCUP Mineralogy of V-type Near Earth Objects 104 Let’s consider an isotropic and homogeneous medium with a thickness dl. Let I0 be the intensity measured of a light beam before crossing that medium, and I1 the intensity of the light beam measured after crossing that medium. In addition, let τ be the attenuation coefficient and c its concentration. By crossing the homogeneous and isotropic medium, the intensity of the light beam decreases by a factor of , where is called the dispersion relation. Hence: Integrating equation 4.17 for a medium of finite thickness gives: The equation 4.18 is called the Beer-Lambert law. However, what are the mechanisms responsible for the absorption? Actually, absorption in the near ultraviolet, visible and near-infrared regions of the spectrum can occur by several mechanisms, but most can be classified as transitions in which single electrons are induced by the radiation to jump from a lower state of energy to one of higher energy. These transitions can be described in terms of the band model of electrons in a solid [166]. In this model, the system has several wide continuous bands in which the electrons can exist, separated by gaps. The highest energy band, in which the electron states are all occupied, is denominated by valence band, while the lowest energy band in which not all states are occupied is called the conduction band. If an electron is excited from the valence band to the conduction band, it leaves an ion with an excess of positive charge behind it. This positively charged electron vacancy is called a hole [166]. This overview of bands occupied by electrons and holes gives rise to a rich variety of ways by which light can be absorbed. In particular, the most commonly process which is responsible for the absorption features in minerals spectra, is the electronic transition between electron shells which are not completely filled in transition elements. Usually, in an isolated ion the orbital states in the partly filled shell are degenerate. Hence, the electrons have the same probability of residing in any state. FCUP Mineralogy of V-type Near Earth Objects 105 But, if the ion is in a solid, the electric field created due to the surrounding ions is not isotropic. As a consequence, this nonisotropy removes the degeneracies and the orbitals in the shell have different energy levels. In this scenario, an electron can move from a lower level of energy to a higher level of energy when a photon is absorbed. This mechanism is usually called the crystal-field absorption [166]. The energy levels are characterized by the valence state of the atom (as for example, Fe2+ or Fe3+), by their coordination number, by the type of ligands formed, or by the metal-ligand interatomic distance. Absorption bands due to crystal-field transitions typically occur in the visible and near-infrared regions. In particular, the Fe2+ bands near 1 and 2 μm and the Fe3+ band near 0.86 μm have important applications in geochemical remote sensing. As an ion is able to produce different absorptions, it is possible to identify specific minerals in spectroscopy. 4.3.1.2. Radiative Transfer Equation The formalism which is usually used to calculate how the intensity of an electromagnetic wave changes due to emission, absorption and scattering processes, as the wave crosses a complex medium is known as the radiative transfer equation. The main hypothesis of this formalism is to suppose that the medium heterogeneities emit and scatter the radiation independently [166]. The radiative transfer equation is written as: 𝛺 𝛺 ∫ 𝛺 𝛺 𝛺 𝛺 𝛺 𝛺 ⁄ 𝛺 4 With several precautions, the theory of the radiative transfer can be applied to two mediums of particular importance in remote sensing: atmospheres and planetary regoliths. However, if the particles do not have a regular shape and if they are not uniformly spaced, the application of this theory can give erroneous results [166]. 4 A deduction of this equation can be found in Hapke (1993) [164], pages 151-155. FCUP Mineralogy of V-type Near Earth Objects 112 where is the intensity of the scattered light. But, as the bidirectional reflectance is defined as the ratio between the intensity of the scattered light and the irradiance J: If , from equation 4.33: Equation 4.34 is known as the Lommel-Seelinger Law. Let us consider now a simple surface composed by particles that scatter light independently and isotropically [166]. The exact solution of this medium was derived by Ambartsumian (1958), based on the fact that adding a new thin layer to this surface will not change the reflectance. The reflectance of this medium is: where is the Ambartsumian-Chandrasekhar H-function, which satisfies the following integral equation: ∫ The multiple scattering reflectance is obtained by subtracting the total reflectance (equation 4.35) to the Lommel-Seelinger Law (equation 4.34): [ ] Hapke (1993) [166] derived an approximate solution for equation 4.36, based on simple assumptions in solving the transfer radiative equation: FCUP Mineralogy of V-type Near Earth Objects 113 where √ . However, Hapke (1993) [166] also derived an accurate approximation for the Ambartsumian-Chandrasekhar function. Equation 4.36 can be rewritten in the following way: [ ∫ ] Solving equation 4.39, Hapke (1993) [166] obtained the following approximation for equation 4.36: [ [ ] [ ( ) ( )]] where √ . More recently, Hapke (2002) [169] presented a better accurate version of the approximate solution for equation 4.36. Linearizing equation 4.36, Hapke (2002) [169] obtained that: [ ( ( ))] where √ √ , where s represents the internal scattering coefficient of the particle near its surface. The bidirectional reflectance of Hapke is obtained by the sum of equations 4.34 and 4.37: [ ] Equation 4.45 is a good approximation for dark surfaces or medium which scatter the light isotropically. However, in laboratorial samples, the reflectance denotes a non linear growth for small values of g, near the opposition. The non linear peak that is evidenced is FCUP Mineralogy of V-type Near Earth Objects 114 called the opposition effect. This phenomenon is only found in dark surfaces which are porous and particulate and where the mutual blocks between the particles cause a shadow that is superior to the wavelength of the incident light beam: Shadow-hiding opposition effect (SHOE). This opposition effect was studied by Hapke (1993) [166]. As only the single scattering is affected, Hapke (1993) [166] introduced a new term in equation 4.34: [ ] where is the term due to the opposition effect and is defined as follows: ( ) Finally, the expression of the Hapke bidirectional reflectance is obtained by replacing equations 4.46 and 4.37 in equation 4.45: [[ ] ] 4.3.1.5. The equivalent slab approximation for QS The internal transmission coefficient, Θ, lists the fraction of light that enters into a particle and the fraction of light that reaches the opposite surface, after transmitted. The internal transmission coefficient is defined as follows: { ( ( ⁄) ⁄( ( ⁄))) [ ]} 6 where represents the diameter of the particle. 