The Chemistry of Extrasolar Planetary Systems

Item Type text; Electronic Dissertation

Authors Bond, Jade

Publisher The University of Arizona.

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Link to Item http://hdl.handle.net/10150/194946 THE CHEMISTRY OF EXTRASOLAR PLANETARY SYSTEMS

Jade Chantelle Bond

A Dissertation Submitted to the Faculty of the

DEPARTMENT OF PLANETARY SCIENCES

In Partial Fulﬁllment of the Requirements For the Degree of

DOCTOR OF PHILOSOPHY

In the Graduate College

THE UNIVERSITY OF ARIZONA

2 0 0 8 2

THE UNIVERSITY OF ARIZONA GRADUATE COLLEGE

As members of the Final Examination Committee, we certify that we have read the dissertation prepared by Jade Chantelle Bond entitled The Chemistry of Extrasolar Planetary Systems and recommend that it be accepted as fulﬁlling the dissertation requirement for the Degree of Doctor of Philosophy.

Date: 31 October 2008 Dante S. Lauretta

Date: 31 October 2008 Michael J. Drake

Date: 31 October 2008 David P. O’Brien

Date: 31 October 2008 Michael R. Meyer

Date: 31 October 2008 Adam P. Showman

Final approval and acceptance of this dissertation is contingent upon the candidate’s submission of the ﬁnal copies of the dissertation to the Graduate College. I hereby certify that I have read this dissertation prepared under my direction and recommend that it be accepted as fulﬁlling the dissertation requirement.

Date: 31 October 2008 Dissertation Director: Dante S. Lauretta 3

STATEMENT BY AUTHOR

This dissertation has been submitted in partial fulﬁllment of requirements for an advanced degree at The University of Arizona and is deposited in the University Library to be made available to borrowers under rules of the Library.

Brief quotations from this dissertation are allowable without special permission, provided that accurate acknowledgment of source is made. Requests for permission for extended quotation from or reproduction of this manuscript in whole or in part may be granted by the head of the major department or the Dean of the Graduate College when in his or her judgment the proposed use of the material is in the interests of scholarship. In all other instances, however, permission must be obtained from the author.

SIGNED: Jade Chantelle Bond 4

ACKNOWLEDGMENTS

Science, like most other activities in life, is a collaborative effort and as such there are many people who deserve thanks for their assistance with producing my dissertation. First and foremost, I’d like to thank my Ph.D. advisor, Dr. Dante Lauretta, for his invaluable support and guidance throughout my time with him. There is simply no way that this work could have been completed without his assistance, both professionally and personally, along with the support of his wife, Kate. For that, I am beyond grateful. David O’Brien deserves a massive thank you for the incredible support and assistance he has provided through out my entire PhD. From first helping me to move in to Hawthorne House to now providing me with countless simulations, his support has been amazing and I am so incredibly grateful for it. Chris Tinney has truly been an outstanding collaborator, supporter and friend throughout the development of my scientific career. His assistance and guidance have been amazing and I am so grateful he took a chance on me. Similarly, Jay Melosh has always been willing to listen and help in any problems I may have had throughout my entire PhD study. His advice and assistance are greatly appreciated. I am very pleased to be able to call both of them friends. I’d also like to thank my entire PhD Committee for their support, ideas, feedback, assistance and tolerating a defence on Halloween! I found their guidance and ideas to be useful in developing this work. A thank you also goes to the entire Nine Circles research group (especially Glinda!) for all of their ideas and assistance and the Anglo-Australian Planet Search (AAPS) group for being so generous with their data. Thank you to Michael Meyer, Cathi Duncan and the LAPLACE group for their finan- cial support, including sending me to conferences in all the cool places! On a special note, I must thank Kathryn Gardner-Vandy, Kelly Kolb, Sarah Horst and Kristin Block for putting up with me, girls nights out, giving me a break and generally being wonderful friends, Eve Berger for going above and beyond, Mike and Jenny Bland and Matt Pasek for simply understanding and helping me to stay sane, David Choi and Colin Dundas for eating Tim Tams and gossiping, the Chamberlain Clan for their Aussie fixes, Bob Marcialis for reading numerous drafts while sending me both sane and insane and Pete Lanagan for being . . . well, Pete. Thank you also to the artists of the comics featured through out this dissertation and their syndications for allowing me to reprint them in this work. Thank you, Hagen, for being so supportive and not grumbling too much at all of the early morning phone calls. Last, but certainly not least, a huge thank you goes to my Mum who always got it, even when no-one else around me did. 5

DEDICATION

For Lillian and Frank Bond. “Now, Voyager, sail thou forth, to seek and ﬁnd.” -Walt Whitman, Leaves of Grass 1900

TABLE OF CONTENTS

LIST OF FIGURES ...... 9

LIST OF TABLES ...... 14

ABSTRACT ...... 16

CHAPTER 1 INTRODUCTION & BACKGROUND ...... 18 1.1 History of Extrasolar Planets ...... 18 1.2 Planetary System Chemical Properties ...... 19 1.3 Terrestrial Planets ...... 21 1.4 Summary of Work ...... 22

CHAPTER 2 R- AND S-PROCESS ELEMENTAL ABUNDANCES IN STARS WITH PLANETS ...... 25 2.1 Introduction ...... 25 2.2 Data ...... 27 2.2.1 Target Stars ...... 27 2.2.2 Spectroscopic Analysis ...... 28 2.3 Results ...... 31 2.4 Host Star Enrichment ...... 32 2.4.1 Enrichment over Solar ...... 32 2.4.2 Enrichment over Non-Host Stars ...... 33 2.5 Elemental Trends ...... 35 2.5.1 Lighter Element Trends ...... 38 2.5.2 Heavy Element Trends ...... 40 2.5.3 Correlation with Planetary Parameters ...... 42 2.5.4 Correlation with Stellar Parameters ...... 42 2.6 Discussion ...... 42 2.7 Summary ...... 45

CHAPTER 3 SOLAR SYSTEM SIMULATIONS ...... 47 3.1 Introduction ...... 47 3.2 Simulations ...... 50 3.2.1 Dynamical ...... 50 3.2.2 Chemical ...... 53 3.2.3 Combining Dynamics and Chemistry ...... 59 7

TABLE OF CONTENTS – Continued

3.2.4 Stellar Pollution ...... 59 3.3 Results ...... 62 3.3.1 Abundance Trends ...... 62 3.3.2 Variations with Time ...... 76 3.3.3 Late Veneer ...... 79 3.3.4 Hydrous Species ...... 80 3.3.5 Volatile Loss ...... 82 3.3.6 Solar Pollution ...... 87 3.4 Discussion ...... 89 3.5 Summary ...... 91

CHAPTER 4 EXTRASOLAR PLANETARY SYSTEM SIMULATIONS . . . . 94 4.1 Introduction ...... 94 4.2 System Composition ...... 96 4.3 Simulations ...... 106 4.3.1 Extrasolar Planetary Systems ...... 106 4.3.2 Dynamical Simulations ...... 110 4.3.3 Chemical Simulations ...... 115 4.3.4 Combining Dynamics and Chemistry ...... 119 4.3.5 Stellar Pollution ...... 119 4.4 Results ...... 122 4.4.1 Dynamical ...... 122 4.4.2 Chemical ...... 132 4.4.3 Stellar Pollution ...... 155 4.5 Implications and Discussion ...... 163 4.5.1 Frequency of Terrestrial Planets ...... 163 4.5.2 Planetary Types ...... 163 4.5.3 Timing of Formation ...... 164 4.5.4 Detection of Terrestrial Planets ...... 165 4.5.5 Hydrous Species ...... 167 4.5.6 Planetary Interiors and Processes ...... 171 4.5.7 Planet Habitability ...... 175 4.5.8 Biologically Important Elements ...... 177 4.5.9 Mass Distribution ...... 178 4.5.10 Stellar Pollution ...... 181 4.6 Summary ...... 181

CHAPTER 5 SUMMARY & CONCLUSIONS ...... 184

APPENDIX A STELLAR PHOTOSPHERIC ABUNDANCES ...... 186 8

TABLE OF CONTENTS – Continued

APPENDIX B SOLAR SYSTEM TERRESTRIAL PLANET ABUNDANCES . . 198

APPENDIX C MIDPLANE TEMPERATURE AND PRESSURE PROFILES . . 246

APPENDIX D HSC CHEMISTRY GAS ABUNDANCES ...... 256

APPENDIX E EXTRASOLAR TERRESTRIAL PLANET ABUDNANCES . . . 266

REFERENCES ...... 353 9

LIST OF FIGURES

1 Calvin and Hobbes, 3/17/1986 ...... 5

1.1 Piled Higher and Deeper, 2/26/2006 ...... 24

2.1 (X/H) vs. (Fe/H) for all elements studied...... 36 2.2 (X/Fe) vs. (Fe/H) for all elements studied...... 37 2.3 (heavy/light) vs. (Fe/H) ...... 41 2.4 Orbital properties of extrasolar planetary systems vs. abundance of the heavy elements...... 43 2.5 Ginger Meggs, 2/13/2008 ...... 46

3.1 Schematic of the results of the simulations of O’Brien et al. (2006). . . . 51 3.2 Radial pressure and temperature proﬁles for the nominal Soalr nebula . . 60 3.3 Normalized abundances for the CJS1 and EJS1 simulations...... 66 3.4 Al/Si v. Mg/Si for all Solar System simulated planets...... 70 3.5 Ca/Si v. Mg/Si for all Solar System simulated planets...... 71 3.6 Oxidation state plot for CJS1 and EJS1 simulated planetary abundances. . 74 3.7 Distribution of mass and its composition...... 77 3.8 Variation in composition with time for the ﬁnal planets produced by the CJS1 and EJS1 simulations...... 78 3.9 Normalized abundances for the CJS1 and EJS1 simulations after material loss in impact events has been incorporated ...... 86 3.10 Pearls Before Swine, 4/2/2007 ...... 93

4.1 Mg/Si vs. C/O for known planetary host stars...... 98 4.2 Mg/Si vs. C/O for host and non-host stars...... 102 4.3 C/O and Mg/Si distributions for host and non-host stars ...... 103 4.4 C/O and Mg/Si distributions for previously published values and this study 105 4.5 Schematic of the Extrasolar Planetary Systems studied...... 106 4.6 Mg/Si vs. C/O for planetary host stars studied...... 110 4.7 Schematic of the results of the dynamical simulations for 55Cancri and Gl777...... 124 4.8 Schematic of the results of the dynamical simulations for HD4203 and HD4208...... 125 4.9 Schematic of the results of the dynamical simulations for HD19994 and HD72659...... 126 10

LIST OF FIGURES – Continued

4.10 Schematic of the results of the dynamical simulations for HD108874 and HD142415...... 127 4.11 Schematic of the results of the dynamical simulations for HD177830. . . 128 4.12 HSC Chemistry output for HD72659 ...... 135 4.13 HSC Chemistry output for HD177830 ...... 136 4.14 HSC Chemistry output for Gl777 ...... 137 4.15 HSC Chemistry output for HD4208 ...... 138 4.16 HSC Chemistry output for 55Cnc ...... 139 4.17 HSC Chemistry output for HD142415 ...... 140 4.18 HSC Chemistry output for HD19994 ...... 141 4.19 HSC Chemistry output for HD108874 ...... 142 4.20 HSC Chemistry output for HD4203 ...... 143 4.21 Schematic of the planetary abundances for Gl777...... 147 4.22 Al/Si v. Mg/Si for planets of Gl777...... 148 4.23 Schematic of the planetary abundances for HD72659...... 149 4.24 Schematic of the planetary abundances for HD4208...... 150 4.25 Al/Si v. Mg/Si for the planets of HD4208 and HD72659...... 151 4.26 Schematic of the planetary abundances for HD177830...... 152 4.27 Schematic of the planetary abundances for 55Cnc...... 156 4.28 Schematic of the planetary abundances for HD142415...... 157 4.29 Schematic of the planetary abundances for HD19994...... 158 4.30 Schematic of the planetary abundances for HD108874...... 159 4.31 Schematic of the planetary abundances for HD42083...... 160 4.32 HSC Chemistry output for HD72659 ...... 169 4.33 HSC Chemistry output for HD4203 ...... 170 4.34 Schematic of extrasolar terrestrial planet interiors...... 174 4.35 Schematic of terrestrial planet interiors for Gl777...... 176 4.36 Solid mass distribution for the Solar-like systems...... 180 4.37 Pooch Cafe, 8/12/2006 ...... 183

5.1 Calvin and Hobbs, 2/11/1993 ...... 185

B.1 Normalized planetary abundances for CJS1 simulation...... 200 B.2 Normalized planetary abundances for CJS2 simulation...... 201 B.3 Normalized planetary abundances for CJS3 and CJS4 simulations. . . . . 202 B.4 Normalized planetary abundances for EJS1 simulation...... 203 B.5 Normalized planetary abundances for EJS2 simulation...... 204 B.6 Normalized planetary abundances for EJS3 simulation...... 205 B.7 Normalized planetary abundances for EJS4 simulation...... 206 B.8 Normalized planetary abundances for CJS1 simulation with volatile loss. . 225 11

LIST OF FIGURES – Continued

B.9 Normalized planetary abundances for CJS2 simulation with volatile loss. . 226 B.10 Normalized planetary abundances for CJS3 and CJS4 simulations with volatile loss...... 227 B.11 Normalized planetary abundances for EJS1 simulation with volatile loss. . 228 B.12 Normalized planetary abundances for EJS2 simulation with volatile loss. . 229 B.13 Normalized planetary abundances for EJS3 simulation with volatile loss. . 230 B.14 Normalized planetary abundances for EJS4 simulation with volatile loss. . 231

C.1 Midplane pressure and temperature profile for 55Cnc...... 247 C.2 Midplane pressure and temperature profile for Gl777...... 248 C.3 Midplane pressure and temperature profile for HD4203...... 249 C.4 Midplane pressure and temperature profile for HD4208...... 250 C.5 Midplane pressure and temperature profile for HD19994...... 251 C.6 Midplane pressure and temperature profile for HD72659...... 252 C.7 Midplane pressure and temperature profile for HD108874...... 253 C.8 Midplane pressure and temperature profile for HD177830...... 254 C.9 Midplane pressure and temperature profile for HD142415...... 255

D.1 Gaseous species output from HSC Chemistry for HD72659 ...... 257 D.2 Gaseous species output from HSC Chemistry for HD177830 ...... 258 D.3 Gaseous species output from HSC Chemistry for Gl777 ...... 259 D.4 Gaseous species output from HSC Chemistry for HD4208 ...... 260 D.5 Gaseous species output from HSC Chemistry for 55Cnc ...... 261 D.6 Gaseous species output from HSC Chemistry for HD142415 ...... 262 D.7 Gaseous species output from HSC Chemistry for HD19994 ...... 263 D.8 Gaseous species output from HSC Chemistry for HD108874 ...... 264 D.9 Gaseous species output from HSC Chemistry for HD4203 ...... 265

E.1 Schematic of planetary composition for Gl777 for disk conditions at 2.5×105 years and 5×105 years...... 267 E.2 Schematic of planetary composition for Gl777 for disk conditions at 1×106 years and 1.5×106 years...... 268 E.3 Schematic of planetary composition for Gl777 for disk conditions at 2×106 years and 2.5×106 years...... 269 E.4 Schematic of planetary composition for Gl777 for disk conditions at 3×106 years...... 270 E.5 Schematic of planetary composition for HD4208 for disk conditions at 2.5×105 years and 5×105 years...... 271 E.6 Schematic of planetary composition for HD4208 for disk conditions at 1×106 years and 1.5×106 years...... 272 12

LIST OF FIGURES – Continued

E.7 Schematic of planetary composition for HD4208 for disk conditions at 2×106 years and 2.5×106 years...... 273 E.8 Schematic of planetary composition for HD4208 for disk conditions at 3×106 years...... 274 E.9 Schematic of planetary composition for HD72659 for disk conditions at 2.5×105 years and 5×105 years...... 275 E.10 Schematic of planetary composition for HD72659 for disk conditions at 1×106 years and 1.5×106 years...... 276 E.11 Schematic of planetary composition for HD72659 for disk conditions at 2×106 years and 2.5×106 years...... 277 E.12 Schematic of planetary composition for HD72659 for disk conditions at 3×106 years...... 278 E.13 Schematic of planetary composition for HD177830 for disk conditions at 2.5×105 years and 5×105 years...... 279 E.14 Schematic of planetary composition for HD177830 for disk conditions at 1×106 years and 1.5×106 years...... 280 E.15 Schematic of planetary composition for HD177830 for disk conditions at 2×106 years and 2.5×106 years...... 281 E.16 Schematic of planetary composition for HD177830 for disk conditions at 3×106 years...... 282 E.17 Schematic of planetary composition for HD55Cnc for disk conditions at 2.5×105 years and 5×105 years...... 283 E.18 Schematic of planetary composition for HD55Cnc for disk conditions at 1×106 years and 1.5×106 years...... 284 E.19 Schematic of planetary composition for HD55Cnc for disk conditions at 2×106 years and 2.5×106 years...... 285 E.20 Schematic of planetary composition for HD55Cnc for disk conditions at 3×106 years...... 286 E.21 Schematic of planetary composition for HD142415 for disk conditions at 2.5×105 years and 5×105 years...... 287 E.22 Schematic of planetary composition for HD142415 for disk conditions at 1×106 years and 1.5×106 years...... 288 E.23 Schematic of planetary composition for HD142415 for disk conditions at 2×106 years and 2.5×106 years...... 289 E.24 Schematic of planetary composition for HD142415 for disk conditions at 3×106 years...... 290 E.25 Schematic of planetary composition for HD108874 for disk conditions at 2.5×105 years and 5×105 years...... 291 13