6 A deduction of this expression can be found on Hapke (1993) [164], page 98. FCUP Mineralogy of V-type Near Earth Objects 115 Let be the mean path length travelled by the light beam. It represents the thickness of a slab, that will have the same value of Θ, when . For the slab: . Expanding equation 4.49 in powers of : ( ⁄) where is the mean distance travelled by all rays of light during a single transit of the particle. If the internal transmission coefficient of a particle near to its surface is equal to zero : (exponential model). If , the scattering occurs between the particles or close to their surfaces. In this case, the internal transmission coefficient is given by: ( √ ) ( √ ) In the case of planetary regolith particles, where , the extinction efficiency is: . Hence, the albedo w, can be defined as follows: According to the equivalent slab model, a good approximation for , when the incident light is unpolarized, consists in replacing the spherical particle by a slab, with adequate optical constants (Fig.4.14). Let and be defined as the external and internal reflectance Fresnel coefficients of a surface, respectively. According to Hapke (1993) [166], an empiric approximation for , can be written as: FCUP Mineralogy of V-type Near Earth Objects 116 The latter equation (4.54) gives good results (Fig.4.15) for and ≲ ≲ . On the other hand, the internal reflectance of a surface can be approximated by: The latter equation (4.55) also gives good results (Fig.4.15) for and . Fig.4.14 – Scheme of the equivalent slab model for [166]. In Fig.4.14 there are two quantities that have not been defined yet: SE and SI. Returning to Fig.4.14, during the first passage through the particle, a fraction of light SI is internally reflected and the remaining 1-SI is refracted through the surface. This process continues as it is evidenced in Fig.4.14. Due to spherical symmetry of the problem, the reflection coefficients are the same for each order of internal reflection. FCUP Mineralogy of V-type Near Earth Objects 117 Fig.4.15 – External (SE) and Internal (SI) surface reflectance coefficients versus the refractive index for . The solid lines denote the exact expressions while the dashed lines denote the approximations results [166]. Hence, the total fraction of light emerging from the surface is: [ ] Due to the spherical symmetry of the problem, . So, from equation 4.56: From equation 4.57, the absorption efficiency can also be obtained: FCUP Mineralogy of V-type Near Earth Objects 118 4.4. Computation of the Hapke Model and justification of the parameters used The Hapke Model is usually cited as a two step model. First, the reflectance spectra are converted into single scattering albedo, where the effects of geometry are removed. Second, the single scattering albedo is converted into optical constants, where the grain size of the sample particles is removed. However, for the purpose of this study, the Hapke Model will be only used to derive the single scattering albedo of the meteorite samples from their reflectance spectra. Hence, only the first step of the Hapke Model will be computed. The radiance coefficient is derived through the bidirectional reflectance expression of Hapke (1993) [166]: [[ ] ] The bidirectional reflectance of Hapke (1993) [166] is derived as function of six parameters: the albedo w, the cosines of the incident and emergence angles ( and , respectively), the opposition effect , the particle phase function and the Ambartsumian-Chandrasekhar functions . The opposition effect can be set to 0 for phase angles greater than 15º. In this study, the phase angles are 30º and so the opposition effect can be neglected [170]. The particle phase function used is the one given by the second order Legendre polynomials: where and , which are the appropriate average values for forward scaterring minerals [170]. is the Ambartsumian-Chandrasekhar function. In this work it will be used the approximation derived by Hapke (1993) [166]: Eq.4.41.This approximation is an accurate approximation derived by Hapke (1993) [166] that was later confirmed to give results close to the exact solutions one’s [171]. FCUP Mineralogy of V-type Near Earth Objects 119 The first step of the Hapke Model described above was implemented using Microsoft Excel [1]. It was shown that the implementation of this model in Microsoft Excel also gives accurate results, when compared to other software implementations, such as Matlab or IDL [1]. In addition, the implementation of this model in Excel can be accessible to all people that have Microsoft Office installed in their computers. The first step of this compositional model was coded using Visual Basic for Applications (VBA), available on Microsoft Excel software [1]. Firstly, the w value is derived through an optimization process to find the radiance coefficient which matches the spectral measurements. This optimization process is done through a table of values of rc and w. For accuracy, rc is calculated for 1000 values of w, varying stepwise from 0 to 1. A linear interpolation between adjacent values is made in order the increase the accuracy. The VLOOKUP() function is used to find the nearest values to the spectral measurements and the corresponding single scattering albedo is output. This process is repeated for each wavelength [1]. In Fig.4.17 a general view of the input spreadsheet of the Hapke Model implementation is presented. Reflectance spectrum of a sample is input into the spreadsheet as two columns of data (columns highlighted in orange), one containing wavelength and the other containing the corresponding reflectance measurement. The ID of the sample is input into the spreadsheet into cells B1, B2 and B3 (cells highlighted in red). Model variables’ are grouped in the green box: green colored cells represent sample specific values; orange colored cells represent assumptions of the Hapke Model. The buttons, displayed in Fig.4.17, are used to run the Hapke Model by steps, which are coded using VBA. The button called “Chart 1: Reflectance vs Wavelength” is used to obtain the plot of the input data. The button called “Clear Chart 1” is used to clear the created plot of the input data. The button called “Solving for Albedo” is used to solve the first step of the Hapke Model, which calculates the single scattering albedo from the input reflectance data, for each wavelength, through the optimization process described above. The button called “Clear Albedo Calculations” is used to clear all the calculations done in the first step of the model. The button called “Chart2: Albedo vs Wavelength” is used to obtain the plot of the derived single scattering albedo versus wavelength. The button “Clear Chart 2” is used to clear the single scattering albedo versus wavelength plot. FCUP Mineralogy of V-type Near Earth Objects 120 Fig.4.17 – Input parameters for the Hapke Model spreadsheet. Red colored cells represent the identification of the sample, orange colored cells represent the input spectra data, green colored cells represent sample specific values and orange colored cells represent assumptions of the Hapke Model. 