LIST OF FIGURES – Continued

E.26 Schematic of planetary composition for HD108874 for disk conditions at 1×106 years and 1.5×106 years...... 292 E.27 Schematic of planetary composition for HD108874 for disk conditions at 2×106 years and 2.5×106 years...... 293 E.28 Schematic of planetary composition for HD108874 for disk conditions at 3×106 years...... 294 E.29 Schematic of planetary composition for HD4203 for disk conditions at 2.5×105 years and 5×105 years...... 295 E.30 Schematic of planetary composition for HD4203 for disk conditions at 1×106 years and 1.5×106 years...... 296 E.31 Schematic of planetary composition for HD4203 for disk conditions at 2×106 years and 2.5×106 years...... 297 E.32 Schematic of planetary composition for HD4203 for disk conditions at 3×106 years...... 298 14

LIST OF TABLES

2.1 Spectral line list used for chemical abundance analysis...... 30 2.2 Mean difference in abundance from previously published values...... 32 2.3 Statistical analysis of abundance distributions...... 34

3.1 Properties of Solar System simulations...... 52 3.2 Chemical species included in calculations of HSC Chemistry...... 55 3.3 HSC Chemistry input values for Solar System Simulations ...... 56 3.4 T50%condensation for the Solar System...... 57 3.5 Percentage of each element assumed to be lost in impact events...... 85 3.6 Mean change in solar photospheric abundances produced by pollution via accretion of solid material...... 88

4.1 Statistical analysis of the host and non-host star distributions of Mg/Si and C/O...... 104 4.2 Orbital parameters of known extrasolar planets ...... 107 4.3 Target star elemental abundances in logarithmic units...... 108 4.4 Target star elemental abundances normalized to 106Si atoms...... 109 4.5 Statistical analysis of the embryo and planetesimal and embryo only simulations ...... 114 4.6 HSC Chemistry input values for extrasolar planetary systems studied. . . 117 4.7 Extrasolar host star accretion rates ...... 120 4.8 Stellar convective zone masses ...... 121 4.9 Properties of simulated extrasolar terrestrial planets...... 129 4.10 Degree of radial mixing for extrasolar planetary system simulations. . . . 130 4.11 Formation time...... 131 4.12 T50% condensation for extrasolar planetary systems studied...... 134 4.13 Change in host star photospheric abundances...... 162

A.1 Stellar abundances of all AAPS target stars...... 187 A.2 Stellar abundances for AAPS target stars normalized to 106 Si atoms. . . 192

B.1 Predicted bulk planetary abundances for the terrestrial planets of the O’Brien et al. (2006) simulations...... 207 B.2 Ensemble-averaged bulk predicted planetary abundances...... 221 B.3 Difference in abundance between t = 2.5×105yr and t = 3×106yr. . . . . 223 B.4 Bulk Planetary Abundances with Volatile Loss...... 232 15

LIST OF TABLES – Continued

E.1 Bulk Planetary Abundances for Gl777...... 299 E.2 Bulk Planetary Abundances for HD4208...... 306 E.3 Bulk Planetary Abundances for HD72659...... 313 E.4 Bulk Planetary Abundances for HD177830...... 320 E.5 Bulk Planetary Abundances for 55Cnc...... 327 E.6 Bulk Planetary Abundances for HD142415...... 331 E.7 Bulk Planetary Abundances for HD19994...... 338 E.8 Bulk Planetary Abundances for HD108874...... 345 E.9 Bulk Planetary Abundances for HD4203...... 349 16

ABSTRACT

This work examines the chemical nature of extrasolar planetary systems, considering both the host star and any potential terrestrial planets located within the system. Extrasolar planetary host stars are found to be enriched over non-host stars in several r- and s-process elements. These enrichments, however, are in keeping with general galactic chemical evolution trends. This implies that host stars have not experienced any unusual chemical processing or pollution and that the observed enrichments are primordial in nature. When combined with detailed chemical models, the dynamical models of O’Brien et al. (2006) are found to produce terrestrial planets with bulk elemental abundances in excellent agreement with observed planetary values. This clearly indicates that the combination of dynamical and chemical modeling applied here is successfully reproducing the terrestrial planets of the Solar System to the first order. Furthermore, these planets are found to form with a considerable amount of water, negating the need for large amounts of exogenous delivery. Little dependence on the orbital properties of Jupiter and Sat- urn is observed for the main rock-forming elements due to the largely homogenous disk composition calculated. The same modeling approach is applied to known extrasolar planetary systems. Ter- restrial planets were found to be ubiquitous, forming in all simulations. Generally, small (< 1ML) terrestrial planets are produced close to their host star with little radial mixing occurring. Planetary compositions are found to be diverse, ranging from Earth-like to refractory dominated and C-dominated, containing significant amounts of carbide material. Based on these simulations, stars with Solar elemental ratios are the best place to focus future Earth-like planet searches as these systems are found to produce the most Earth- like terrestrial planets which are located within the habitable zones of their systems and 17 containing a significant amount of water. C-rich planets, although unusual, are expected to exist in ∼20% of known extrasolar planetary systems based on their host star photospheric compositions. These planets are unlike any body we have previously observed and provide an exciting avenue for future observation and simulation. 18

CHAPTER 1

INTRODUCTION & BACKGROUND

1.1 History of Extrasolar Planets

Extrasolar planets (planets orbiting stars other than our own Sun) are a relatively new branch of the astronomical and planetary sciences. After the discovery of (the now de- moted) Pluto in 1930 (see Tombaugh (1946) for a review), planet finding activities appeared to have reached an end for the foreseeable future. Between 1930 and 1992, several brown dwarfs were discovered orbiting other solar-type stars (Henry and Kirkpatrick, 1990). Commonly referred to as “failed stars”, brown dwarfs are low-mass celestial objects (M≥10MJUP) that formed by stellar processes but did not obtain the critical mass required to sustain nuclear burning within their core. As such, they were widely believed to be too large to be planets in the classical sense. Other claims for planetary detections were also made during this period (Strand, 1944; Reuyl and Holmberg, 1943; van de Kamp, 1963, 1969) but these were never independently verified or were later shown to be false, produced by timing artifacts or instrumentation errors. It wasn’t until 1992 that the first confirmed detection of an extrasolar planet occurred when two bodies were found to be orbiting the millisecond pulsar PSR 1257+12 (Wol- szczan and Frail, 1992; Backer et al., 1992). Pulsars had not previously been thought to be likely hosts for extrasolar planets due to their formation mechanism. Pulsars are rapidly rotating neutron stars produced during supernova events. They are exposed to ob- servers when the outer layers of the original star are removed. The resulting reduction in mass of the host star was believed to be significant enough that the gravitational attraction between the host star and any orbiting planets would be reduced to the level where the planet would no longer be gravitationally bound to the remaining neutron star, and would 19 therefore be lost. The first detection of an extrasolar planet orbiting a solar-type star occurred two years later in 1994 with the announcement of a planetary body orbiting 51 Paegasi (Mayor and Queloz, 1995). This detection was later confirmed in 1995 (Marcy and Butler, 1996) and began a virtual “planet-rush”. Over the next 18 months, we progressed from knowing of no other planets orbiting solar-type stars to having detected nine extra-solar planetary systems by the end of 1996 (although two companions were later revised to possible brown dwarf status because of their high mass). As of September 2008, we currently know of 294 planets orbiting a total of 252 solar-type stars (there are 30 multiple planet systems). The vast majority of these detections have occurred via the radial velocity method (for a summary of this method, see Schneider (1999)), although other methods such as that of transiting photometry and microlensing are becoming increasingly important in future planet searches as we seek to detect terrestrial-sized planetary bodies and utilize space- based observing programs such as Kepler. Based on the detection statistics of the current planet search programs, Marcy and Butler (1998) estimated that the Milky Way could be “home” to up to 10 billion planets. While this rough estimate is by no means a definitive number, it does give us some idea of the scale of the planet search that we have begun. Upon examination, the known extrasolar planetary systems can be seen to be vastly different to our own solar system in terms of both the bulk chemistry of the host star and the orbital properties of the companion planets. Unravelling the causes of these two differences have proven to be the most intriguing aspects of current extrasolar planetary studies.

1.2 Planetary System Chemical Properties

Extrasolar planetary host stars are known to be somewhat chemically anomalous. They have been found to be enriched in iron as compared to other average non-host ﬁeld stars (Gonzalez, 1997, 1998; Butler et al., 2000; Gonzalez and Laws, 2000; Gonzalez et al., 20

1999; Gonzalez and Vanture, 1998; Santos et al., 2000, 2001, 2004; Gonzalez et al., 2001; Smith et al., 2001; Reid, 2002; Fischer and Valenti, 2005; Bond et al., 2006). At present, this observed chemical anomaly is the only externally observable correlation between the properties of a star and the presence of a planetary companion. Furthermore, they are potentially also enriched in other lighter, key planet building elements (Mg, Si, O), (Gonzalez and Vanture, 1998; Gonzalez et al., 2001; Santos et al., 2000; Bodaghee et al., 2003; Fischer and Valenti, 2005; Beirao˜ et al., 2005; Bond et al., 2006), although not with the same statistical significance as has been observed for iron. However, in addition to these simple elemental enrichments, several extrasolar planetary systems are abnormal in key cosmochemical ratios. Several systems are found to have C/O ratios at or above unity or Mg/Si ratios almost a factor of two higher than in our Solar System. If these ratios are primordial, established in the giant molecular cloud from which these systems subsequently formed, then the chemistry of the planet building material and ultimately any planets within these systems will be drastically different to that of our own Solar System. The exact origin of this observed elemental enrichment has been the subject of much debate. The two main hypotheses for the observed metalicity trend are the ‘pollution’ model whereby metal rich material is added to the outer envelope of the star during the planetary formation process and the ‘primordial’ model in which the gas cloud from which the star formed was already enriched prior to stellar formation occurring. Propo- nents of the primordial model cite the fact that approximately 6 M⊕ of iron would need to be systematically added to the host star’s photosphere after the dissipation of the protoplanetary disk in order to produce the observed metallicity trend (Murray et al. 2001). However, no clear correlation between the metallicity of the host star and the orbital parameters of the remaining planetary companions can be seen. It is difficult to think of any feasible process by which the presence of the remaining planetary companion could have caused the accretion of another approximately Jupiter-mass planet, independent of its final orbital parameters. Furthermore, it is even more challenging to explain why almost 21 all stars currently known to have planetary companions would have accreted significant amounts of iron during the planetary formation process while only a small percentage of non-host stars show a similar level of enrichment. However, neither of these arguments clearly excludes the pollution hypothesis with proponents suggesting that an enrichment in lighter elements, and especially in 6Li, would favor the pollution theory (Santos et al. 2001). Furthermore, enrichment is not essential for planetary formation as indicated by the existence of a planetary hosts with iron depletions relative to the Sun (such as the known host star HD155358 which has [Fe/H] = −0.68). As a result, no clear consensus has emerged as to which explanation is believed to be correct and resolving this issue and examining its implications on planetary formation and composition is one of the outstanding questions within extrasolar planetary science.

1.3 Terrestrial Planets

The idea of terrestrial planets existing within known extrasolar planetary systems isn’t new. In fact, Earth-sized bodies are expected to be more prevalent throughout the universe than their more massive gas giant companions (Marcy et al., 2000). However, they remain well beyond current detection levels. While we have been able to detect several “super Earths” (bodies with masses between 1L and 10 ML) (Rivera et al., 2005; Beaulieu et al., 2006; Lovis et al., 2006; Udry et al., 2007), we are still unable to examine bodies with masses comparable to that of the Earth. Only Kepler and possibly Darwin will be able to determine if such bodies do exist within other systems. Dynamical simulations have shown that several systems contain a region of orbital space in which such bodies could potentially exist for long periods of time (e.g. Barnes and Raymond 2004; Raymond and Barnes 2005; Asghari et al. 2004). Others have gone one step further and simulated actual terrestrial planet formation within these systems (e.g. Raymond et al. 2005; Mandell et al. 2007). Based on these simulations, terrestrial planets have been found to form in a wide variety of dynamical systems, indicating 22 that they are indeed likely to exist in the majority of known extrasolar planetary systems. Only Raymond et al. (2006) has modeled terrestrial planet formation within specific systems, finding that terrestrial planets could form in one of the systems simulated (55Can- cri), while small bodies comparable to asteroid sized objects would be stable in another (HD38529). Studies have also suggested that a significant portion of these planets will reside in the habitable zone of their stellar system (Mandell et al., 2007). However, no previous studies have considered the chemical composition of these potential planets. Given the wide variety of host star compositions observed, it is likely that terrestrial planets will display a similar variety of compositions. Of particular interest are those systems with exceptionally high C/O values. Such systems will contain C-based species such as SiC, TiC and graphite as the main planet forming material, therefor producing C-rich planets unlike anything we have previously observed. These variations will have drastic implications on a wide variety of planetary properties, including planetary processes (such as volcanism and tectonics), our ability to detect and observe extrasolar terrestrial planets and their ability to host life (as we currently know it). While the detection of a true Earth analog extrasolar terrestrial planet remains the “holy grail” of extrasolar planetary searches, the potential diversity in planetary compositions is a fasci- nating aspect of planetary studies.

1.4 Summary of Work

Given the anomalies observed within the extrasolar host star population, it is thus natural to wonder about other aspects of the extrasolar planetary system. Why are extrasolar planetary host stars enriched above other non-host field stars? Given the unusual and often extreme orbital parameters of the known giant extrasolar planets, could terrestrial planets still form in these systems? Would such planets be stable on geologic timescales? If the chemistry of the host star is unusual (as compared to other non-host field stars), then what is the chemical composition of the planets in the system like? Are all of the bodies 23 iron rich? Would we see any planets with a composition similar to that of the Earth? Would it be chemically possible for life as we know it to develop within such a system? This work begins to answer some of these questions, focussing specifically on the nature of the observed elemental enrichments in known host stars, the potential for the formation of extrasolar terrestrial planets and the chemical abundances of these planets. In the second chapter, we determine the stellar abundances of five elements produced by the rapid (r-) and slow (s-) processes, in addition to three lighter elements, in the known host and non-hosts stars observed as part of the Anglo-Australian Planet Search. This is the first study of its kind to consider elements beyond the iron peak. Such elements are produced in specific yet vastly different formation environments and as such provide us with information about the evolutionary history of the material that has been incorporated into extrasolar planetary systems. In the third chapter, we determine the detailed bulk composition of the terrestrial planets produced in dynamical simulations for the Solar System. These predicted compositions are compared to the bulk compositions of the actual terrestrial planets as an examination of the chemical validity of the dynamical simulations currently being used to model terrestrial planet formation. This study is the first to undertake such a high resolution and detailed chemical test of dynamical models of terrestrial planet formation. Finally, in the fourth chapter, we expand this work out to consider extrasolar terrestrial planets. Identical dynamical formation and chemical composition simulations are undertaken for nine known extrasolar planetary systems, varying in both their dynamics and chemical compositions. While other studies have considered the potential of terrestrial planet formation in extrasolar planetary systems (e.g. Raymond et al. 2006), this is the first time that the dynamics of formation have been simultaneously considered with the chemistry of the associated material. These studies will assist us in obtaining a better sense of both the dynamical and the chemical nature of terrestrial planet formation, both within our Solar System and within known extrasolar planetary systems. 24

Figure 1.1: “Piled Higher and Deeper” by Jorge Cham. www.phdcomics.com. Reprinted with permission. Originally published 2/26/2006. 25

CHAPTER 2

R- AND S-PROCESS ELEMENTAL ABUNDANCES IN STARS WITH PLANETS

2.1 Introduction

Extrasolar planets are known to preferentially orbit metal-enhanced stars. Chemical analyses of known host stars have shown that they appear to be metal enriched compared to a sample of “average” F, G and K stars not known to harbor planets (Gonzalez, 1997, 1998; Gonzalez and Laws, 2000; Gonzalez et al., 1999; Gonzalez and Vanture, 1998; Santos et al., 2000, 2001, 2003; Gonzalez et al., 2001; Smith et al., 2001; Reid, 2002; Fischer and Valenti, 2005; Bond et al., 2006). In addition to this metal enrichment, other elements have also been shown to exhibit similar trends, although not as statistically significant or with such large host and non-host differences (Gonzalez et al. 2001; Santos et al. 2000; Bodaghee et al. 2003; Fischer and Valenti 2005; Bond et al. 2006). However, the vast majority of abundance studies completed so far have focussed on elements with atomic number (Z)≤30 (i.e. those located before the iron stability peak). These elements are produced by a variety of processes (for example alpha particle addition, the CNO cycle and stellar burning reactions) in stellar interiors during main sequence evolution. Elements located beyond the iron peak, however, are produced via neutron-capture reactions, specifically the rapid (r-) and slow (s-) processes. Here rapid and slow refers to the speed of neutron capture with respect to the β-decay rate of the nuclei. In the r-process, neutron capture occurs before β-decay of the unstable nuclei can occur. Alternatively, in the s-process neutron capture occurs less frequently, thus allowing the nuclei to undergo β-decay before capturing another neutron and effectively allowing the nuclide to remain within the valley of β stability. Due to the different neutron fluxes required for each process (∼1023 neu- 26 trons cm−2 s−1 for the r-process vs. ∼105 neutrons cm−2 s−1 for the s-process (Clayton, 1968), the r- and s-processes occur in different stellar environments. There is some debate as to exactly where the r-process occurs (see e.g. Qian 2004), but it is believed to primarily occur in type II supernova events. The s-process is thought to produce its heav- ier elements (such as Ba and Ce) in the interior of lower mass AGB stars and its lighter elements (such as Sr, Y and Zr) during the He-burning stages of stellar evolution for larger mass stars (Reddy et al., 2003). Due to these vastly different settings, the abundances of the heavy elements produced by these processes can provide information on the history of the material later incorporated into both the host star and planets themselves, as well as testing models of galactic chemical evolution (e.g. Fenner et al. 2006 and Lanfranchi et al. 2008). These elements have been part of several previous spectroscopic studies of stars without planets (e.g. Edvardsson et al. 1993; Reddy et al. 2003; Allende Prieto et al. 2004; Bensby et al. 2005; Reddy et al. 2006), however only Ba (Huang et al., 2005) and Eu (Gonzalez and Laws, 2007) have been specifically examined in a small number of known extrasolar planetary host stars. Studies of stellar abundances have become critical to our understanding of planetary formation processes. However, in spite of significant advances in atmospheric models, stellar interiors and atomic line data in recent years, the measurement of stellar abundances is still an intricate process. In particular, choices made in the way an analysis is carried out can result in systematic abundance differences on scales similar to the effects we would most like to understand in the stars themselves. Additionally, in many cases there is no consensus on the “correct” choice in these techniques. To make headway in this field, therefore, it is essential to perform abundance studies in a manner immune to such systematic errors, by analyzing both samples of interest, and control samples, in an identical manner. For the abundances of exoplanetary host stars, that means both host and non-host stars must be analyzed in an identical manner. This is the primary goal of this current study. Abundances of five r-and s-process elements, in addition to three other lighter ele- 27 ments, were derived for all the planet-hosting and non-planet-hosting stars in the Anglo- Australian Planet Search (AAPS) data set with viable template spectra (28 hosts, 108 non-hosts), so that robust conclusions can be reached about the differences in their elemental abundances. The inclusion of these heavy elements begins to extend the spectroscopic studies of known host stars beyond the iron peak, thus continuing the search for additional chemical anomalies within these systems, while also providing an independent check of nucelosynthesis models and previously published abundances and trends for known host stars. The lighter elements selected for study (O, Mg and Cr) are included so as to complement the previous study of Bond et al. (2006) while also providing valu- able information as to the cause of the metal-enrichment commonly seen in known host stars. The majority of this chapter appeared as Bond et al. (2008) Beyond the Iron Peak: r- and s-Process Elemental Abundances in Stars with Planets, Astrophysical Journal, 682, 1234-1247.

2.2 Data

2.2.1 Target Stars

The F- and G-type stars, observed at the 3.9m Anglo-Australian Telescope (AAT) since January 1998 as part of the AAPS program were selected for this present study (Butler et al., 2001, 2002; Tinney et al., 2001, 2002, 2003, 2005, 2006; Jones et al., 2002b,a, 2006; Carter et al., 2003; McCarthy et al., 2004; Jenkins et al., 2006; Wright et al., 2007; O’Toole et al., 2007). As at January 2007, 31 stars present within the AAPS sample were known to be planet hosts, of which 28 had spectra useful for the purposes of this study.

0 Stars known to be young (age < 3 Gyr), active (logR HK > -4.5) or with other stars within 500 are rejected from the AAPS search. For a more detailed description of the data and details of the criteria applied to the AAPS target stars, the reader is referred to Butler et al. (1996) and Tinney et al. (2005). Of the 28 host stars considered here, 26 have had some 28 common elemental abundances previously determined by other authors.

2.2.2 Spectroscopic Analysis

The method utilized in this study closely follows that outlined in Bond et al. (2006) who studied Fe, C, Na, Al, Si, Ca, Ti and Ni. Spectra encompassing the entire visible spectrum from 4820 to 8420A˚ with a signal-to-noise ratio (S/N) between 200 and 300 per spectral pixel at resolution λ/∆λ ≈ 80000 were obtained via the University College London Echelle Spectrograph (UCLES) using the 31.6 line/mm echelle grating as part of the AAPS program. The raw data were reduced and processed so as to produce a one-dimensional spectra, suitable for spectral analysis. This study does differ slightly from Bond et al. (2006) in the method used to determine the equivalent widths of absorption lines. In Bond et al. (2006), equivalent width estimates were obtained via direct integration over the line. In this study we follow the method of other similar studies (eg. Santos et al. 2000; Gonzalez et al. 2001; Santos et al. 2001) and make Gaussian line fits to the spectra using the IRAF task splot in the package noao.onedspec (due to difficulties in automating the previous script). This slight difference in methodology does not produce any significant difference in the final abundances obtained. Five heavy elements were analyzed for the first time as part of this study and they were primarily selected based on their process of production. Y, Zr and Ba are all produced primarily by the s-process while Eu is primarily produced by the r-process and Nd is an almost even mix between the two based on Solar System abundances (Arlandini et al. 1999; Simmerer et al. 2004). The line list utilized in this study is shown in Table 2.1 and is derived from Gilli et al. (2006) (for Mg and Cr), Reddy et al. (2003)(for Y, Zr, Ba, Eu and Nd) and Den Hartog et al. (2003) (for Nd). Additional Ba lines have been used in other studies, but were neglected here as they gave consistently lower abundances by approximately 0.60 dex. While the use of more lines in determining an abundance should produce a more robust value, the cause of this offset could not be determined and as 29 such the single stronger line at 6496.91A˚ was used as this line produced a Ba abundance close to the reported solar value for a solar spectrum. Atomic parameters for each line were obtained from the NIST Atomic Spectra Database, Version 3.0 (for O, Mg and Cr), Pitts and Newsom (1986) and Hannaford et al. (1982) (for Y), Reddy et al. (2003) (for Ba), Den Hartog et al. (2003) (for Nd) and the Kurucz atomic line database1 (for Zr and Eu). All of the atomic parameters applied here have also been utilized in previous studies and produce solar elemental abundances well within errors of those published elsewhere, when applied to a solar spectrum, thus giving us confidence in applying them here. Elemental abundances were obtained via standard local thermodynamic equilibrium analysis, as has been done by previous studies (see Santos et al. 2000; Gonzalez et al. 2001; Santos et al. 2001; Bond et al. 2006). A revised version of Sneden’s (1973) MOOG abundance code entitled width6 (Ryan 2005, personal communication) was once again used in conjunction with a grid of Kurucz (1993) ATLAS9 atmospheres2 to obtain the final elemental abundances. As all of the non-host, and all but 7 of the host stars, had been the subject of an earlier study (Bond et al. 2006), previously published stellar atmospheric parameters were used. For those stars without previously determined values, we followed the same method as used in Bond et al. (2006) to obtain the values and refer the reader to the paper for more details. As the OI triplet lines are known to suffer from non-LTE

2 effects, the corrections of Takeda (2003) for ξt=1km/s and log g=4.0 cm/s were applied to obtain our ﬁnal O abundances. When applied to a solar spectrum, this method produced abundances in agreement with those of Asplund et al. (2005).

Errors were obtained via sensitivity studies. Teff , log g, [Fe/H] and microturbulence were varied in turn by a specified amount (±100 K for Teff , ±0.3 dex for log g, ±0.3 dex for metallicity and ±0.05 dex for microturbulence) and the resulting variation in the elemental abundance was determined. Thus the final error was obtained by summing in quadrature the sensitivity errors, continuum placement error (typically 0.05 dex) and stan-

1http://www.pmp.uni-hannover.de/cgi-bin/ssi/test/kurucz/sekur.html 2http://kurucz.harvard.edu or http://www.stsci.edu/hst/observatory/cdbs/k93models.html 30

Table 2.1: Spectral line list used for chemical abundance analysis. Atomic parameters are also provided. See text for references.

λ Log gf χl (A)˚ (eV) Mg I 5711.09 -1.71 4.35 6318.72 -1.99 5.11 Cr I 5304.18 -0.69 3.46 5312.87 -0.56 3.45 5318.81 -0.69 3.44 5783.09 -0.50 3.32 5783.89 -0.29 3.32 O I 7771.94 0.37 9.15 7774.17 0.22 9.15 7775.39 0.002 9.15 Ba II 6496.91 -0.41 0.60 Y II 4854.87 -0.01 0.99 4900.12 -0.09 1.03 5087.43 -0.17 1.08 5200.42 -0.49 0.99 5402.78 -0.63 1.84 Zr II 5112.28 -0.59 1.66 Eu II 6645.13 0.20 1.38 Nd II 4914.18 -0.70 0.38 5234.19 -0.51 0.55 31 dard deviation of each mean abundance (where the elemental abundance was determined from more than one spectral line):

2 2 2 2 2 2 2 σﬁnal = σstd + σcontinuum + σTeﬀ + σlogg + σ[Fe/H] + σmicro (2.1)

2.3 Results

The stellar elemental abundances are shown in Table A.1(in the standard astronomical logarithmic form) and Table A.2 (in the more cosmochemically useful form with abundances normalized to 106 Si atoms). The notation of − for an abundance indicates that a value could not be obtained from the spectrum due to noise. Additionally, for ease of comparison in Section 5.1, C and Si stellar elemental abundances, along with the normalized C abundances, for all target stars are presented in Tables A.1 and A.2. The C and Si abundances were previously published in Bond et al. (2006). Other authors have previously determined abundance values for several of the elements also studied here with many of these abundances differing from our values. Differ- ences between the present study and that of others is not a signiﬁcant issue for the primary thrust of this paper, which is to compare host and non-host stellar abundances which have been measured in an identical fashion. However, in the interests of completeness it is noted that Mg abundances were determined for 29 common stars (20 hosts) by Beirao˜ et al. (2005), Cr in 28 stars (20 hosts) and Mg in 29 stars (20 hosts) by Gilli et al. (2006), Cr in 25 stars (16 hosts) by Bodaghee et al. (2003), O in 8 stars (4 hosts) by Ecuvillon et al. (2006a), O and Eu in 4 host stars, Cr and Mg in 1 host star by Gonzalez and Laws (2007) and O in 1 host star by Santos et al. (2000). The mean differences between our values and those previously published are shown in Table 2.2 (for those samples having more than one common star) with the difference being deﬁned as the abundance from this 32

Table 2.2: Mean difference in abundance between values determined in the present study and previously published values. Difference is deﬁned as Abundancethis study - Abundancepreviously published. Study Element Difference Sample Size Beriao˜ et al. (2005) Mg −0.18± 0.02 29 Gilli et al. (2006) Cr 0.00± 0.02 28 Mg −0.18± 0.02 29 Bodaghee et al (2003) Cr 0.02± 0.01 25 Ecuvillon et al (2006) O 0.11± 0.07 8 Gonzalez & Laws (2007) O −0.03± 0.04 4 Eu −0.22± 0.16 4

study minus the published abundance. Generally, the results presented here can be seen to be in agreement with those previously published for Cr and O with a signiﬁcantly larger mean difference occurring for Mg and and a large deviation occurring for Eu. The differences between our abundances and those previously published are believed to be due to the use of a smaller number of lines in determining the abundance (for Mg), the use of different methods (for O and Eu), the use of different atomic parameters (for Eu) and the use of different non-LTE corrections (for O).

2.4 Host Star Enrichment

2.4.1 Enrichment over Solar

The mean and median abundances, standard deviation and the difference between the host and non-host stars for all target stars can be seen in Table 2.3 for each element. The quoted uncertainties are the standard error in the mean, and the median uncertainty from the algorithm of Kendall et al. (1987)3. The mean values of this study for the known host stars are all consistent to within the 1σ value of those listed by Beirao˜ et al. (2005) and Gilli et al. (2006). The data show that in general known extrasolar planetary host stars

3For a distribution with N values, the error in the median is the range in values on either side of the √ median which contains ( N)/2 values 33 differ only slightly from the mean solar abundance patterns with all of the median abundances being well within 1σ of the solar abundance - as concluded by previous studies (eg. Ecuvillon et al. 2004; Bodaghee et al. 2003). This is reassuring as the Sun is itself is obviously a planetary host star with abundances enhanced over those of most other stars in the solar neighborhood (based on the abundances of Asplund et al. 2005). In many respects, therefore, the Sun is not a typical ﬁeld star, based on its abundances and its multiple planetary companions. The largest enrichment over solar is seen in Nd and Zr, with Cr showing a smaller enrichment and Mg showing minimal enrichment. Eu showed the largest depletion relative to solar values, with Y and Ba also showing mild to moderate depletions. Only O produced a mean abundance equal to the Solar abundance. Similarly, the non-host stars can also be seen be depleted when compared to solar abundances for almost all of the elements studied, with the largest depletion being −0.16 for Y and Eu. It is also worth noting that three of the ﬁve heavy elements examined show a mean depletion relative to solar for both the host and non-host stars. Of these three, two are produced by the s-process (Y & Ba) with the remaining element (Eu) produced by the r-process.

2.4.2 Enrichment over Non-Host Stars

A more powerful comparison is obtained by comparing the host and non-host populations to each other. On doing so, it can be seen that host stars are systematically enriched over non-host stars in all elements studied. The enrichment ranges in size from 0.06 (for O) to 0.11 for (for Cr and Y) (see Table 2.3). This difference between the host and non-host populations can also be seen in the results of the Kolmogorov-Smirnov (K-S) statistical test. Designed to test whether two distinct populations were drawn from the same parent sample, the K-S test determines the distance between the cumulative probability distribution function of the two populations. The probability of the two samples being drawn from the same parent dataset is returned. The abundances determined here showed a signiﬁcant difference between the host and 34

Table 2.3: Statistical analysis of abundance distributions. Difference is deﬁned as host star abundance − non-host star abundance. Numbers in parentheses indicate the sample size. K−S result is the % result of a Kolmogorov-Smirnov statistical test to determine the probability of host and non-host stars having the same parent sample based on their abundance distribution. Non-Planetary Planetary Hosts Difference K−S Hosts Result O: Mean −0.06± 0.02 (90) +0.00± 0.03 (27) 0.06 8.8 +0.04 Median −0.05± 0.02 +0.00−0.03 0.05 Std. Dev. 0.15 0.17 Mg: Mean −0.09± 0.01 (90) −0.02± 0.03 (28) 0.07 3.18 +0.04 +0.06 Median −0.10−0.01 +0.01−0.07 0.11 Std. Dev. 0.14 0.16 Cr: Mean −0.03± 0.02 (90) +0.08± 0.03 (28) 0.11 0.55 +0.04 +0.03 Median +0.00−0.01 +0.12−0.02 0.12 Std. Dev. 0.16 0.17 Y: Mean −0.16± 0.01 (90) −0.05± 0.03 (27) 0.11 0.01 Median −0.15± 0.03 +0.00± 0.02 0.15 Std. Dev. 0.14 0.16 Zr: Mean −0.03± 0.02 (85) +0.06± 0.03 (26) 0.09 2.5 +0.03 +0.03 Median −0.01−0.02 +0.05−0.06 0.06 Std. Dev. 0.15 0.16 Ba: Mean −0.11± 0.02 (85) −0.01± 0.04 (26) 0.10 1.8 +0.04 Median −0.11± 0.02 +0.05−0.05 0.16 Std. Dev. 0.18 0.19 Eu: Mean −0.16± 0.02 (73) −0.10± 0.03 (26) 0.06 3.84 +0.03 +0.04 Median −0.17−0.04 −0.09−0.05 0.08 Std. Dev. 0.14 0.15 Nd: Mean +0.00± 0.02 (84) +0.06± 0.03 (27) 0.06 2.15 +0.01 Median +0.01± 0.03 +0.09−0.05 0.08 Std. Dev. 0.14 0.15 35 non-host samples with probabilities of the same parent sample of stars for the two populations ranging from 0.01% for Y to 8.8% for O (see Table 2.3). This supports our claim that host star elemental abundances are signiﬁcantly different to those of non-host stars and furthermore that host stars are enriched over non-host stars for all elements studied.