4.5. Compositional Determination using the Solver tool in Microsoft Excel The mineralogical composition of a sample can be derived by fitting its single scattering albedo curve through the linear combination of reasonable end-member minerals and minimizing the sum of the residuals’s squares χ2 (i.e. by minimizing the difference between the single scattering albedo’ and the modelled curves through the method of least squares). But, it is important to note that any mixing model is strongly dependent on the choice of end-members. If end-members are not constrained, then derived model abundances must be viewed as an estimate of possible composition and not as a unique determination. This process was also implemented using Microsoft Excel software (authors: Teresa M. Seixas and M. A. Salgueiro da Silva, Physics and Astronomy Department, Faculty of Sciences, University of Porto, 2010). A resulting demo spreadsheet is shown in Fig.4.18. FCUP Mineralogy of V-type Near Earth Objects 121 Fig.4.18 – General overview of the demo spreadsheet developed in Microsoft Excel, to derive the mineralogical composition of laboratory mixtures. In Fig.4.18 the highlighted cells in orange represent the identification of the samples. The single scattering albedo for the sample and end-members must be placed in the columns which are highlighted in yellow and purple, respectively. The range of wavelength must be placed in the column highlighted in cyan. Highlighted in grey there are two initial estimates of the end-members proportions presented in the sample. The sum of the residuals is highlighted in green. In column H, the modelled single scattering albedo is derived. In column I, the residuals are derived. In column J, the square of the residuals is calculated. The derivations presented in cells highlighted in grey and columns H, I and J are changed each time the Solver is applied until the minimal solution (i.e. the lowest of the residuals’ sum) is found. FCUP Mineralogy of V-type Near Earth Objects 128 Fig.5.7 – Plot of the residuals of the sample XP-CMP-010 (C1XP10). All single scattering albedo of the laboratory mixtures with a grain size proportion between 0 and 45μm were well fitted. The only exception is the case of the sample XP-CMP-014, as constrained by the residuals in Fig.E.4. For these seven laboratory mixtures, their nominal and derived compositions are plotted in Fig.5.8. In addition, the problem stated above about the derived composition for the sample XPCMP-014, seems to be supported by Fig.5.8. It is evidenced a little discrepancy between the nominal and derived compositions for this sample. In the case of the laboratory mixtures with a grain size proportion between 45 and 75μm, only the single scattering albedo’ of the samples XT-CMP-033 (Fig.E13), XT-CMP-035 (Fig.E17) and XT-CMP-038 (Fig.E23) seem to be well fitted as it constrained by their residuals. In all the other cases, the absorption feature found near 2 μm is not perfectly fitted, as evidenced by the samples residuals’. However, the absorption feature found near 1 μm seems to be well fitted in all the samples. For these seven laboratory mixtures, their nominal and derived compositions are plotted in FigS.5.95.11. FCUP Mineralogy of V-type Near Earth Objects 129 Fig.5.8 – Plot of clinopyroxene (Cpx) versus orthopyroxene (Opx) nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 0 and 45μm. Fig.5.9 - Plot of orthopyroxene (Opx) versus olivine nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. FCUP Mineralogy of V-type Near Earth Objects 130 Fig.5.10 - Plot of anorthosite versus olivine nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. Fig.5.11 - Plot of anorthosite versus orthopyroxene (Opx) nominal and derived compositions for all seven laboratory mixtures with a grain size proportion between 45 and 75μm. FCUP Mineralogy of V-type Near Earth Objects 131 According to Fig.5.9, the derived mineralogical compositions, for orthopyroxene and olivine, seem to be close to the nominal compositions. However, only the case of the sample XT-CMP-034 shows a little discrepancy between the derived and nominal compositions. Actually, the single scattering albedo of this sample is not well fitted, as it is also constrained by its residuals’ plot (Fig.E16). In addition, this discrepancy is also evidenced in Fig.5.10 and Fig.5.11. In all the others samples, the derived mineralogical compositions, for anorthosite and olivine (Fig.5.10) and for anorthosite and orthopyroxene (Fig.5.11), seem to be close to the nominal compositions. In all laboratory mixtures with a grain size proportion between 70 and 145 μm the single scattering albedo seems to be well fitted, as evidenced by their residuals. However, in some samples, the residuals show some variance near the 1 μm absorption band. This can be due to the effects of the grain size proportion: in fact, the grain size proportion of the chosen end-members is 70 to 125 μm, inferior to the ones of the selected samples. Actually, it is also important to quote the effect of the grain size proportion of the end-members in the derivation of the mineralogical composition of the mixtures. The grain size proportion of the end-members may be similar to the ones of the selected laboratory mixtures. However, in this case none end-members with a grain size proportion between 70 and 145 μm were found in the RELAB database. Only end-members with close values of the grain size proportion between 70 and 125 μm were found. In Fig.5.12 and 5.13 the effect of the grain size proportions are evidenced. In Fig.5.12, the reflectance spectra of all three clinopyroxenes, selected from the RELAB database, are plotted. It can be seen that as the grain size proportion increases, the reflectance spectra decreases. In Fig.5.13, the single scattering albedo versus wavelength is plotted, for all three clinopyroxenes. It can be also seen that as the grain size proportion increases, the derived albedo decreases. Hence, if the grain size proportions