2.5 Elemental Trends

Plots of the current results are shown in Figures 2.1 and 2.2 - in Figure 2.1 [X/H] versus [Fe/H] (as more commonly used in previous studies of planet host star abundances) is presented, while in Figure 2.2 [X/Fe] versus [Fe/H] (as usually analyzed in cosmochemical studies) is shown. Two signiﬁcant outliers can be seen - one host and one non-host sit- ting below the general trend for O, Cr and Mg. These stars are HD142415 ([Fe/H]=0.02, host) and HD199288 ([Fe/H]=0.04, non-host). These stars can be seen to be depleted (compared to solar abundances) in the majority of elements studied here except for Fe, suggesting that they have formed from Fe-rich precursor material based on the assumption that it is easier to enrich one element than it is to deplete seven other elements. The fact that HD142415 is also mildly enriched in both Ba and Y (with no Zr abundance available) could also possibly indicate that the material had been processed through an s-process environment, either an AGB star or the He-burning stage of a larger mass star. Fe and Fe precursors: The [X/H] trends in Figure 2.1 are in agreement with the understanding we currently have about the nucelosynthetic origin of the elements. All of the elements located before the Fe peak (here O, Mg and Cr) increase linearly with increasing [Fe/H] with Pearson product-moment correlation coefﬁcients (r) above 0.70 for both host and non-host stars. This can be easily understood as the stellar evolutionary processes that serve to increase the amount of Fe present in later generations of stars also produce the pre-iron peak elements in various amounts. Thus as the stellar Fe content increases, so too would the amount of pre-iron peak elements (neglecting any unusual mixing or other nebula interactions). 36

0.5 0.5 0.5

0 0 0

-0.5 -0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5 -0.5 0 0.5

0.5 0.5 0.5

0 0 0

-0.5 -0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5 -0.5 0 0.5

0.5 0.5

0 0

-0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5

Figure 2.1: [X/H] vs. [Fe/H] for all elements studied.Open squares represent non-host stars and ﬁlled squares represent host stars. Typical error bars are shown in the upper left of each panel. Numerical values are provided in Table A.1.Left Column: O, Mg and Cr. Center Column: Y, Zr and Ba. Right Column: Eu and Nd. 37

0.5 0.5 0.5

0 0 0

-0.5 -0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5 -0.5 0 0.5

0.5 0.5 0.5

0 0 0

-0.5 -0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5 -0.5 0 0.5

0.5 0.5

0 0

-0.5 -0.5

-0.5 0 0.5 -0.5 0 0.5

Figure 2.2: [X/Fe] vs. [Fe/H] plots for all elements studied. Open squares represent non- host stars and ﬁlled squares represent host stars. Typical error bars are shown in the upper left of each panel. Numerical values are based on those shown in Table A.1 (in the [X/H] form). Left Column: O, Mg and Cr. Center Column: Y, Zr and Ba. Right Column: Eu and Nd. 38

s-process: The r- and s-process elements, however, are less certain. All three s-process elements (Y, Zr and Ba) still display to varying degrees the same trend of increasing abundance with increasing [Fe/H] as the pre-iron peak elements. One possible explanation for this observed trend is that the increase in s-process elemental abundances is due to the increase in the number of seed nuclei (e.g. Fe atoms) available. Due to the nature of the s-process, it is reliant on the sufficient availability of seed nuclei to be able to proceed. Thus as the metallicity increases, so too does the abundance of s-process species. r-process: Unlike the s-process elements, the r-process element (Eu) and the mixed source element (Nd) do not display a strong correlation with increasing [Fe/H]. Obser- vations of metal poor stars have shown that the abundances of s-process elements such as Y and Ba decrease faster with metallicity than the abundances of r-process elements such as Eu (Spite and Spite 1978). This has been attributed to a lack of appropriate seed nuclei inhibiting the s-process significantly more than the r-process. We are alternatively extending this into the metal-rich regime to conclude that the r-process is not as reliant on the presence of elements such as Fe, thus explaining its lack of a strong dependance upon metallicity. From Figure 2.2, it can be seen that the overall [X/Fe] trends visible here are in good agreement with those identified by Bodaghee et al. (2003) and Gilli et al. (2006) (which are discussed in more detail below). They can also be seen as extensions into the high metallicity region of those trends identified by Reddy et al. (2006).

2.5.1 Lighter Element Trends

In addition to examining the nature of the general trend of increasing elemental abundances with increasing metallicity,basic determinations about the nature of the various nucelosynthesis processes occurring within the precursors to these systems can be made by considering the more subtle, second order trends present within the data. O: [O/Fe] displays a weakly correlated decreasing trend with increasing [Fe/H] for both the host stars (slope=−0.16, r=−0.22) and non-host stars (slope=−0.24, r=−0.32). 39

Previous studies have hinted at the possibility of a plateau starting at [Fe/H]∼0 (Reddy et al., 2003, 2006). However, the overlap between the present sample and those previously published is not large enough to allow us to undertake any meaningful comparison. It is also worth noting that while the solar C/O ratio is 0.54 (Asplund et al., 2005), using the [C/H] values previously published by Bond et al. (2006), the C/O values of the host stars studied here range from 0.40 to 0.89. This variation has the potential to greatly impact the C chemistry of the proto-stellar disc and thus also any terrestrial planets forming in the system. This issue will be considered in more detail in Chapter 4. Mg: [Mg/Fe] can be seen to display a weakly correlated trend of decreasing with increasing [Fe/H] values producing r values of −0.32 (host stars) and −0.44 (non-host stars). The distribution observed for Mg is in good agreement (for the [Fe/H] regions in common) with those of previous studies who observed a decrease in [Mg/Fe] with increasing [Fe/H] up to [Fe/H]∼−0.10 before both distributions plateaued (Beirao˜ et al. 2005; Gilli et al. 2006; Reddy et al. 2006). Also of interest is the Mg/Si ratio of our target stars as it has the potential to greatly affect the chemical evolution of any terrestrial planets forming in the system. A high

Mg/Si ratio will result in all of the available Si forming olivine ((Fe,Mg)2SiO4) with excess Mg still being available, while a low Mg/Si ratio will result in all of the available

Mg forming enstatite (MgSiO3) with excess Si forming SiO2. Utilizing the Si abundances previously published in Bond et al. (2006), the host stars of this study were found to have Mg/Si ratios ranging from 0.46 to 1.26, resulting in potentially large variations in the nature of any terrestrial planets forming in these systems. This issue will be considered in more detail in Chapter 4. Fe Group (Cr): From Figure 2.2, it can be seen that (with the exception of one outlier) [Cr/Fe] displays no significant trend, remaining largely unchanged over all of the [Fe/H] values considered here. This is in good agreement with Reddy et al. (2003); Bensby et al. (2005) and Gilli et al. (2006). The lack of a statistically significant trend with [Fe/H] is confirmed by the low r value of −0.04 for hosts and 0.00 for non-hosts. 40

2.5.2 Heavy Element Trends

From Figure 2.2, it can be seen that all of the heavy elements display varying degrees of a weakly negative to non-existent correlation with [Fe/H]. The r values produced range from −0.22 (Y and Ba) to −0.56 (Nd) indicating that the correlation is not strong. These trends are in agreement with Allende Prieto et al. (2004) and Bensby et al. (2005) over the range of metallicity values in common with this study. We do differ slightly from some previous studies in that we observed a stronger decrease in [Y/Fe] with increasing [Fe/H] than was observed by Bensby et al. (2005) and disagree with previous studies who observed no significant trend with [Fe/H] for both [Ba/Fe] and [Nd/Fe] (Reddy et al. 2003). We do observe trends in both of these samples and the difference is attributed to the fact that we are examining a different metallicity region to that of Reddy et al. (2003). Our population is largely concentrated in the region of [Fe/H]>0, while the sample of Reddy et al. (2003) almost exclusively has values of [Fe/H]<0. However, it is worth restating that the general trends in the neutron capture elements identified here are the same as those previously identified for other solar-type stars. This implies that the host stars examined in this paper follow the same trends as other field stars but with a bias towards the high metallicity region. Additionally, the high degree of scatter observed in these samples is also in agreement with previous studies. Also of interest is the ratio of the heavy to light s-process elements as each are thought to be produced in slightly different stellar settings. The abundance of the heavy s-process element Ba to the light s-process elements Y and Zr is shown in Figure 2.3 as [heavy/light] where:

[Y/H] + [Zr/H] [heavy/light] = [Ba/H] − 2

From Figure 2.3, it can be seen that there is no dependance on metallicity, with values scattering around the solar value (0.0 by deﬁnition). This is in agreement with Reddy 41

Figure 2.3: [heavy/light] vs. [Fe/H] plots for the five heavy elements examined as part of this study. For the definition of [heavy/light], please see the text. Open squares indicate non-host stars, filled squares indicate host stars and dashed lines indicate solar values. 42 et al. (2003). Thus we support their conclusion that the neutron exposure in AGB stars is independent of the metallicity of the star itself, assuming (as Reddy et al. 2003 did) that AGB stars are the primary source of both the heavy and the light s-process elements.

2.5.3 Correlation with Planetary Parameters

Figure 2.4 shows plots of [X/H] against planetary parameters (Msin i, semi-major axis a, eccentricity and planetary period) for the 5 heavy elements considered in this study. HD164427 was omitted as its companion is believed to be a brown dwarf, not a gas giant. As can be seen visually and by the low r values (all ≤0.48 with most <0.15), no statistically signiﬁcant correlations exist between these abundances and any of the orbital parameters. This agrees with previous studies of other elements (e.g. Reid, 2002; Santos et al., 2003; Fischer and Valenti, 2005).

2.5.4 Correlation with Stellar Parameters

It is well known that the stellar atmospheric parameters (specifically Teff and log g) have the potential to drastically alter photospherically determined stellar abundances. As such, we examined the abundances presented here and as both samples produced r2 correlation coefficients less than 0.5, we concluded that there is no statistically significant trends present with either the stellar Teff or log g values.

2.6 Discussion

The host stars studied here do not signiﬁcantly deviate from previously established galactic chemical evolution trends. Instead, they can be regarded as being extensions of many of those trends into metallicities above solar. This lack of deviation from previously known trends strongly suggests that while they are more metal enriched than other stars not known to host planets, the host stars themselves have not systematically undergone any extraordinary chemical processing during their growth and evolution (in agreement 43

10 10

5 5

0 0

5 5

0 0

0.8 0.8

0.4 0.4

0 0 4000 4000

0 0

-0.6 -0.4 -0.2 0 0.2 0.4 -0.6 -0.4 -0.2 0 0.2 0.4

Figure 2.4: Orbital properties of extrasolar planetary systems vs. abundance of the heavy elements. The companion to HD164427 has been omitted as its mass makes it a likely brown dwarf. Values for planetary parameters were obtained from Butler et al. (2006). Left: s-process elements Y (triangles), Zr (squares) and Ba (crosses). Right: r- and mixed process elements Eu (triangles) and Nd (squares). From top to bottom: Msini of the planet, orbital semi-major axis of the planet, eccentricity of the orbit of the planet and period of the planet. 44 with Robinson et al. 2006). The conclusion that planetary hosts stars have essentially undergone “normal” stellar evolution may indeed suggest that planetary formation is a normal result of the star formation process. Of course, this does not exclude planet formation at lower metallicity values (as planets have been detected orbiting stars with metallicities significantly below solar) nor does it guarantee planet formation at high metallicity values. There are two primary hypotheses that have been offered to explain the observed high metallicity trend. The first is the “pollution” model which posits metal-rich material being added to the photosphere as a consequence of planetary formation (Laughlin, 2000; Gonzalez et al., 2001; Murray et al., 2001). The second explanation is commonly referred to as the primordial model and it suggests that the gas cloud from which these systems formed was metal enriched, resulting in the star itself being enriched in the same elements (Santos et al. 2001). Our conclusion that these host stars exhibit normal chemical evolution trends and that they are simply the metal-rich members of field star population lends support to the primordial model. One would expect that pollution of the stellar photosphere would produce deviations from the galactic evolutionary trends. To date no such trends have been observed. Additionally, we also observe that the abundances of the more volatile elements (such as O) increase with increasing metallicity for both the host and non-host stars. This would not be the case in the pollution model as it is likely that only the more refractory elements (such as Fe and Ni) would remain in the solid form as they migrated towards the star (and thus be deposited in the stellar photosphere) while the more volatile elements would be evaporated before they could be incorporated into the stellar photosphere. As such, we would expect to see enrichment only in the refractory elements and not the volatile elements if the pollution model is accurate. Furthermore, the fact that we observe no trends with metallicity in the orbital parameters of the remaining planets is difficult for the pollution model to explain. It is hard to imagine a situation whereby almost all planetary host stars accreted a significant amount of material during planetary formation without affecting the orbital parameters of the remaining planets. For these rea- 45 sons, we agree with previous studies (such as Santos et al. 2001, 2003, 2005 and Fischer and Valenti 2005) and support the primordial model for explaining the metal enrichment in planetary host stars. The abundances reported here also impact on terrestrial planet formation and evolution within these systems. Those with low Mg/Si ratios will have terrestrial planets dominated by enstatite (MgSiO3) (with a small amount of Mg-rich olivine also present) with other Si-based species also available (a composition similar to the Earth’s crust), while those with high Mg/Si ratios will have olivine-dominated planets with other Mg- rich species also present (a composition similar to the Earth’s mantle). Similarly, a high C/O ratio will result in planets with greatly increased C contents due to solid C being incorporated into the planet itself. While the detailed consequences of these examples have not yet been fully examined, it is conceivable that any terrestrial planets forming in these systems could differ from currently known terrestrial planets in terms of their rheology (thus possibly affecting the tectonics of such a planet) and the nature of volcanic activity, based on the varying silica contents of the magma. The full implications of such a chemical composition for the evolution of the terrestrial planets themselves is the subject of ongoing research and will be discussed in Chapter 4.

2.7 Summary

Elemental abundances for 8 elements, including 5 heavy elements produced by the r- and s-processes, have been presented for 28 planetary host stars and 90 non-host stars from the AAPS. Although the elemental abundances of the planetary host stars are only slightly different from solar values, the host stars are enriched over the non-hosts stars in all elements studied with the mean difference varying from 0.06 dex to 0.11 dex. Additionally, the trends of the abundances (both [X/H] and [X/Fe]) with [Fe/H] were considered and found to be largely in keeping with known galactic chemical evolution trends. This implies that these systems have followed normal evolutionary pathways and 46 are not signiﬁcantly or unusually altered. This leads us to conclude that not only are the abundance trends we are observing primordial in origin and represent the initial composition of the gas nebula that produced the star and its planets but that planetary formation may also be a natural companion to the evolution of stellar material.

Figure 2.5: GINGER MEGGS Dist. by Atlantic Feature Syndicate/United Feature Syndi- cate, Inc. Originally published 2/13/2008. 47

CHAPTER 3

SOLAR SYSTEM SIMULATIONS

3.1 Introduction

Before we can consider extrasolar terrestrial planets, we must first examine the terrestrial planets of our own Solar System. Terrestrial planet formation, both in terms of the dynamics and chemistry involved, is still not fully understood. Dynamically, basic planetary formation is described through the planetesimal theory (see Chambers (2004) for a detailed review). This theory sees planetary formation occurring through three main steps. Initially, dust settles into the midplane and accretes to form planetesimals, the first solid bodies of the system. This stage is followed by the collisional accretion of planetesimals to produce planetary embryos. The growth of the embryos occurs initially via runaway growth (where an increasing geometric cross-section and gravitational field allow for the accretion of an ever increasing number of planetesimals) before transitioning to oligarchic growth (where neighboring embryos grow at similar rates). Finally, as numbers of planetesimals decrease, the interaction between embryos becomes the dominant factor as they perturb each other onto crossing orbits, thus producing accretion via violent collisions. Many attempts have been made to simulate the third stage of terrestrial planet formation described above (e.g. Kominami and Ida 2002, 2004; Chambers and Wetherill 1998; Chambers 2001; Raymond et al. 2004). However, these simulations have had limited success and no simulation has been able to exactly reproduce the terrestrial planets of the Solar System in terms of their number, masses and orbital parameters. For example, direct N-body simulations have produced systems with approximately the correct number of planets yet with an orbital excitation greater than that observed for the Solar System 48

(Chambers and Wetherill, 1998; Chambers, 2001; Raymond et al., 2004). Similarly, simulations incorporating tidal torques (e.g. Kominami and Ida 2002, 2004) have produced too many planets but with more favorable excitation levels. Recent developments in modeling have occurred with the incorporation of dynamical friction, the process whereby equipartitioning of energy between low mass planetesimals and larger embryos results in reduced relative velocities for the embryos, thus increasing their probability of accreting. Dynamical friction has been shown to be a viable mechanism to reduce the dynamical excitation levels of the final planets to better agree with observed values within the Solar System (Levison et al., 2005). The highest resolution simulations incorporating dynamical friction that have been completed to date are those of O’Brien et al. (2006). These simulations represent a factor of ∼5 increase in the number of gravitationally interacting bodies compared to most other previous simulations of this type (e.g. Chambers 2001), and because of the large number of small planetesimals in the simulation, dynamical friction is more accurately treated than in previous simulations. The terrestrial planets produced by these simulations are a significantly better fit to the observed properties of the terrestrial planets of the Solar System. There are fewer planets (averaging 3 to 3.5 planets per simulation) and the dynamical excitation of the planetary systems is comparable to that of the actual terrestrial planets in the Solar System. Fur- thermore, the accretion timescales of the planets simulated are in agreement to within a factor of two with the accretion timescale for the Earth (∼10-30 Myr) as obtained from 182Hf-182W dating. Previous simulations (e.g. Chambers and Wetherill 1998; Chambers 2001; Raymond et al. 2004) produced accretion timescales larger than that of the Earth by a factor of four or more. However, more recent results from Touboul et al. (2007) imply an increased accretion timescale for the Earth, nearing a value of 60Myr. This longer timescale is consistent with the results of O’Brien et al. (2006). Thus it can be seen that the O’Brien et al. (2006) simulations have provided us with the most realistic and feasible models of terrestrial planet formation completed to date and represent a significant step towards understanding the last stage of the planetary formation. 49

However, the vast majority of terrestrial planet formation studies completed to date have been limited (at best) in their simultaneous chemical composition studies. Several previous studies have estimated the amount of potentially hydrated material that may have accreted into the terrestrial planets (e.g. Raymond et al. 2004; O’Brien et al. 2006) while O’Brien et al. (2006) also considered the delivery of siderophile-rich late veneer material needed to account for the siderophile element budget of the Earth’s mantle (Chou, 1978). However, detailed examination of the bulk chemical composition of the resulting planets as a test of dynamical simulations has never been thoroughly explored. This combination of the two approaches to produce a comprehensive model of both the chemistry and dynamics of terrestrial planet formation is essential to determine how well current numerical simulations reproduce not only the masses and dynamic state of the terrestrial planets, but also their bulk composition. The present study represents a first step towards such a model. I have derived detailed bulk elemental abundances for all of the terrestrial planets formed in the simulations of O’Brien et al. (2006). These predicted compositions are compared to the bulk compositions of the actual terrestrial planets as an examination of the chemical validity of the dynamical simulations currently being used to model terrestrial planet formation. This approach allows me to not only examine the bulk elemental composition of the final planets but to also study the compositional evolution with time. Additionally I also investigate the delivery of hydrated material, which has important implications for the development of habitable terrestrial planets, along with the composition of the “late veneer” material accreted by the planets. Finally, the amount of material accreted by the Sun as “pollution” during the terrestrial planet formation process and the resulting changes in stellar photospheric abundances are also examined. This study is the first to attempt to produce a comprehensive model of both the chemistry and dynamics of terrestrial planet formation, and represents a major improvement in modeling both the dynamical and chemical formation of terrestrial planets. 50

3.2 Simulations

3.2.1 Dynamical

The eight SyMBA sympletic N-body integrator simulations of O’Brien et al. (2006) were utilized in this study. Each simulation consists of an equal distribution of mass between Mars-mass embryos (0.0933ML, 25 per simulation) and planetesimals 1/40th the size of the embryos (0.00233ML, ∼1000 per simulation). The initial distribution of bod-

P P −3/2 P ies is constrained by the disk surface density proﬁle (r)= 0(r/1AU) with 0 = 8 gcm−2 (Chambers, 2001) and bodies initially located between 0.3 and 4.0 AU (O’Brien et al. 2006). Four simulations were run with Jupiter and Saturn in low-eccentricity, low- inclination orbits as predicted by the Nice Model (Gomes et al., 2005; Levison et al., 2005; Morbidelli et al., 2005) (hereafter termed ‘Circular Jupiter and Saturn’, CJS) and four with Jupiter and Saturn in their current, slightly eccentric orbits (hereafter termed ‘Eccentric Jupiter and Saturn’, EJS). Multiple terrestrial planets formed within 250Myr in all simulations. The general system structure is shown in Figure 3.1 and the planetary orbital properties are shown in Table 3.1. It should be noted that the formation of Mer- cury analogs is not currently possible in the simulations of O’Brien et al. (2006) as all of the embryos begin with a mass almost twice that of Mercury. Additionally, the planetesimals are not gravitationally interacting with each other, preventing formation of a small ‘Mercury’ via planetesimal accretion. Finally, perfect accretion is assumed throughout the simulations, preventing giant impact events from striping off the crust and mantle of a differentiated proto-Mercury, a scenario that widely suggested to explain Mercury’s current high density. As the simulations, results and their implications have already been discussed in great detail by previous publications, the reader is referred to O’Brien et al. (2006) for further discussion. 51

Final Planetary Systems

CJS1 1.14 0.81 0.78 CJS2 0.44 0.35 1.20 0.79 CJS3 0.76 1.57 0.54 CJS4 1.30 1.42

EJS1 0.59 0.89 0.47 EJS2 0.35 0.74 0.81 EJS3 0.67 0.95 0.45 0.10 EJS4 0.220.76 0.14 0.98

SS 0.05 0.81 1.00 0.11

0 0.5 1 1.5 2 2.5 Semi-Major Axis

Figure 3.1: Schematic of the results of the simulations of O’Brien et al. (2006). The horizontal lines indicate the variation between aphelion and perihelion. The vertical lines indicate variation in distance from the midplane due to the inclination of the planet. Nu- merical values represent the mass of the planet in Earth masses. The Solar System (SS) is shown for comparison. 52

Table 3.1: Properties of simulated terrestrial planets produced in the Solar System simulations of O’Brien et al. (2006). Numbering starts at 4 and increases with increasing distance from the Sun. CJS denotes the results of the circular Jupiter and Saturn simulations while EJS indicates the results of the eccentric Jupiter and Saturn simulations.

Planet M a e i (ML) (AU) (◦) CJS1−4 1.15 0.63 0.05 4.36 CJS1−5 0.81 1.21 0.06 4.85 CJS1−6 0.78 1.69 0.04 2.06

CJS2−4 0.44 0.55 0.05 2.61 CJS2−5 0.36 0.69 0.06 3.54 CJS2−6 1.20 1.10 0.02 0.38 CJS2−7 0.80 1.88 0.04 1.83

CJS3−4 0.77 0.62 0.05 1.54 CJS3−5 1.57 1.14 0.06 2.11 CJS3−6 0.55 2.09 0.06 2.07

CJS4−4 1.31 0.66 0.10 0.60 CJS4−5 1.43 1.54 0.08 3.64

EJS1−4 0.59 0.56 0.03 1.69 EJS1−5 0.89 0.84 0.03 1.24 EJS1−6 0.48 1.29 0.03 1.55

EJS2−4 0.35 0.50 0.04 1.61 EJS2−5 0.74 0.75 0.02 1.07 EJS2−6 0.82 1.17 0.02 0.71 EJS2−7 0.10 3.19 0.24 14.95

EJS3−4 0.68 0.58 0.04 2.12 EJS3−5 0.45 1.52 0.04 3.20 EJS3−6 0.96 0.96 0.02 2.10 EJS3−7 0.11 2.08 0.13 7.83

EJS4−4 0.77 0.68 0.03 1.48 EJS4−5 0.23 0.48 0.06 2.29 EJS4−6 0.14 1.07 0.09 3.82 EJS4−7 0.99 1.33 0.02 1.59 53

3.2.2 Chemical

The chemical composition of material within the disk is assumed to be determined by equilibrium condensation within the primordial solar nebula. This assumption is a reasonable starting point as numerous analyses of primitive chondritic material have shown that their bulk elemental abundances are smooth functions of their equilibrium condensation temperature, as determined for a solar composition gas at low pressure (Davis, 2006). Several possible explanations for this pattern have been proposed, with the most widely accepted being that the chemistry of their parent planetesimals was established by midplane temperature and pressure profiles (Cassen, 2001). Alternatively, the sluggish nature of reaction kinetics within the cooler regions of the outer nebula may also be a possible cause for the observed trend. Although the exact cause of the primitive chondritic pattern is still unclear, I am able to utilize the fact that equilibrium condensation temperature has been shown to be an excellent proxy for determining the bulk elemental abundance of rock-forming elements within the early Solar System. Furthermore, obser- vational evidence of these equilibrium compositions is still seen today, preserved as the thermal stratification of the asteroid belt (Gradie and Tedesco, 1982). The mineralogy of primitive chondritic meteorites is similar in most respects to that predicted by equilibrium condensation (Ebel, 2006). Thus the assumption of equilibrium composition is a valid starting point for this type of study. This assumption in turn implies that the primary controls on embryo and planetesimal compositions are the radial midplane temperature and pressure gradients encountered by the nebular material. In order to determine the equilibrium composition of the solid material, I utilized the commercial software package HSC Chemistry (v. 5.1). HSC Chemistry determines the equilibrium chemical composition of a system by iteratively minimizing the system’s Gibbs free energy, using the GIBBS equilibrium solver described in White et al. (1958). This software has been successfully used in recent studies of solar nebula chemistry (Pasek et al., 2005) and supernova stellar outflows, producing compositions that corre- 54 late with observed mineralogy in presolar interplanetary dust particles (Messenger et al., 2005). As such, I feel confident in applying it to this current study. Several limitations have been encountered with the program, primarily the limited number of elements and species possible to incorporate in a simulation and its inability to handle species in solid solutions such as pyroxene and olivine, two of the major rock-forming minerals on the terrestrial planets. End member species can and have been considered but the non-ideal solid solution interaction between them is not currently possible. These limitations will not significantly affect the final conclusions of this current study but will be the subject of future work. The list of solid and gaseous species included in the HSC Chemistry calculations are shown in Table 3.2. Our present simulations incorporate 14 major rock-forming elements (C, N, O, Na, Mg, Al, Si, P, S, Ca, Ti, Cr, Fe and Ni), along with H and He, which dominate the partial pressures of protoplanetary disks. Current solar photospheric abundances are utilized as a proxy for the initial composition of the solar nebula. All inputs are assumed to initially be in their gaseous, elemental forms and are then cooled and allowed to react. Solar elemental abundances for all elements were taken from Asplund et al. (2005) and are shown in Table 3.3. Although there has been some recent discussion about revising the Solar C and O abundances (e.g. Socas-Navarro and Norton, 2007), I have adopted the most recent and widely accepted values for this current study. The 50% condensation temperature of each element in a gas with a Solar photosphere composition as obtained from HSC Chemistry is shown in Table 3.4. It should be noted that although O is present in high temperature condensates and silicate species, it does not obtain 50% condensation until water ice condenses within the system, hence its low 50% condensation temperature. Additionally, the 50% condensation temperatures from the models of Lodders (2003) are also provided. The average difference between this study and that of Lodders (2003) is just +11.75K (defined as Tpresent study -TLodders), with the largest difference being 83K for P. The excellent agreement between the temperatures obtained in this study and those of Lodders (2003) implies that the chemical models utilized in this study are accurately 55

Table 3.2: Chemical species included in the equilibrium calculations of HSC Chemistry.

Gaseous Species Al CrO MgOH PN AlH CrOH MgS PO AlO CrS N PS Al2O Fe N2 S AlOH FeH NH3 S2 AlS FeO NO SN C FeOH NS SO CH4 FeS Na SO2 CN H Na2 Si CO H2 NaH SiC CO2 HCN NaO SiH CP HCO NaOH SiN CS H2O Ni SiO Ca HPO NiH SiP CaH HS NiO SiP2 CaO H2S NiOH SiS CaOH He NiS Ti CaS Mg O TiN Cr MgH O2 TiO CrH MgN P TiO2 CrN MgO PH TiS

Solid Species Al2O3 FeSiO3 CaAl2Si2O8 C CaAl12O19 Fe3P NaAlSi3O8 SiC Ti2O3 Fe3C Cr2FeO4 TiC CaTiO3 Fe Ca3(PO4)2 TiN Ca2Al2SiO7 Ni FeS AlN MgAl2O4 P Fe3O4 CaS Mg2SiO4 Si Mg3Si2O5(OH)4 MgS MgSiO3 Cr H2O Fe2SiO4 CaMgSi2O6 56

Table 3.3: HSC Chemistry input values for Solar System Simulations. All values are in moles. All species are assumed to initially be in their elemental and gaseous forma with no other species present within the system. Element Abundnace H 1.00 ×1012 He 8.51 ×1010 C 2.45 ×108 N 6.03 ×107 O 4.57 ×108 Na 1.48 ×106 Mg 3.39 ×107 Al 2.34 ×106 Si 3.24 ×107 P 2.29 ×105 S 1.38 ×107 Ca 2.04 ×106 Ti 7.94 ×104 Cr 4.37 ×105 Fe 2.82 ×107 Ni 1.70 ×106

reproducing the initial equilibrium composition of the solar disk. In order to relate the chemical abundances to a spatial location within the original disk, I used midplane pressure and temperature values obtained from the axisymetric α viscosity disk model of Hersant et al. (2001). The Hersant et al. (2001) model is a two- dimensional time-dependent turbulent accretion disk model incorporating vertical disk structure, turbulent pressure and self-gravity. As for other standard disk models, a Kep- lerian rotation law, hydrostatic equilibrium and energy balance between viscous heating and cooling due to radiative losses are assumed (Hersant et al., 2001). The effects of irradiation from the central star are neglected, as are other disk features such as inner holes and shadow zones. To deﬁne a “nominal” disk model, Hersant et al. (2001) restrict the initial disk mass to be 0.3MJ or less, in accordance with the gravitational instabil- ity models of Shu et al. (1990). Similarly, angular momentum is assumed to have been 57

−4 Table 3.4: T50% condensation for a gas with Solar photosphere composition and at P = 10 bar. Initial phase for each element is also listed. Values are provided from the models of Lodders (2003) for comparison.

Element T50% condensation (K) Initial Phase This Study Lodders (2003) Species Formula

Al 1639 1665 Hibonite CaAl12O19 C <150 40 Methane Clathrate CH4.7H2O Ca 1527 1505 Hibonite CaAl12O19 Cr 1301 1291 Metallic Chromium Cr Fe 1339 1328 Metallic Iron Fe Mg 1339 1327 Spinel MgAl2O4 Na 941 953 Albite NaAlSi3O8 Ni 1351 1348 Metallic Nickel Ni O 180 179 Hibonite CaAl12O19 P 1309 1226 Schreibersite Fe3P S 658 655 Troilite FeS Si 1329 1302 Gehlenite Ca2Al2SiO7 Ti 1580 1573 Perovskite CaTiO3 58 transferred via turbulence to the formation region of the outermost giant planet (Neptune) in 2.5×105 years so as to allow for solid/gas decoupling and core formation to occur. Initial disk temperature is taken to be greater than 1000K within 3AU from the Sun in order to produce the crystalline silicate distributions observed in meteorites and comets by Bockelee-Morvan´ et al. (2002). Finally, deuterium abundances produced by the model must be in agreement with the values reported for LL3 meteorites at 3AU (the predicted location of formation) and cometary values obtained for Halley, Hyakutake and Hale- Bopp. Based on these restrictions, Hersant et al. (2001) find a “nominal” disk model is produced by a stellar mass accretion rate of 5×10−6MJyr−1, an initial disk radius of 17 AU and an α value of 0.009. In the current study, I limit the midplane conditions to be those produced by the “nominal” model. Previous work has successfully applied this model to constrain the bulk composition of Jupiter (e.g. Pasek et al. 2005). Pressure and temperature values were determined with an average radial separation of 0.03AU throughout the study region. The midplane temperature and the pressure, and thus also the equilibrium composition of solids present within the disk, changes with time as the disk evolves. In order to capture this effect in our simulations, I determined an ensemble of predicted planetary compositions constrained by predicted disk conditions at multiple time periods. At the earliest, I use the temperature and pressure profiles obtained for disk conditions at t = 2.5×105yr, where conditions were first determined to be suitable for solids to be present across the entire radial region being modelled in the dynamic simulations. I end with disk conditions at t = 3×106yr, the average lifetime for the protoplanetary gas disks (Haisch et al., 2001). Five cases between these end points are considered (for disk conditions at t = 5×105yr, 1×106yr, 1.5×106yr, 2×106yr, 2.5×106yr). Thus I determine the planetesimal and embryo compositions for a total of seven cases, covering the entire range of times during which embryos and planetesimals could potentially form, either early or late in the lifetime of the disk. Midplane P and T profiles for each set of simulations are shown in Figure 3.2. Note that the timescale used for the evolution of the disk conditions is not 59 coupled with the timescale of the dynamical simulations. Rather, chemical compositions are simply determined for a range of disk mid-plane conditions and no time-variation in disk conditions over the duration of the dynamical simulation is incorporated.