are not taken into account the derived mineralogical composition may be completely meaningful. FCUP Mineralogy of V-type Near Earth Objects 132 Fig.5.12 – Reflectance spectra of all three clinopyroxenes obtained from the RELAB database. Fig.5.13 – Single scattering albedo versus wavelength for all three clinopyroxenes. FCUP Mineralogy of V-type Near Earth Objects 133 For all these seven laboratory mixtures, their nominal and derived mineralogical compositions are plotted in Fig.5.14. As it is evidenced by Fig.5.14, the derived mineralogical compositions for the orthopyroxene seem to be underestimated in some cases. Fig.5.14 – Plot of clinopyroxene (Cpx) versus orthopyroxene (Opx) nominal and derived compositions for all fourteen laboratory mixtures with a grain size proportion between 0 and 45μm and 70 and 145μm. As supported by the results of this test, the applied procedure to derive the mineralogical composition of samples seems to give good results. Hence, this procedure fulfilled the confidence tests. 5.2. Mineralogy of HED meteorites and of V-type Asteroids In this section it is aimed to analyse spectroscopic features of HED meteorites and compare with the V-type asteroids. To obtain information about the mineralogical composition of asteroids, studying their reflectance spectra in the visible and NIR is a powerful technique. In this spectral range (0-4-2.6 μm), there are several absorption features which can be easily identified to infer about the presence of several minerals in the surface of an asteroid. Through a medium resolution spectrograph, the FCUP Mineralogy of V-type Near Earth Objects 134 reflectance spectrum of an asteroid can be obtained. As the obtained spectrum is due to the Sun’s reflected light from the asteroid, it must be corrected by extracting the Sun light component. The resulting spectrum only has information about the asteroid’s surface. For the purpose of this study, a set of five V-type asteroids were obtained from “The MIT-UH-IRTF Joint Campaign for NEO Spectral Reconnaissance” [175], available online at http://smass.mit.edu/minus.html. It is an ongoing joint observing program for routine measurement of Near Earth Objects (NEO) spectra, conducted by MIT, University of Hawai and the NASA Infrared Telescope Facility (IRTF). A low-to-medium resolution NIR spectrograph and imager called SpeX, is used to obtain 0.8 to 2.5 microns spectra of NEO. When available, visible wavelength data from SMASS survey are also included. Normalization is at 0.55 μm, when visible wavelength data are available. If not, normalization is made near 1.21 μm. The selected set of asteroid is listed in Table D3. When needed, the geometric albedo values listed in Table D3 were used to “unnormalize” the asteroid spectra. The reflectance spectra of the selected five V-type asteroids are presented in Fig.5.15. In addition, a set of ten Howardites, ten Eucrites and ten Diogenites were obtained from the RELAB database [173]. Their reflectance spectra range in wavelength between 0.3 and 2.6 μm, and were obtained with a 0.005 μm sampling resolution. The thirty selected HED meteorites are listed in Table D4. The reflectance spectra of the selected HED meteorites are presented in Fig.5.16. All selected asteroids (Fig.5.15) have the pyroxene and possibly olivine and plagioclase characteristics absorption features of the HED meteorites (Fig.5.16). The pyroxene features are due to the presence of Fe2+ ions in the M2 crystallographic site. All selected V-type asteroids (Fig.5.15) and HED meteorites (Fig.5.16) have strong pyroxene bands centered near 0.9 μm and 1.9 μm. All of the four V-type near-Earth Vestoids (meaning 3908 Nyx, 4055 Magellan, (5604) 1992 FE and (6611) 1993 VW) have strong absorption edges. 3908 Nyx, 4055 Magellan and (5604) 1992 FE have peaks near 1.3 μm, while (6611) 1993 VW have a peak near 1.4 μm. The 1.2 μm absorption band, attributed to Fe2+ in the M1 crystallographic site, is not visually distinctive in the selected V-type asteroids. In addition, the absorption feature found near 0.9 μm, can also be indicative of the presence of olivine. The olivine spectrum is FCUP Mineralogy of V-type Near Earth Objects 135 dominated by a complex absorption feature centered near 1.0 μm. For further details about the crystallographic sites and spectroscopic features of pyroxene, olivine and plagioclase, check section B.3. Fig.5.15 - Reflectance spectra of 4 Vesta and the four near-Earth vestoids. The reflectance spectra are incrementally shifted vertically by 0.5. FCUP Mineralogy of V-type Near Earth Objects 136 Fig.5.16 - Reflectance spectra of the selected ten diogenites (D), ten eucrites (E) and ten howardites (H). Eucrites’reflectance spectra are shifted vertically by 1 unity. Howardites’reflectance spectra are shifted vertically by 2 units. There are some interesting techniques to analyze the spectroscopic features in the reflectance spectra of asteroids and compare them with meteorites. One interesting technique is to compare meteorites spectra with asteroids spectra. The first step is to identify the mineralogy of an asteroid, by comparing their reflectance spectra with ones’ of meteorites taken from reflectance spectra databases. Hence, the reflectance spectra of HED meteorites were compared with the ones’ of the selected V-type asteroids (Figs.E41E45). FCUP Mineralogy of V-type Near Earth Objects 137 The mineralogies of the five selected V-type asteroids are quite similar to the ones of the HED meteorites (Fig.5.17). As it is quoted in the literature, 4 Vesta has a mineralogical composition similar to howardites [176]. This fact is constrained by Fig.5.17. In addition, the mineralogical compositions of the near Earth vestoids (6611) 1993 VW and (5604) 1992 FE were identified as being close the ones’ of howardites. 