3.2.3 Combining Dynamics and Chemistry

In order to combine these two different modeling approaches, I assume that each planetary embryo and planetesimal of the dynamical models retains the chemical composition in equilibrium with the nebula in the region that it first formed. The bulk compositions of the final planets are simply the sum of each object they accrete. By tracing the origin of each embryo and planetesimal incorporated into the final planets of the O’Brien et al. (2006) dynamical simulations, and calculating the chemical composition of those bodies based on their original locations, I constrain the bulk composition of the final terrestrial planets. This procedure was completed for all planets using a small perl script at each of the seven different time steps simulated within the chemical models.

3.2.4 Stellar Pollution

The extent of stellar pollution produced by terrestrial planet formation was measured by determining the amount and composition of material accreted by the Sun during the formation process and the resulting photospheric elemental abundance changes such an addition would produce. Any solid material migrating to within 0.1AU from the Sun is assumed to have accreted onto the Solar photosphere through gravitational attraction. This material is then assumed to have been uniformly mixed throughout the solar photosphere and convective zone. Granulation within the photosphere and gravitational settling and turbulence within the convective zone are neglected from this present study, primarily as there is still signiﬁcant debate about and very little consensus on the exact nature of these processes and their effects on speciﬁc elements. However, it is expected that each of these processes would reduce the amount of accreted material present within the upper layers of the Sun, thus reducing the photospheric elemental abundance. As such, the 60

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure 3.2: Radial midplane pressure and temperature profiles for the Solar nebula obtained from Hersant et al. (2001) for “nominal” conditions. “Nominal” refers to the con- ˙ −6 J −1 ditions of M = 5×10 M yr ,Rinital = 17 AU and α = 0.009. Top Panel: Pressure profiles for each of the 7 disk conditions considered. Bottom Panel: Temperature profiles for each of the 7 disk conditions considered. Times shown indicate evolutionary age of the disk, not the dynamical models. 61 resulting abundance changes determined in this study represent the maximum possible values that I would expect to observe in the Solar spectrum. Additionally, at earlier times, the mass of the solar convective zone would have been greater than the current value of 0.03 MJ, again resulting in a smaller change in photospheric abundance than predicted by this current study. The mass of each element accreted by the Sun was determined in the same way as described in section 3.2.3 for terrestrial planets. The resulting photospheric abundance is given by:

  f  X  [X/H] = log J (3.1) fX,

where [X/H] is the resulting abundance of element X after accretion of terrestrial planet material, fX is the mass abundance of element X in the Solar photosphere after J accretion and fX, is the initial mass abundance of element X in the Sun before accretion J (from Murray et al. 2001). fX, values were obtained by utilizing the solar abundances of Asplund et al. (2005) and a current solar convective zone mass of 0.03MJ (Murray et al., 2001). The present approach only addresses pollution by the direct accretion of planetesimals and embryos during the current simulations, yet pollution may also occur both before planetesimal accretion (and thus the simulation) begins and through the accretion of dust produced by the impact and collisional events during the formation process. Pollution by these processes is believed to be small, resulting in at most a factor of two increase in the amount of material added to the solar photosphere. 62

3.3 Results

3.3.1 Abundance Trends

The bulk elemental abundances for all terrestrial planets for each set of disk conditions examined are shown in Table B.1 (as elemental wt% of the planet). Additionally, the Mg/Si value for each planet is also shown in Table B.1. The average simulated terrestrial planet abundances are in reasonable agreement with the abundances of their actual Solar System counterpart. The ensemble-averaged (i.e. averaged over all disk conditions considered) mean elemental abundances for all of the terrestrial planets are listed in Table B.2. The simulated abundances are all within 10 wt% of the observed terrestrial planet values of Morgan and Anders (1980) (Venus), Kargel and Lewis (1993)(Earth) and Lodders and Fegley (1997) (Mars) for the key planet-building elements (Mg, Si, O, Fe) and display smaller deviations for the other elements (within 6 wt% for S, within 2 wt% for all other elements). However, it can be seen from Table B.1 that the bulk elemental abundances of the final terrestrial planets vary significantly with disk conditions, especially for certain elements such as O, Al and Ca. The difference in planetary composition between the chemical simulations undertaken for disk conditions at t = 2.5×105yr and t = 3×106yr is shown in Table B.3. As expected, the relative abundances of the most refractory elements (Al, Ca, Ti) are observed to decrease for disk conditions at later times while the more volatile elements (O, Na, H) all increase in abundance. The largest variation occurs in the O abundance, with simulations for disk conditions at t = 3×106yr producing terrestrial planets with up to 30.79 wt% more O than identical simulations for disk conditions at t = 2.5×105yr. These compositional variations can be directly attributed to the changing conditions of the disk itself. As the temperature and pressure at the mid-plane decrease with time, the equilibrium composition at a specific radius also changes. Thus the assumed composition of the planetesimals and embryos also change, in turn producing variations in the final planetary abundances. Consequently, as time progresses and the disk cools, the planetary 63 composition can be seen to decrease in abundance of the most refractory elements (Al, Ca) and increase in the more volatile elements (O, H) as volatile, hydrated species become stable throughout an increasingly larger fraction of the disk. A more accurate picture of the predicted planetary abundances is obtained by examining the abundances produced for each specific set of disk conditions simulated. Figure 3.3 shows the Si and planet-normalized bulk elemental composition for the planets produced in the CJS-1 and EJS-1 simulations for all seven disk conditions. The simulated planets were normalized to Venus, Earth and Mars on the basis of their orbital properties, primarily their semimajor axis. Normalized abundances were obtained for each simulated planet via the relation:

(X/Si) Normalized abundance for element X = simulated (3.2) (X/Si)observed

where (X/Si)simulated are the abundances of element X and Si for the planet produced by the present simulations while (X/Si)observed are the previously published abundances of X and Si for Venus, Earth or Mars. Reference Solar System planetary abundances were taken from Morgan and Anders (1980) (Venus), Kargel and Lewis (1993)(Earth) and Lodders and Fegley (1997) (Mars). All three suites of planetary reference values are themselves based (to some degree) on modeling. Bulk Earth values were obtained by extrapolation of Best Bulk Silicate Earth (Best BSE or BBSE) abundance values (as obtained from previously published studies of mantle xenoliths, pyrolite and basaltic elemental ratios), the observed volatility trend for volatile lithophile elements and the known physical and seismic properties of the interior of the Earth (Kargel and Lewis, 1993). This approach, although necessary, does introduce uncertainties especially when calculating the composition of the core as partioning ratios are not well known for all elements. Furthermore, the use of a volatility trend requires knowledge of the condensation temperature of the elements. While such 64 information is well known for several species (such as Fe, Mg and Si with errors on the order of 3% of the condensation temperature), larger degrees of uncertainty remain for more volatile elements. This uncertainty produces errors in the predicted condensation temperatures of ±14% of the condensation temperature itself and results in final elemental abundance errors of 5-10% of the actual abundance for most elements, reducing to less than 1% for Fe, Mg and Si. Of the elements considered in the current study, those with the greatest overall uncertainty are O, S and C. Although not quantified by Kargel and Lewis (1993), the uncertainty in O is derived from the fact that the O abundance was obtained based on the valence states of cations. This approach is predicated on the assumption of an abundance of 15% of all Fe present is in the form Fe3+. As such, any variation within the assumed valence state distribution will induce a variation in the O abundance of Earth. C and S abundances similarly suffer from large uncertainties (∼25% of the given abundance) due to variations in the published BSE values, variations in the predicted condensation temperatures and the partioning coefficients for each of these elements. Bulk Mars abundances were taken from Lodders and Fegley (1997) and are based on the combination of known meteorite material (H, CV and CI meteoritic material) required to reproduce the oxygen isotopic abundances observed for the SNC meteorites. Variations in the elemental compositions within each class of meteorite are thus the main source of uncertainty in the current estimates. Lodders and Fegley (1997) estimate their final abundance errors to be ±10% for all elements studied. Finally, the marked lack of data regarding the elemental abundances of Venus resulted in the necessary use of the purely theoretical bulk planetary abundances of Morgan and Anders (1980) for this study. Morgan and Anders (1980) adopt a similar approach to the one used here and assume that the composition of solid material initially present within the system is controlled by equilibrium condensation. The exact predicted composition was obtained from the Ganapathy-Anders 7-component model (Morgan et al., 1978). This model determines the composition of a body based on the amount of early condensate (i.e. first solids present within the disk), metal, silicate, troilite, FeO, MnO and 65

Cr2O3 believed to be incorporated into a final planetary body. Bulk elemental abundances are thus inferred from general cosmic proportions as compared to an “index element” for each component. As such, general and solar elemental ratios are assumed to be homogenous and remain constant for the entire planetary system. For the simulations of Morgan and Anders (1980), the limited information available on the index elements for Venus was supplemented by geochemical restrictions from the Earth, Moon and chondritic meteorites to obtain predicted bulk planet abundances. As such, large uncertainties exist in all elements studied. As quantified uncertainty factors are not provided by Morgan and Anders (1980), it is impossible to gauge the errors on their predicted abundances. How- ever, they are believed to be considerable, resulting in elemental abundances for Venus that should be taken as a guide only. From Figure 3.3 it can be seen that although no single simulation exactly reproduced the compositions of the terrestrial planets, the time when the adopted disk conditions produced simulated compositions that most closely resemble those of the terrestrial planets (the ‘best fit’ time) is 5×105yr, based on the comparatively small deviations produced in the bulk elemental abundances. For disk conditions at later times, significantly larger enrichments can be seen in O, Na and S for all planets. Although not shown here, disk conditions at this same time also produces the best agreement between simulated and observed abundances for the other six simulations studied and as such future discussions will focus on the compositions produced by the disk conditions at this time. This result does not imply that the material which formed the terrestrial planets actually condensed out of the Solar nebula exclusively at 5×105yr. Rather it is simply the time at which the disk conditions and resultant snapshot of the chemistry of the Solar disk (as used here) best reproduced the expected abundances. Similar plots for the other six simulations studied are displayed in Figures B.2 - B.7. Numerically, for simulations based upon disk conditions at 5×105yr, abundances agree with Solar System values to within 1 wt% for all elements studied except for Mg (up to 2 wt% deviation), Fe and O (up to 2.5 wt% deviation) and S (up to 5 wt% deviation). In 66

10 10 CJS1-4 (Venus) EJS1-4 (Venus)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS1-5 (Earth) EJS1-5 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS1-6 (Mars) EJS1-6 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure 3.3: Normalized abundances for CJS1 and EJS1 simulated terrestrial planets. The terrestrial planet each simulation is normalized to is shown in parentheses. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. Left: CJS1 terrestrial planets. Right: EJS1 terrestrial planets. Reference Solar System planetary abundances were taken from Morgan and Anders (1980) (Venus), Kargel and Lewis (1993)(Earth) and Lodders and Fegley (1997) (Mars). 67 terms of normalized values, the predicted abundances are within a factor of 0.5 for almost all elements studied. Enrichments can be seen in Na (up to 4× the expected value) and S (up to 8× the expected value) and are discussed below. These abundances represent an excellent agreement with observed Solar System values, indicating that these terrestrial planet formation models are producing terrestrial planets with both orbital properties and bulk elemental abundances comparable to those of the terrestrial planets of the Solar System. Despite the excellent agreement between the observed and predicted abundances, the planet-normalized spider plots of Figure 3.3 display several systematic deviations from the expected terrestrial planet abundances. Specifically, considerable deviations can be seen in P (Venus analogs only), Na and S for all simulations. This is partly due to the uncertainties in the expected planetary abundances for these elements as was previously discussed. Additionally, as P, Na and S are some of the most volatile in this system, it is believed that the observed enrichments are a result of the fact that I am currently not considering volatile loss during the accretion process. Both the dynamical and chemical simulations assume that in any given impact, perfect merging of the two bodies occurs and all mass is retained. However, such violent impacts are known to both melt and eject a considerable portion of both the target body and the impactor (such as in the moon forming impact (Hartmann and Davis, 1975; Cameron and Ward, 1976)). This has the potential to drastically alter the bulk elemental abundances of a planet and it is expected that a significant fraction of the most volatile species in the system will be lost during the accretionary process, thus reducing the observed enrichments. This effect has not been incorporated into our simulations, resulting in volatile enriched final planetary bodies. The results of a first order approach to include the abundance effects of volatile loss are discussed in Section 3.3.5. Some degree of radial compositional variation is captured by the current models. Sim- ulated planets deemed to be Venus and Earth analogs were found to produce equivalent normalized abundances for the major rock forming elements when normalized to the el- 68 emental abundances of Venus and Earth. Differences in the quality of the fit occur for P, Na and S and are due to uncertainties in the abundances of Venus. The Mars analogs are found to have a slightly better fit when normalized to the elemental abundances of Mars, as compared to those of Venus and Earth. This difference is minor with Earth abundances arguably producing an equivalent normalization. However, given the small normalized radial compositional differences observed for Venus, Earth and Mars, it is not surpris- ing that the approach adopted in this study has been unable to complectly reproduce the expected radial variations. In addition to matching the bulk elemental abundances, simulations such as these should also reproduce the key geochemical ratios for the terrestrial planets. Here I have considered the values of Mg/Si, Al/Si and Ca/Si. All three of these values differ by chemically significant amounts from the expected planetary values for all planets produced. For simulations run with disk conditions at times of 5×105 years and beyond, the planetary ratios are identical to the solar input ratios. This effect is produced by the fact that for disk conditions at t = 5×105 years, all the Mg, Al, Ca and Si has condensed out over the study region. This results in the solid species (and thus also the planets produced) possessing the same chemical ratios as present in the initial solar nebula. For example, for Hersant et al. (2001) disk conditions at t = 5×105 years, Mg/Si, Al/Si and Ca/Si ratios in the solid material reach Solar values by 0.5AU, interior to the primary feeding zone for the formation of Earth and Mars. Unfortunately, both the Mg/Si and Al/Si values are lower than is observed for the Earth but above current Martian values (see Figure 3.4) while the Ca/Si values are above those of both Earth and Mars (but in agreement for Venus). Thus for the majority of disk conditions presently considered I am producing planets that have compositional ratios in between those of Earth and Mars and are slightly enriched in Ca. Variations in these elemental ratios from solar ratios are produced in simulations based on disk conditions at 2.5×105 years (see Figures 3.4 and 3.5). At this time, the inner disk (within approximately 1AU) is dominated by Al and Ca rich species such as spinel

(MgAl2O4) and gehlenite (Ca2Al2SiO7), resulting in planets with high Al/Si and Ca/Si 69 values and low Mg/Si values as can be seen in Figures 3.4 and 3.5. Thus it can be seen that there is a temporal variation in the chemical ratios of the planets produced. This further supports our choice of disk conditions at t = 5×105 years as producing the “best fit” abundances as although the ratios discussed here do not precisely agree with the observed values, the deviation is significantly smaller than for earlier disk conditions. The fact that the current chemical ratios do not agree with observed planetary values implies two possible solutions. The first is that the dynamical simulations are not forming planets from material sourced from the correct region of the disk. There is a small radial region within the disk where the values of Mg/Si, Al/Si and Ca/Si are in agreement with those of Earth. This region occurs between temperatures of 1352 and 1305K, corresponding to a radial location of 0.61 to 0.68AU for disk conditions at t = 2.5×105 years and 0.12 to 0.13AU for disk conditions at t = 3×106 years for the disk model of Hersant et al. (2001). Thus one possible solution for the current difference may be that the dynamical models need to form Earth from material located within this region. However, it is ex- ceedingly difficult to imagine how such a formation process would occur as it requires the movement of a large amount of material over a relatively long distance through the disk. Current formation simulations do not produce this degree of radial mixing and it is questionable whether sufficient material would be located within the region, making it unlikely that such a scenario is feasible. On the other hand, the problem may lie in the disk models I am currently using to obtain radial P and T profiles. Although there are variations in disk models (see Boss (1998) for a review), I feel confident in our current model as similar P and T profiles have been produced by other studies (Cassen, 2001) and the current model has been successfully utilized by other chemical studies (e.g. Pasek et al. 2005). Furthermore, planetary compositions produced at these temperatures would not be in agreement with observed bulk planetary elemental abundances. Finally, observations have suggested that the temperature at 1AU in the disk of young stellar objects is less than 400K (Beckwith et al., 1990), again making it unlikely that it would be possible to produce the disk conditions required to produce the necessary Mg/Si, Al/Si and Ca/Si 70

1.4 Earth fractionation line

1.2

1.0

0.8 ratio) ght 0.6

0.4 Mg/Si (wei

0.2 Mars fractionation line 0.0 0.0 0.2 0.4 0.6 0.8 1.0

Al/Si (weight ratio)

Figure 3.4: Al/Si v. Mg/Si for all simulated terrestrial planets. Black circles indicate values for disk conditions at t = 2.5×105 years while red circles indicate values for disk conditions at t = 5×105 years. Values at all other times are concentrated at the 5×105 years values and are not shown for clarity. Earth values are shown in green and are taken from Kargel and Lewis (1993) and McDonough and Sun (1995). Martian values are shown in pink and are taken from Lodders and Fegley (1997). 71

1.1

1.0

0.9

0.8

0.7 Mg/Si (weight ratio)

0.6

0.5 0.0 0.2 0.4 0.6 0.8 1.0 1.2 Ca/Si (weight ratio)

Figure 3.5: Ca/Si v. Mg/Si for all simulated terrestrial planets. Black circles indicate values for disk conditions at t = 2.5×105 years while red circles indicate values for disk conditions at t = 5×105 years. Values at all other times are concentrated at the 5×105 years values and are not shown for clarity. Earth values are shown in green and are taken from Kargel and Lewis (1993) and McDonough and Sun (1995). Martian values are shown in pink and are taken from Lodders and Fegley (1997).Venus values are shown in light blue and are taken from Morgan and Anders (1980). 72 values at 1AU. The second possibility is that the composition of solids within the disk are controlled by disequilibrium condensation. Disequilibrium condensation refers to the chemical model in which once a solid has condensed, it is removed from the system and may no longer interact with the remaining gas. Such a process may occur if condensation occurs rapidly enough to for large bodies to grow quickly and thus shield the interiors from further equilibrium reactions (Cowley, 1995). Under such conditions, secondary condensates would have Mg/Si values above average (Sears, 2004), thus possibly increasing the Mg/Si value to the required Earth value over the primary feeding zone in the O’Brien et al. (2006) simulations. Additionally, observations of protoplanetary disks have shown that much of the chemistry within the the disk itself is in disequilibrium (Bergin et al., 2007), further supporting our suggestion of disequilibrium chemistry as the controlling chemical factor. The apparent role of disequilibrium processes does not invalidate our initial assumption of equilibrium controlled abundances. The assumption of equilibrium is acceptable for determining bulk elemental abundances but ﬁner details of planetary composition will need to incorporate a greater interaction between a variety of equilib- riums and disequilibrium processes. Furthermore, an equilibrium-based predictions for simulations of this type need to be calculated initially to provide a baseline for future disequilibrium studies. The oxidation state of the planets is also of interest. As I am currently calculating predicted bulk elemental abundances and not detailed mineralogies for the simulated planets, estimates of the bulk oxidation state of the planet are obtained by calculating the oxidation state of the embryos and planetesimals before accretion occurred based on their equilibrium compositions. The resultant oxidation states are shown in Figure 3.6. The current simulations are producing very reduced planetary compositions for the earliest set of disk conditions studied, containing large amounts of metallic Fe and little FeO. For disk conditions at later times, the compositions become more oxidized, evolving through the oxidation state of the H-type meteorites to ﬁnish closer to the redox state of the CR mete- 73

orites, primarily due to the increased amount of magnetite (Fe3O4) and fayalite (Fe2SiO4) accreted by the planets. For midplane conditions at t = 2.5×105 years, fayalite is only present in significant amounts beyond 1.07AU while magnetite is only present beyond 4.82AU, beyond the dynamical simulation range. For conditions at t = 3×106 years, fayalite is present beyond 0.2AU with magnetite now present beyond 0.95AU. Thus for later disk conditions, the feeding zones of the terrestrial planets (and thus the planets themselves) are more oxidized. In addition to the observed temporal variations, migratory processes within the disk can also act to alter the oxidation state of the solid material. For example, the redistribution of water within the inner 5AU of the disk over time via diffusion and advection has been found to significantly alter the oxidation state of the disk itself (Pasek et al., 2005; Cyr et al., 1999), producing both reducing and oxidizing regions of various widths and locations (depending on the initial conditions). As a result, the redox state of the solid material should be drastically altered as the disk evolves. It is also worth mentioning that the extremely reduced nature of the solid material may also be increased to some extent by our present inability to simulate the olivine and pyroxene solid solutions. This acts to limit the amount of Fe that can be oxidized and incorporated into silicate species, thus producing more reduced compositions. The effect of this approach is believed to be relatively minor. Comparison of simulated oxidation values to those of the terrestrial planets is difficult, either due to lack of direct information (Venus) or a varying redox state (Earth and Mars). Earth has a reduced metallic core, slightly oxidized mantle and extremely oxidizing crust while Mars may be composed of a relatively reduced core and highly oxidized crust. Furthermore, it is believed that the redox state has varied through time (e.g. Galimov 2005). For example, in order for core formation to occur, a reduced mantle composition is required. However, the mantle must have achieved its present oxidation level relatively early in the Earths history as little temporal variation is observed in the oxidation state of basalts younger than ∼3.5 - 3.9 billion years (Delano, 2001). Thus the generally reduced 74

EL K CR CV

Figure 3.6: Oxidation state plot for CJS1 and EJS1 simulated planetary abundances. Squares indicate CJS1 values. Circles indicate EJS1 values. Values are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. Approximate regions for several meteorite groups are also shown. 75 nature of the planets produced in the current simulations is in excellent agreement with the hypothesis of an initially reduced Earth and Mars before planetary processes altered the redox state of their upper layers. The exact oxidation state of the primitive planets and thus the best fit conditions implied by it are still unclear. Although differences in the orbital and dynamical properties of the final planets are produced by the two different types of simulations considered here (O’Brien et al., 2006), the differences between the CJS and EJS simulations in terms of both their bulk elemental abundances and their geochemical ratios for rock forming elements are negligible. Neglecting the small volatile rich planet formed in the asteroid belt in simulation EJS3, comparable bulk compositional trends can be seen in both the CJS and EJS simulations. This indicates that in the case of terrestrial planet formation within the Solar System, the bulk composition of the final terrestrial planets formed is not highly dependent on the orbital properties of Jupiter and Saturn. Once again, this is a result of the fact that the majority of planet forming elements have fully condensed out of the nebula by 1253K, corresponding to midplane radii within 0.75AU from the host star for the models of Her- sant et al. (2001) and do not change greatly in their relative weight percentage values over the remainder of the simulation region. As a result the majority of the solid mass within the present simulations has solar elemental abundances. This is consistent with the idea that the equilibrium chemical composition of the Solar nebula is not highly zoned in that it does not contain radially narrow zones of vastly different compositions. Instead, the vast majority of the solar disk is dominated by material composed of pyroxene (MgSiO3), olivine (Mg2SiO4) and metals (Fe and Ni). Figure 3.7 shows the weight percentage (in solid material) of the key planet building elements O, Fe, Mg and Si, in addition to the mass distribution of the dynamical simulations. It can be seen that the majority of mass is expected to have similar relative elemental abundances. The small variations observed in the weight percent vales are caused by increases in the solid portion of disk mass generated by the condensation of FeS at 1AU and serpentine (Mg3Si2O5(OH)4) at 3.5AU. Thus small variations in the feeding zones produced by different orbital properties of Jupiter 76 and Saturn within the CJS and EJS simulations are unlikely to produce major differences in final bulk composition. Difference between the two types of simulation emerge when I consider the delivery of hydrous species (Section 3.3.4).

3.3.2 Variations with Time

Figure 3.8 shows the variation with time for weight percentage values of several key planet building elements in the planets produced by the CJS1 and EJS1 simulations. It can be seen from Figure 3.8 that very little variation in composition (<5 wt%) occurs during the formation process with the bodies obtaining their ﬁnal compositions relatively early during accretion and displaying only minor deviations over time. The same trend is also observed for the six other simulations not shown. This implies that the planets within these simulations formed homogenously (i.e. from material with similar composition to the ﬁnal planet). If formation did occur in this manner, then the presence of a 800- 1000km deep magma ocean would be required during core formation in order to produce the siderophile abundances observed in the crust and upper mantle of the Earth (Drake, 2000). However, it is likely that the homogenous accretion observed here is due to the ‘snapshot’ approach I have taken when determining the composition of solid material within the disk. In these simulations, I have considered the composition determined by disk conditions at just seven discrete times. In reality, it is quite likely that the composition of solid material will change over time as it undergoes migration and experiences other disk and stellar processes (such as the redistribution of water (Cyr et al., 1999)). Such changes are not captured in our current approach. As such, more detailed chemical models incorporating temporal variations in the composition of solid material are required in order to test the homogenous accretion produced here. 77

Si Mg

Figure 3.7: Distribution of solid mass and its relative composition within the Solar disk at 5×105 years. Top: Composition (in wt%) for the solid material within the disk for O (black), Fe, (red), Mg (green), Si (blue) and Al (yellow). Bottom: Initial distribution of mass within the dynamical simulations of O’Brien et al. (2006). 78

40 40 CJS1-4 EJS1-4 O O 30 30 Fe Fe 20 Si 20 Si wt % wt %

10 Mg 10 Mg

0 0 050 100 150 200 250 050 100 150 200 250 40 40 CJS1-5 EJS1-5 O O 30 30 Fe Fe 20 Si 20 Si wt % wt %

10 Mg 10 Mg

0 0 050 100 150 200 250 050 100 150 200 250 40 40 CJS1-6 EJS1-6 O O 30 30

Fe Fe 20 20

wt % Si wt % Si

10 Mg 10 Mg

0 0 050 100 150 200 250 050 100 150 200 250 Time (Myr) Time (Myr)

Figure 3.8: Temporal variation in the elemental abundances of the ﬁnal terrestrial planets produced by the CJS1 and EJS1 simulations. Variations in compositions are due to the accretion of embryos and planetesimals throughout the dynamical simulations. Black indicates the abundance of O, red indicates the abundance of Fe, blue indicates the abundance of Si and green indicates the abundance of Mg. Left: CJS1 simulation results. EJS1: EJS1 simulation results. 79

3.3.3 Late Veneer

An alternative hypothesis to homogenous accretion is heterogeneous accretion where the composition of material accreted changes significantly with time. Under this hypothesis, siderophile distribution within the Earth is explained via accretion of a “late veneer” (Chou, 1978). Referring to the last ∼1% of mass accreted by the Earth, this material would need to be siderophile-rich and highly oxidized, containing essentially no metallic iron (Drake and Righter, 2002). Although carbonaceous chondrites are sufficiently oxidized and contain limited metallic Fe, their Os isotopic abundances are not in agreement with values required for the late veneer. On the other hand, ordinary chondrites do posses Os isotopic ratios of the correct value but are not oxidized. As such, we currently have no samples in the meteorite record that represent a possible source for the late veneer material (Drake and Righter 2002). As expected from the previously observed homogenous accretion, the late veneer of the present simulations is similar in composition to the final planetary abundance. Here the late veneer is taken to be the material accreted after the last impact by a projectile with mass > Membryo. This material is not highly oxidized and is primarily composed of olivine, pyroxene, metallic iron, troilite (FeS), diopside (CaMgSi2O6), nickel and albite

(NaAlSi3O8). This composition is similar to that of the ordinary chondrites. Furthermore, as noted in O’Brien et al. (2006), the planets produced in the EJS simulations accrete an average of ∼10% of the ﬁnal planetary mass as a late veneer, an order of magnitude above the predicted amount. Thus it can be seen that the current simulations do not successfully reproduce the late veneer material. However, as discussed in sections 3.3.1 and 3.3.2, migration of material and temporal variations in the composition of solid material accreted have the potential to drastically alter the composition of the late veneer. Such variations are not currently captured by the current simulations. 80

3.3.4 Hydrous Species

Differences between the CJS and EJS simulations emerge when considering the delivery of hydrous species to the ﬁnal planets. In the present simulations, “hydrous species” refers to water ice and the aqueous alteration product serpentine (speciﬁcally clinochrysotile

Mg3Si2O5(OH)4). As stated in O’Brien et al. (2006), the planets formed in the EJS simulations do not accrete significant amounts of volatile rich material from beyond 2.5AU (with the exception of the same planet formed in the EJS3 simulations as previously men- tioned). Thus for disk conditions before 1×106 years, these planets are found to contain no hydrous species. For disk conditions after this time, however, all terrestrial planets produced in the EJS simulations are found to have significant water components, although generally less than those of the CJS terrestrial planets. The appearance of water and serpentine at later times is largely due to the migration of the ice line to be located within the region from which material is being obtained for the dynamical simulations. The planets formed via the CJS simulations, however, all contain significant amounts of serpentine and water ice for disk conditions at t = 5×105 years. At this time, seven planets include hydrous material (<0.008ML), incorporated via the accretion of one to four planetesimals each containing a minor amount of serpentine. Only one planet contained a significant amount of serpentine (0.02ML). All eight planets accreted the hydrous material before the late veneer and during the stage where large, violent impacts were occurring. As such, it is expected that a significant component of this hydrous material would be vaporized during later impact events and subsequently lost from the final planetary body or incorporated into the planetary core. Only four planets (including two Earth analogs) accreted serpentine (0.001 - 0.003ML) as part of their late veneer. Given the relatively late delivery of this material, it is believed that the majority would be retained during later planetary processing. As one would intuitively expect, the amount of hydrous material accreted by a planet increases with increasing planetary semimajor axis. Assuming that all of the hydrogen accreted as hydrous species is converted to water, 81 the current simulations are producing planets containing 0.6 to 24.8 Earth ocean masses of water for disk conditions at 5×105 years. If I assume that Venus initially possessed a similar amount of water to the Earth while Mars contained 0.06 - 0.27 Earth oceans of water (Lunine et al., 2003), then it can be seen that for all of the terrestrial planets considered, the present simulations are producing planets with sufficient water to avoid the need to invoke other large-scale exogenous water delivery sources such as cometary impacts. Of course, these values should be considered to be extreme upper limits on the amount of total water present within the planet as some will undoubtedly be lost by photodissociation, Jean’s escape and (definitely in the case of Earth) by the formation of organic species which contain large amounts of H. Additionally, primordial accretion of water (as predicted here) would result in H being accreted to the core of the planet via element partioning, thus acting to decrease the amount of water available on the planetary surface (Okuchi, 1997). Although values for the amount of water expected to be lost by these processes are unknown, it is still likely that a significant portion of the initial water will be retained within or on the crust. Temporal variations in the composition of the disk are likely to increase the amount of hydrous material accreted by the planets due to water-rich material being stable over a drastically larger fraction of the disk. These time-dependent variations can be observed in my current approach as simulations undertaken for conditions at 3×106 years produce terrestrial planets containing up to 1200 Earth ocean masses of water, well above the levels observed for simulations under disk conditions at 5×105 years. My current approach also does not account for the possibility of water delivery via adsorption of water onto solid grains later incorporated into planetesimals (Drake and Campins, 2006). As all of the solid material considered in the current simulations would be bathed in H2O vapor, it is possible that a significant amount of water could be delivered to the final planets via adsorption that is not presently accounted for. Thus it appears likely that the terrestrial planets form “wet” with a sizeable portion of their primordial water delivered as a natural result the planetary formation process. 82

This conclusion does not consider the resulting D/H ratio of the accreted water, a key constraint on any water delivery hypothesis, as isotopic abundances can not be determined in our current approach. Primordial D/H ratios are known to vary greatly, from 2−3×10−5 for protostellar hydrogen (Lecluse´ and Robert, 1994) to 9×10−5 for aqueous inclusions in meteorites (Deloule and Robert, 1995). Similar variations can be seen in solid bodies with carbonaceous chondrites containing values close to that of the Earth (1.5×10−4) while Mars appears to be enriched in D with D/H values around 3×10−4 (Drake and Righter, 2002). Thus it is difﬁcult to make detailed predictions about the possible D/H ratios of the simulated planets. However zeroth order predictions can be made if I assume that the D/H ratio decreases linearly with increasing semimajor axis from the observed Martian value to that of Jupiter’s atmosphere (2.6×10−5). As all of the hydrous species in the current simulations are produced beyond 3.6AU (for disk conditions at t = 5×105 years), I can assume that the D/H ratio will be less than 1.4×10−4. Thus it is expected that the D/H ratios for the simulated planets will be less than the currently observed planetary values. However, given the large degree of processing this material is expected to experience both during and after planet formation, it is likely that the planetary D/H values will increase as H is preferentially lost from the system.