4055 Magallen seems to have a similar mineralogy to eucrites. Finally, 3908 Nyx were identified to be similar to the ones’ of diogenites. This fact seems to contradict the lack of identified vestoids with mineralogies similar to diogenites in the NEA [177]. Also, this method provided us possible particle sizes on the selected asteroids’ surfaces, as they are currently unknown. Fig.5.17 - Reflectance spectra of the five V-type asteroids versus respectively HED with the closest mineralogy. A vertical shift of 1 unit is introduced between different groups of spectra. FCUP Mineralogy of V-type Near Earth Objects 144 where, (5.3) Parameters Ak (k=0,1,2,….,n) and the compositional mineral proportions are fitted in the Solver application. In this study, a third degree (n=3) polynomial will be considered. In the modeling process, it must be guaranteed that the sum of the endmembers minerals compositions (xi) is 1. In order to better fit the individual selected Vtype asteroids and HED meteorites, the piogenite and enstatite end-members with different granulometries, were considered in the modeling process. The derived compositions for the selected Diogenites, Eucrites and Howardites are listed in Tables D6, D7 and D8, respectively. In addition, the derived compositions for the selected Vtype asteroids and their meteoritic analogs are listed in Table D9. In this study, the term “meteoritic analog” refers to a certain meteorite which is inferred, from Fig.5.17, to have the closest mineralogy to a certain V-type asteroid. In all cases, the results of the lowest χ2 found, is reported. In Fig.E46 the best obtained fit of diogenites, meaning the fit of the sample MP-TXH-068-A, and its residuals plot are presented. In Fig.E47 the worst obtained fit of diogenites, meaning the fit of the sample MP-TXH-071-A, and its residuals plot are presented. In Fig.E48 the best obtained fit of eucrites, meaning the fit of the sample MB-TXH-069-B, and its residuals plot are presented. In Fig.E49 the worst obtained fit of diogenites, meaning the fit of the sample MP-TXH-070-D, and its residuals plot are presented. In Fig.E50 the best obtained fit of diogenites, meaning the fit of the sample MP-TXH-053-A, and its residuals plot are presented. In Fig.E51 the worst obtained fit of diogenites, meaning the fit of the sample MP-TXH-068-D, and its residuals plot are presented. In Fig.5.23, it is shown the best fit results of the mixing model simulations of the five selected V-type asteroids. In addition, in Fig.E52 to E56 their obtained fits and respectively residuals plot and background curve are presented. It can be seen from Fig.5.23, that these results are good matches for all of the selected asteroids. In Fig.5.24, the reflectance spectrum of 4 Vesta and the reflectance spectra of the chosen appropriate end-members are presented. The appropriate end-members minerals chosen to fit the reflectance spectrum of 4 Vesta are listed in Table D9. FCUP Mineralogy of V-type Near Earth Objects 145 Fig.5.23 – Best fit model results of the selected V-type asteroids. Data points represent the single scattering albedo of the samples, while the black lines represent the best obtained fits. A vertical shift of 0.5 units is introduced between different asteroids/meteorites albedos. On the other hand, if the particle grain sizes (di) and each mineral density (ρi) are known, the respectively HED meteorites and V-type asteroids mass fractions can be derived through the following expression: ∑ The derived mass fractions of selected Diogenites, Eucrites and Howardites are listed in Tables D10, D11 and D12, respectively. In addition, the derived mass FCUP Mineralogy of V-type Near Earth Objects 146 fractions for the selected V-type asteroids and their meteoritic analogs are listed in Table D13. Fig.5.24 – Plot of the reflectance spectrum of 4 Vesta with its appropriate end-members minerals found in this study. All spectra are normalized by the reflectance value at 1.2μm. In Fig.5.25 is presented a plot of the derived high-Ca pyroxenes vs derived lowCa pyroxenes in this study, for all the selected HED meteorites and V-type asteroids. The derived mass fractions indicate that the surface mineralogies of the studied V-type asteroids are indeed composed by a mixture of pyroxenes (Fig.5.25). In addition, the mineralogies of the selected HED meteorites are composed by a mixture of pyroxenes and in some cases some traces olivine are also present. It is also important to quote that the derived mass fractions of the selected HED meteorites and V-type asteroids may not be unique, as mixing models are highly on the choice of endmembers. FCUP Mineralogy of V-type Near Earth Objects 147 Fig.5.25 - Plot of the low-Ca pyroxenes content vs high-Ca pyroxenes content, of all thirty HED meteorites and all five V-type asteroids. The derived mass fractions of diogenites indicate that their mineralogies are mainly dominated by magnesian orthopyroxenes, consistent with the mineralogical composition found in the literature for this type of meteorites. Chromite is a minor phase being found in only three samples: MP-TXH-081-A, MP-TXH-088-A and MBTXH-095-B. Troilite is an accessory phase, which was found in only one sample: MBTXH-074-B. Anorthite was also found in one sample: MB-TXH-067-D. It is also important to quote that the added low-degree polynomial to the modelling process, can model the presence of chromite, troilite or silica, as these minerals do not have any absorption feature as function of the wavelength. FCUP Mineralogy of V-type Near Earth Objects 148 The derived mass fractions of eucrites indicate that their mineralogies are mainly dominated by piogenite and enstatite. Minor phases of chromite were found in two samples: MB-TXH-069-B and MB-TXH-096-D. Accessory phases such as troilite and phosphates were also found. Troilite were found in three samples: MB-TXH-097-A, MB-TXH-096-D and MB-TXH-070-D. Phosphates were found in only one sample: MPTXH-084-A. These mineralogies are in agreement with the literature, however fayalitic olivine and Ca-rich plagioclase (anorthite) were not found in these samples, according to the results of this study. In addition, as it was referred above in the case of diogenites, the added low-degree polynomial can model the presence of minerals such as troilite or chromite. The derived mass fractions of howardites indicate that their mineralogies are mainly dominated by diogenitic and eucritic material (Fig.5.25), being consistent to the mineralogies found in the literature. These results strengths the relationship between HED meteorites and V-type Near Earth Asteroids. Indeed, they strength the relationship between 4 Vesta, HED meteorites and Vestoids (Fig.5.25). As pointed out in the literature and confirmed in this study, 4 Vesta has a reflectance spectrum that can be matched very well with the spectrum of a howardite [176]. The derived mineralogies of 4 Vesta and of howadirte MP-TXH-083-A also tend to strength this idea. As constrained by the derived mineralogies of 4055 Magellan, (5604) 1992 FE and (6611) 1993 VW and their postulated HED meteoritic analogs in this study, all of them have pyroxene mineralogies consistent with eucrites or howardites (Fig.5.25) [177], supporting the links established in Fig.5.17. Finally, the only near-Earth vestoid with pyroxene mineralogy consistent with diogenites is 3908 Nyx (Fig.5.25). This evidence contradicts the fact that none near Earth vestoids have pyroxene mineralogies consistent with diogenites [177]. In addition, as it can be seen from Fig.5.17, the reflectance spectrum of 3908 Nyx can be matched by the spectrum of a diogenite (MP-TXH-071-A). As none of the studied near-Earth Vestoids are composed only by one type of pyroxenes, this is suggestive of the absence of a quickly cooling of their surfaces. The wide range of the studied V-type Near Earth Asteroids mineralogies points towards one direction: they are not originated from different parent bodies, but they FCUP Mineralogy of V-type Near Earth Objects 149 almost certainly come from different layers of 4 Vesta, which was excavated by a huge collision. In fact, 4 Vesta has a huge crater in the South Polar Region, probably originated by a big collision. This collision eventually ejected a big amount of material into 4 Vesta neighborhood. Some ejected fragments were re-accumulated and formed small bodies, which show a howardite type spectra (i.e composed by eucritic and diogenitic material). These bodies would have formed the howardite type family of asteroids. If this collision excavated enough material until it reached inner layers close to the mantle of 4 Vesta, traces of olivine may be present in some Vestoids and HED meteorites. If true, this evidence may favor the ocean magma model of 4 Vesta interior [178]. Actually, in the near-Earth Vestoid (6611) 1993 VW, a small amount of fayalitic olivine was found. If the latter assumption is true, then (6611) 1993 VW can be originated from the inner layers close to the mantle of 4 Vesta. Despite being suggested in the literature that Vestoids may come from different layers of 4 Vesta, the near-Earth Asteroids 3908 Nyx and 4055 Magellan can be ejected fragments of the disruption of a 4 Vesta clone and not 4 Vesta itself [179]. This can be a good hypothesis, but such a clone body is still unknown. These two nearEarth Vestoids are also referred as being dynamically too young to have been originated in the same event as the Vesta family did [180]. However, the derived mineralogies of the studied near-Earth Vestoids do not point towards this direction. The silicate mineralogy inferred for the studied V-type Near Earth Asteroids requires that their parent body or parent bodies suffered early igneous differentiation in their histories [181], producing a basaltic surface lithology. In fact, the presence of highCa pyroxenes in asteroids can act as a tracer of their igneous history. Several minerals, such as troilite, high-Ca orthopyroxenes and plagioclase were all preferentially incorporated in early partial melts from a chondritic precursor and depleted from residues of partial melting [181]. Hence, the parent body of HED meteorites and Vestoids, which is reported as being 4 Vesta, may be a differentiated asteroid. FCUP Mineralogy of V-type Near Earth Objects 150 6. Conclusion In this work, a set of five V-type asteroids and thirty HED meteorites reflectance spectra were obtained from public databases. These spectra were modelled to test the hypothesis that HED meteorites may be linked to Vestoids and 4 Vesta. Their spectra were modelled through the application of a mixing model, based on the Hapke theory, with reasonable end-members pointed out in the literature. All end-member spectra were obtained from the RELAB public database and were constrained to samples with grain sizes less than 125 μm. The Hapke Model was coded using Visual Basic for Applications, available in Microsoft Excel software, to derive the single scattering albedo of the selected samples. Then, a Solver application was developed to model the derived single scattering albedo of the chosen samples, using reasonable endmembers choices, by minimizing the sum of squares of the residuals. Firstly, to gain confidence, this procedure was applied to a set of laboratory mixtures of minerals, containing low-Ca pyroxenes, high-Ca pyroxenes and olivine. Their reflectance spectra were obtained from the RELAB database. With the exception of few cases, the derived compositions almost matched the nominal compositions. Then, fulfilled the confidence test, this procedure was applied to the selected samples of V-type asteroids and HED meteorites. The results of this study strengthened the hypothesis of a genetic linkage between V-type asteroids and HED meteorites. In addition, they indicate that the studied near-Earth Vestoids may come from different layers of 4 Vesta or another, still unknown, 4 Vesta-like body. According to the detected presence of fayalitic olivine in (6611) 1993 VW, this near-Earth Vestoid can come from the inner layers close to the mantle of 4 Vesta. The presence of olivine found in (6611) 1993 VW and some HED meteorites, may favor the ocean magma computational model of 4 Vesta interior. These near-Earth Vestoids were probably originated from a huge collision suffered by 4 Vesta or a 4 Vesta-like body. In this study, it was possible to identify a near-Earth Vestoids with pyroxene mineralogy consistent with diogenites: 3908 Nyx. This evidence confirmed that it was possible to link a near-Earth Vestoids with diogenites. However, the apparent difficulty in linking diogenites with V-type asteroids could not be confirmed or refuted by the FCUP Mineralogy of V-type Near Earth Objects 151 results of this study. The remaining studied near-Earth Vestoids were identified as having pyroxene mineralogies consistent with eucrites and howardites, confirming what is pointed out in literature [177]. On the other hand, the pyroxene mineralogy of 4 Vesta was confirmed as being consistent with howardites. The derived silicate mineralogy of the studied near-Earth Vestoids indicates that their parent body (probably 4 Vesta) may be a differentiated asteroid [181]. It is hoped in the future to apply this procedure to study other types of asteroids and meteorites or even comets. These studies can give us clues about the processes that took place in the infant Solar System. So far, the mineralogical composition and internal structure of the primitive Earth is still unknown. Understanding several processes that took place in the infant Solar System can teach us crucial steps of its formation history. Due to the recently increasing number of discovered exoplanets, the question of how do planetary systems formed comes to our mind. It is interesting to note that to study these planets, we still need to study their host star. However, due to the advances in technology, in a close future, several Earth-like exoplanets will be probably found. By studying their mineralogies, this will probably leave us one step closer to understand the formation of Earth-like planets. In the zoo of known exoplanets, by understanding the formation of their planetary systems, Humankind will be one step closer to understand the nature and origin of life. 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Abstracts of Poster Communications in Conferences F.1.1. Abstract of the Poster Communication in the IJUP11 Conference (17-19 February 2011) Laboratorial characterization of meteorites S. Batista 1 and T. M Seixas 1 1 Department of Physics and Astronomy, Faculty of Science, University of Porto, Portugal. Asteroids represent the sole surviving population of early inner solar-system planetesimals [1-2]. They have not formed a planet due to Jupiter’s gravity influence and fragmentation produced by several collisions with other asteroids. Nowadays, asteroids provide our only in situ record of the conditions and processes that the inner portions (~1.8-3.5 AU) of the late solar nebulae and the infant solar system have experienced [1]. Several advances in our knowledge of the diversity of asteroid mineralogies have been made since the beginning of the use of charge-coupled device (CCD) detectors to obtain the visible and nIF reflectance spectra of small objects [3]. The bulk composition, mineralogy and petrology of a meteorite are functions of the original bulk composition of its parent body and the conditions of heating and melting that it has experienced during its formation. Our aim in this work is to study several meteorites of different types for investigating meteorite processes and the associated thermophysical history. In a first step, thin sections of the meteorites will be prepared for optical mineralogy studies and electron microprobe analysis. These thin sections will be taken from a layer below the exposed fusion crust. In a second step, powder samples are to be prepared by milling slices of the original meteorites for X-ray diffractometry. In a third step, thin sections previously prepared will be analysed by optical mineralogy using a petrographic microscope, allowing the identification of the component minerals and their microtexture. Electron microprobe analysis in the same samples will help to complement this information by providing chemical composition profiles along selected paths such as grains boundaries, shearing planes and chondrules internal structures. Finally, X-ray diffractometry on powder samples will be performed in order to determine the crystal structure of the constituent phases and to discriminate eventual polymorphic phases. FCUP Mineralogy of V-type Near Earth Objects 257 References: [1] Gaffey, M.J., Cloutis, E.A., Kelley, M.S., Reed, K.L. (2002), Mineralogy of asteroids. In: Bottke, W.F., Cellino, A., Paolicchi, P., Binzel, R.P. (Eds.), Asteroids III. University of Arizona Press, Tucson, pp. 183–204. [2] Safronov, V. S. (1969), Evolution of the Protoplanetary Cloud and Formation of the Earth and Planets, Moscow: Nauka; English transl. NASA TTF-677 [1972]. [3] Burbine T. H., McCoy T. J., Meibom A., Gladman B., Keil K. (2002), Meteoritic parent bodies: Their number and identification. In: Bottke, W. F. Jr., Cellino, A., Paolicchi, P., Binzel, R. P. (Eds), Asteroids III. University of Arizona Press, Tucson, pp. 653–667. FCUP Mineralogy of V-type Near Earth Objects 258 F.1.2. Abstract of the Poster Communication in the XXIENAA Conference (7-10 September 2011) LABORATORIAL CHARACTERIZATION OF A CAMPO DEL CIELO METEORITE S. Batista1, T. M. Seixas1, M. A. Salgueiro da Silva1, A. B. Lopes2, J. M. Vieira2, B. Almeida2, D. Flores3 and F. Noronha3 1 Department of Physics and Astronomy, Faculty of Science, University of Porto, Rua do Campo Alegre 687, 4169-007 Porto, Portugal 2 Department of Ceramics and Glass Engineering/CICECO, University of Aveiro, Campus de Santiago, 3810-193 Aveiro, Portugal 3 Department of Geosciences, Environment and Land Management, Faculty of Science, University of Porto, Rua do Campo Alegre 687, 4169-007 Porto, Portugal Session: Sun, stars and planets — Communication: Poster Abstract: During the last few years, the number of known meteorites is rapidly growing up through the discovery of large concentrations of meteorites in cold [1] and hot deserts [2]. Meteorites are bodies that were fragmented due to several collisions from their parent body or were completely shielded by cosmic rays from their source body [3]. A parent body, which can be an asteroid or a comet, is the body from which the meteorite acquired its current chemical composition and mineralogical characteristics [3]. The bulk composition, mineralogy and petrology that a meteorite acquires, are functions of the original bulk composition of its parent body and the conditions of heating and melting that it has experienced during its formation. Mineralogy, texture and chemical composition of meteorites are essential for their classification and provide us important clues about the planet formation processes as well as important keys towards the main processes occurred on the evolution of the solar nebula [4]. Nowadays, asteroids represent the sole surviving population of early inner solar-system planetesimals [5-6]. Asteroids provide our only in situ record of the conditions and processes that the inner portions (~1.8-3.5 AU) of the late solar nebulae and the infant solar system have experienced [4]. Several advances in our knowledge about the diversity of asteroidal mineralogies have been made since the beginning of the use of charge-coupled device (CCD) detectors to obtain the visible and NIF reflectance spectra of small objects [3]. Our aim in this work is to study a meteorite recovered in Campo del Cielo (Argentina) for investigating meteorite processes and the associated thermophysical history. For this purpose, thin sections were taken from a layer below the exposed fusion crust and were prepared for optical mineralogy and electron microprobe analysis. Exploratory XRD measurements were performed at room temperature, in order to identify the component mineral phases. Here we report results of this study including a preliminary classification of this meteorite. FCUP Mineralogy of V-type Near Earth Objects 259 References: [1] Sephton, M. A.; Bland, P. A.; Pillinger, C. T.; Gilmou, I., The preservation state of organic matter in meteorites from Antarctica., Meteoritics & Planetary Science, Vol. 39, No. 5, p.747-754. [2] Bischoff, A. & Geiger, T. (1995), Meteorites for the Sahara: Find locations, shock classification, degree of weathering and pairing., Meteoritics (ISSN 0026-1114), vol. 30, no. 1, p. 113-122. [3] Burbine T. H., McCoy T. J., Meibom A., Gladman B., Keil K. (2002), Meteoritic parent bodies: Their number and identification. In: Bottke, W. F. Jr., Cellino, A., Paolicchi, P., Binzel, R. P. (Eds), Asteroids III. University of Arizona Press, Tucson, pp. 653–667. [4] P. Vernazza, R. P. Binzel, A. Rossi, M. Fulchignoni, M. Birlan, Solar wind as the origin of rapid reddening of asteroid surfaces, Nature 458 (2009) 993-95. [5] Gaffey, M.J., Cloutis, E.A., Kelley, M.S., Reed, K.L. (2002), Mineralogy of asteroids. In: Bottke, W.F., Cellino, A., Paolicchi, P., Binzel, R.P. (Eds.), Asteroids III. University of Arizona Press, Tucson, pp. 183–204. [6] Safronov, V. S. (1969), Evolution of the Protoplanetary Cloud and Formation of the Earth and Planets, Moscow: Nauka; English transl. NASA TTF-677 [1972]. FCUP Mineralogy of V-type Near Earth Objects 260 F.1.3. Abstract of the Poster Communication in the IJUP12 Conference (22-24 February 2012) The early stages of the Solar System Evolution: constraints from the small icy worlds S. Batista1,2,3 R. Albuquerque1,3 and T. Seixas1,3 1 Department of Physics and Astronomy, Faculty of Sciences, University of Oporto, Portugal 2 Department of Mathematics, Faculty of Sciences, University of Oporto, Portugal 3 Center for Geophysics of the University of Coimbra, University of Coimbra, Portugal The birth environment of the Sun has been subject of several studies [1]. It known that stars born in clusters [2], as it is thought the Sun also did [3]. Nowadays, the architecture of the Solar System is well known: its eight planets and the small icy worlds, which are sculpted by the force of gravity and by the phenomenon of resonances [4]. However, understanding several steps in the early Solar System is a puzzling task. In this scenario, asteroids provide our only in situ record of the conditions and processes that the inner portions ( 1.8-3.5 AU) of the late solar nebula and the infant solar system have experienced [5]. In addition, comets can also play an important role in this scenario. However, the only way that Humankind can access this material is thought meteorites, i.e extraterrestrial material that comes from the small icy worlds such as asteroids or comets. Trying to link a specific meteorite with a particular asteroid or class of asteroids can help us to understand a bit better the early stages of our Solar System. Actually, as the number of known exoplanets has recently been increasing in the last decade [6], understanding the steps evolved in the formation and evolution of the Solar System can then be extrapolated to understand others planetary systems. Our aim in this work is to try to establish possible genetic linkages between asteroids and meteorites [7]. We hope to link a specific meteorite to a certain parent or source body. For this purpose, the reflectance spectra of meteorite samples will be compared to spectra of known asteroids available at the Brown University RELAB database. References: [1] Adams, F. C. (2010), The Birth Environment of the Solar System, ARAA, 48, 47-85 [2] Lada, C. J. and Lada, E. A. (2003), Embedded Clusters in Molecular Clouds, ARAA, 41, 57-115 [3] Bobylev, V., V., Bajkova, A. T., Myllari, A. and Valtonen, M. (2010), Searching for possible siblings of the sun from a common cluster based on stellar space velocities, Astronomy Letters, 37, 8, 550-562 [4] Jewitt, D. (2009), Icy bodies in the New Solar System, Proceedings of the International Astronomical Union, IAU Symposium, 263, 3-16 FCUP Mineralogy of V-type Near Earth Objects 261 [5] Gaffey, M.J., Cloutis, E.A., Kelley, M.S., Reed, K.L. (2002), Mineralogy of asteroids. In: Bottke, W.F., Cellino, A., Paolicchi, P., Binzel, R.P. (Eds.), Asteroids III. University of Arizona Press, Tucson, pp. 183–204. [6] Butler, R. P., Wright, J. T., Marcy, G. W., Fischer, D. A., Vogt, S. S., Tinney, C. G., Jones, H. R. A., Carter, B. D., Johnson, J. A., McCarthy, C., Penny, A. J. (2006), Catalog of nearby Exoplanets, ApJ, 646, 1, 505-522 [7] Burbine T. H., McCoy T. J., Meibom A., Gladman B., Keil K. (2002), Meteoritic parent bodies: Their number and identification. In: Bottke, W. F. Jr., Cellino, A., Paolicchi, P., Binzel, R. P. (Eds), Asteroids III. University of Arizona Press, Tucson, pp. 653–667. FCUP Mineralogy of V-type Near Earth Objects 262 F.2. Abstract of the Oral communication in the Física 2012 Conference (6-8 September 2012) MINERALOGY OF V-TYPE NEAR EARTH OBJECTS S. Batista1,3, T. Seixas1,3, M. Silva1,3 e E. Alves2,3 1 Departamento de Física e Astronomia da Faculdade de Ciências da Universidade do Porto, Rua do Campo Alegre, 687, 4169-007 Porto 2 Departamento de Ciências da Terra da Faculdade de Ciências e Tecnologia da Universidade de Coimbra, Largo Marquês de Pombal, 3000-272 Coimbra 3 Centro de Geofísica da Universidade de Coimbra, Av. Dr. Dias da Silva, 3000-134 Coimbra [email protected] ABSTRACT In this work, we present part of our study on the possible genetic linkage between HED meteorites and V-type asteroids through the analysis of asteroid NIR reflectance spectra taken from the publically available MIT-UH-IRTF Joint Campaign for NEO Reconaissance and correlated spectra of HED meteorites and intimate mineral mixtures from the RELAB database. Five V-type asteroids were considered: 4 Vesta, 3908 Nyx, 4055 Magellan, 5379 Abehiroshi and 6611 (1993 VW). We report on the spectral analysis performed through the Hapke radiative transfer model and the resulting estimation of asteroids surface composition, mineralogy and grain size. FCUP Mineralogy of V-type Near Earth Objects 263 F.3. Extracurricular Internship PEEC (3rd edition) Programa PEEC-Programa de Estágios da FCUP (2010-11) Plano de Estágio (resumo) Estagiário: Sérgio Batista Tutor científico: Teresa M. Seixas (DFA – FCUP) Duração: 6 meses Laboratorial characterization of meteorites In this work we intend to study several meteorites of different types for investigating meteorite formation processes and the associated thermophysical history. The experimental work is to be performed at facilities of the universities of Porto (UP) and Aveiro (UA) according to the following sequence of activities: a) Meteorite samples preparation Thin sections of the meteorites will be prepared for optical mineralogy studies, electron microprobe analysis (UA). The thin sections will be taken from a layer below the exposed fusion crusts. For X-ray diffractometry, powder samples are to be prepared by milling slices of the original meteorites. b) Meteorite samples characterization Thin sections previously prepared will be analysed by optical mineralogy using a petrographic microscope (UA, UP). This will allow the identification of the component minerals and their microtexture. Electron microprobe analysis (UP, UA) of the same samples will help to complement this information by providing chemical composition profiles along selected paths such as grain boundaries, shearing planes and chondrules internal structures. FCUP Mineralogy of V-type Near Earth Objects 264 X-ray diffractometry (UA) on powder samples will be performed in order to determine the crystal structure of the constituent phases and to discriminate eventual polymorphic phases.