3.3.5 Volatile Loss

Although the loss of volatile material in large impact events is believed to be significant, detailed studies of such a loss during terrestrial planet formation have not been undertaken. As such, I am limited to making only first-order approximations of the amount of each element lost from the final planet due to impacts. To do this, I need to determine both the amount of the final body that is molten and/or vaporized after each impact and the amount of each element that would be lost from the molten phase. Using equation 9 from Tonks and Melosh (1993), the volume of melt produced by the initial shockwave of 83

an impact, Vm, is given by:

Ã !3/2 ρP vi Vm = Vproj 1/2 m (3.3) ρtcos i vi

where Vproj is the volume of the projectile, ρP is the density of the projectile, ρt is the density of the target, i is the impact angle (as measured from the vertical), υi is the impact m speed and υi is the minimum impact speed needed to produce melting. Thus it can be seen that to ﬁrst order:

3/2 Vm ∝ VprojρPvi (3.4)

Since the fraction of the planet that is molten (fm) is also directly proportional to the volume of the planet that is molten, I have:

3/2 fm ∝ VprojρPvi (3.5)

Thus it is possible to determine the fraction of the ﬁnal planet that is molten after each

−1 impact by scaling the fraction given in Tonks and Melosh (1993) for a υi = 15 kms impact with the Vproj, ρP and υi values determined by the O’Brien et al. (2006) models. m Note that this approach assumes constant values for ρt, i and υi . Although these values will undoubtedly change, I am unable to provide any limitations for them at this time. As such, I adopt the values from Tonks and Melosh (1993). The amount of each element assumed to be lost from the molten phase is based on the condensation temperature of the element. The most volatile element present in solid form 84

(C, primarily present as CH4.7H2O and polyaromatic hydrocarbons in the Solar nebula) was set to lose 100% of its mass from the melt. For all other elements, the percentage lost was determined from the observed elemental depletions in ordinary chondrites (as compared to CI chondrites) with respect to the 50% condensation temperature as taken from Davis (2006). CI chondrites are widely believed to be the most primitive and un- altered meteorites. Thus the elemental depletions relative to CI chondrites observed in other meteorite classes are believed to be due to processing of material within the solar protoplanetary disk and as such serves as an excellent proxy for the loss of volatile material. The percentage of each element assumed to be lost in each impact event is shown in Table 3.5. Note that the condensation temperature of O is taken to be the condensation temperature of silicate, the dominant form of O throughout the majority of the disk. The amount of each element lost was then determined after each impact event, based on the

υi value of the projectile itself. All lost volatile material was assumed to be permanently removed from the planet. This approach is obviously based on several broad assumptions and can only provide order of magnitude approximations for the loss of volatile material. In addition to as-

1/2 suming details about each impact event (such as ρt, cos i), this approach assumes that the target body has completely cooled and resolidiﬁed between impacts (i.e. each impact occurs with two cold solid bodies). A hotter body produces a larger melt fraction (Tonks and Melosh, 1993) and is expected to result in greater loss of material from each impact event. Furthermore, I currently only consider melting produced by the initial shockwave as it moves through the body. Detailed hydrocode simulations need to be undertaken in order to examine how volatile losses would vary under more realistic conditions including differentiation (or lack thereof) of the target body, atmospheric losses, impacts into a still- molten embryo and recapturing of re-condensed species, as well as determining the fate of the ejected material. Furthermore, the current dynamical models assume perfect accretion for each of the impact events. To obtain more realistic simulation conditions, loss of volatile material needs to be incorporated into the simulated impact events. However, this 85

Table 3.5: The percentage of each element assumed to be lost from the melt produced by individual impact events. Values are based on depletions observed in meteorite samples and are taken from Davis (2006).

Element Tcond (K) % Lost C 78 100.00 N 131 97.69 H 182 92.90 S 704 49.90 Na 958 32.92 P 1248 16.69 Cr 1296 14.32 O 1316 13.37 Ni 1353 11.64 Fe 1357 11.45 Mg 1397 9.65 Si 1529 4.16 Ti 1593 1.75 Ca 1659 0.00 Al 1677 0.00

approach is currently computationally very demanding and is not yet feasible. As expected, loss of material through impacts reduced the amount of volatile species present in the final planetary bodies but did not significantly alter the abundance of more refractory species. Table B.4 shows the abundances of the simulated planets after impact- induced elemental loss has been incorporated. This loss of volatile elements can best be seen in the planet normalized abundances shown in Figure 3.9. Normalized abundances for the other six simulations are shown in Figures B.8 - B.14. Clear reductions in the amount of volatile elements (specifically Na and S) can be seen, while negligible changes are produced in the abundances of the more refractory elements (such as Fe, Cr and Mg). These reductions in the abundance of the most volatile elements are unable to produce final planetary abundances in exact agreement with the observed planetary abundances within the Solar System but they do represent a substantial improvement, and indicate 86

10 10 CJS1-4 (Venus) EJS1-4 (Venus)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS1-5 (Earth) EJS1-5 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS1-6 (Mars) EJS1-6 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure 3.9: Normalized abundances for CJS1 and EJS1 simulated planets showing abundances after volatile loss was considered. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. Reference Solar Sys- tem planetary abundances were taken from Morgan and Anders (1980) (Venus), Kargel and Lewis (1993)(Earth) and Lodders and Fegley (1997) (Mars). Left: CJS1 terrestrial planets. Right: EJS1 terrestrial planets. 87 that impact-induced melting and vaporization (and the associated loss of material) is an important factor in determining the bulk elemental abundances of volatile species within a planet. The average fraction of the planet melted in each individual impact event is < 5% for both simulations (2.8% for the CJS simulations, 3.5% for the EJS simulations). Only one simulation (EJS1) produced an impact event large enough to melt and/or vaporize the entire target body. This event was the 10th impact occurring on this body and thus occurred early in the dynamical formation simulation (t = 9.72×106 years for the dynamical simulation). As such, it is likely that although the entire body would have been disrupted by the event, a solid body may have reformed from the remaining material afterwards. More detailed simulations are required to determine the full effects of large-scale impacts such as this one on the terrestrial planet formation process.

3.3.6 Solar Pollution

Insigniﬁcant amounts of pollution of the Solar photosphere occurred during the terrestrial planet formation simulations. A maximum of 0.135ML of solid material was added to the Sun in the CJS simulations, while the EJS simulations contributed a maximum of 1.11ML. The ensemble-averaged resulting solar abundances are shown in Table 3.6. The addition of solid material generated in the CJS simulations produced no observable enrichment in the Solar spectrum. The EJS simulations, however, did produce a minor enrichment of up to 0.02 dex for Ti and Al, 0.01 dex for C, N, Na, Mg, Al, Si, P, S, Ca, Cr, Fe and Ni and no enrichment in O. Nonetheless, this enrichment is not large enough to be deﬁnitively detected with current spectroscopic studies as many such studies produce stellar abundances with errors equivalent to or larger than the expected enrichment (e.g. ±0.03 for Fischer and Valenti (2005)). As such, any Solar pollution produced by terrestrial planet formation with the Solar System is believed to be negligible. 88

Table 3.6: Mean change in solar photospheric abundances produced by pollution via accretion of solid material during terrestrial planet formation. All solid material migrating to within 0.1AU from the Sun during the simulations of O’Brien et al. (2006) is assumed to be accreted. Change is deﬁned as Abundanceafter planet formation - Abundancebefore planet formation.

Element Simulation CJS-1 CJS-2 CJS-3 CJS-4 EJS-1 EJS-2 EJS-3 EJS-4 Mg 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 O 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 S 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Fe 0.00 0.00 0.00 0.00 0.00 0.01 0.01 0.01 Al 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Ca 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Na 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Ni 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Cr 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 P 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Ti 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 Si 0.00 0.00 0.00 0.00 0.01 0.01 0.01 0.01 89

3.4 Discussion

My simulations successfully produce terrestrial planets that are in excellent agreement with the terrestrial planets of the Solar System, in terms of both their dynamics and their bulk elemental abundances. Although current simulations are unable to capture the finer details of planetary formation (such as the amount and composition of the late veneer material), the success of these simulations on a broader scale provides us with increased confidence in the dynamical models of O’Brien et al. (2006). Furthermore, it also serves to validate the approach utilized here to combine detailed dynamical and chemical modeling together. This will allow for reliable application of this approach not only to other dynamical models but also to extrasolar planetary systems. The final bulk elemental compositions of the simulated terrestrial planets for rock forming elements are not strongly constrained by the orbital properties or evolution of Jupiter and Saturn as can be seen by the strong similarities between the CJS and EJS simulations. The only significant differences occur in the amount of water rich material accreted onto the final planets. This suggests that the bulk chemical evolution of the terrestrial planets is to a large extent independent of the evolution of the giant planets. This conclusion, however, is only valid for late stage and in situ formation after the giant planets have formed and undertaken the vast majority of their migration. Simulations are currently running to examine the effects on planetary composition of formation occurring during giant planet migration. As Jupiter and Saturn are not believed to have undergone extensive migration based on recent studies (Gomes et al., 2005; Levison et al., 2005; Morbidelli et al., 2005), the issue of migration is not a significant one for the Solar System. However, it will be an important issue for extrasolar planetary systems where many are thought to have experienced large amounts of migration. The delivery of water and other volatile species to the final terrestrial planets appears to be a normal outcome of terrestrial planet formation, implying that it is normal for such planets to accrete significant amounts of water. As such, it negates the need for large-scale 90 delivery of water and other volatile species by such exotic processes as cometary impacts. Later planetary processing (especially within the mantle) would have resulted in dis- sociation of some of the primordial water, causing a significant amount of H to migrate to the core of the Earth or alternatively to be lost from the planet. This migration, however, would have left a large amount of OH in the mantle, thus increasing its redox state. This process allows us to explain the evolution of the Earth’s redox state over time from the initially reduced state produced by the current simulations to the present stratified redox state. This conclusion, however, requires a more detailed experimental understanding of the efficiency of H partioning within the mantle and core of the Earth, along with the evolution of mantle processes and mixing. It is interesting to note the order of magnitude difference between the CJS and EJS simulations in the amount of solid material added to the Sun during terrestrial planet formation. This may potentially have great bearing on extrasolar planetary host stars as many known extrasolar planets are currently in eccentric orbits. If such an orbit always results in a greater amount of solid material being accreted by the host star, then it may give more weight to the pollution hypothesis which has previously been suggested to explain the observed [Fe/H] enrichment (e.g. Laughlin 2000; Gonzalez et al. 2001; Murray et al. 2001). However, no correlation has been found between known planetary eccentricity and the metallicity of the host star (e.g. Reid, 2002; Santos et al., 2003; Fischer and Valenti, 2005; Bond et al., 2006), suggesting that this is not a strong effect. I determined that a negligible amount of solid material was added to the Solar photosphere during terrestrial planet formation and that no observable elemental enrichment would be produced. While similar simulations are needed for extrasolar planetary systems, the current simulations support the conclusion that the observed metal enrichment in extrasolar host stars is primordial in origin, established in the giant molecular cloud from which these systems formed as has been concluded by other studies (such as Santos et al. 2001, 2003, 2005 and Fischer and Valenti 2005). Of course, this does not rule out the possibility that extrasolar planetary host stars may have in fact accreted a giant planet during the formation 91 and migration process and thus display higher levels of pollution. Host star pollution for extrasolar planetary systems is discussed further in Chapter 4. Finally, biologically important elements are obviously of great interest, especially for the Earth. Of the six major biogenic elements (H, C, N, O, S and P), four are accreted in excess by the planets during their formation (H, O, S and P). Only C and N are not accreted during planetary formation. Both species are primarily present as solids within an equilibrium Solar nebula as clathrates (methane or ammonia trapped in a water ice lattice) and organics. As these species only form in the outermost regions of the disk, they could only be delivered to the final planets via cometary and meteorite impacts, migration or temporal variations within the disk or some combination of all three. Therefore delivery of material from the outer regions of the disk are necessary in the current models for life to be able to develop. Assuming the outer regions of the disk are the source for C and N, all of the biologically important elements are accreted in the form believed to be required required for the early evolution of life. C and N are reduced in their clathrate forms and O would be present in several forms (oxides, silicates and hydrous material). S is primarily accreted in its reduced form as troilite (FeS) while P accreted both as schribersite (Fe3P) and phos- phates, both of which can be utilized by early life. Therefore our results are in agreement with current predictions for chemical requirements for the evolution of early life on Earth.

3.5 Summary

Bulk elemental abundances have been determined for the simulated terrestrial planets of O’Brien et al. (2006). These abundances are in excellent agreement with observed planetary values, indicating that the models of O’Brien et al. (2006) are successfully producing planets comparable to those of the Solar System in terms of both their dynamical and chemical properties, adding greater weight to their predictive properties. Simulated redox states are also in agreement with those predicted for the early Earth. Although differences 92 do exist between the observed and predicted geochemical ratios, these are believed to be a result of our assumption of equilibrium controlled compositions. Additionally, the current simulations are unable to successfully reproduce the accretion of a late veneer of material by the early Earth, in terms of both the chemistry and, in the case of the EJS simulations, the amount. Significant amounts of water are accreted in the present simulations, implying that the terrestrial planets form “wet” and do not need significant water delivery from other sources. N and C, however, do still need to be delivered to an early Earth by some other process in order for life to develop. Additionally, the bulk elemental abundances of the final planets in the current simulations are not strongly dependent on the orbital properties of the giant planets with the CJS and EJS simulations both producing comparable results. This suggests that although the orbits of Jupiter and Saturn are of great dynamical importance to the evolution of the terrestrial planets, they may not exert such a large influence over the chemical evolution of the same planets in late stage in-situ formation. Finally, the pollution of the outer layers of the Sun via solid material during planetary formation produces a negligible photospheric elemental enrichment. Assuming similar levels of pollution in other planetary systems, this in turn implies that the high metallicity trend observed in extrasolar planetary systems is in fact primordial. 93

Figure 3.10: PEARLS BEFORE SWINE °c Stephan Pastis/Dist. by United Feature Syn- dicate, Inc. Originally published 4/2/2007. 94

CHAPTER 4

EXTRASOLAR PLANETARY SYSTEM SIMULATIONS

4.1 Introduction

Extrasolar terrestrial planets are a tantalizing prospect. Given that the number of planets in the galaxy is expected to correlate inversely with planetary mass, it is expected that Earth-sized terrestrial planets are much more common than giant planets (Marcy et al., 2000). Although still undetectable by current exoplanet searches, the possibility of their existence in extrasolar planetary systems has been examined by several authors. Many such studies have focussed on the long term dynamical stability of regions within the planetary system where such planets could exist for geologic timescales (Barnes and Raymond, 2004; Raymond and Barnes, 2005; Asghari et al., 2004). Several systems have been found to posses such regions (e.g. Barnes and Raymond 2004), indicating that if they are able to form, terrestrial planets may still be present within extrasolar planetary systems. Additionally, many of these systems appear to be ‘packed’, containing no available space in which another planet could be inserted and still be dynamically stable (Barnes and Raymond, 2004). If this same packing principle holds true for other systems then it can be utilized to predict regions in which planets are likely to be detected. Analyses of this nature are of great interest to future planet search missions as they assist in constrain- ing future planet search targets. However, they provide little insight into the formation mechanism of such planets and do not necessarily indicate the presence of a terrestrial planetary companion. Few other studies have gone one step further and undertaken detailed simulations of terrestrial planet formation within specific systems. Raymond et al. (2005) considered terrestrial planet formation in a series of hypothetical ‘hot Jupiter’ simulations and found 95 that terrestrial planets can indeed form in such systems provided the ‘hot Jupiter’ is located within 0.5AU from the host star. Furthermore, such planets may even have water contents comparable to that of the Earth. Terrestrial planets have been found to form even in simulations of systems which have undergone large-scale migration of the known giant planet (Mandell et al., 2007). Terrestrial planets were found to form both exterior and interior to the giant planet after migration has occurred and many were located within the habitable zone of the host star. As many extrasolar planets are believed to have experienced such a migration, it is encouraging to still be able to form terrestrial planets within these systems. To date, only Raymond et al. (2006) have undertaken terrestrial planet formation simulations for specific planetary systems. They considered four known planetary systems and found that terrestrial planets could form in one of the systems (55Cancri). Small bodies comparable to asteroid sized objects would be stable in another (HD38529). Although such simulations have not been undertaken for a large number of planetary systems, early studies have indicated that approximately one-third of extrasolar planetary systems may be ‘habitable’, containing a terrestrial planet located within the habitable zone of the system (Mandell et al., 2007) and that many more may still contain other terrestrial bodies. An even more intriguing question beyond whether or not such terrestrial planets could exist within these systems is their potential chemical composition. Extrasolar planetary host stars are already known to be chemical unusual (Gonzalez, 1997, 1998; Butler et al., 2000; Gonzalez and Laws, 2000; Gonzalez et al., 1999; Gonzalez and Vanture, 1998; Santos et al., 2000, 2001, 2004; Gonzalez et al., 2001; Smith et al., 2001; Reid, 2002; Fischer and Valenti, 2005; Bond et al., 2006) (see Chapter 2), displaying systematic enrichments in Fe and smaller, less statistically significant enrichments in other species such as C, Si, Mg and Al (Gonzalez and Vanture, 1998; Gonzalez et al., 2001; Santos et al., 2000; Bodaghee et al., 2003; Fischer and Valenti, 2005; Beirao˜ et al., 2005; Bond et al., 2006). Given that these enrichments are primordial in origin (Santos et al., 2001, 2003, 2005; Fischer and Valenti, 2005), it is thus natural to assume that the planet forming ma- 96 terial within these systems will be similarly enriched. Hints of such a correlation between transiting giant planets and stellar metallicity have been observed (Guillot et al., 2006; Burrows et al., 2007). Consequently, it is possible that terrestrial extrasolar planets may have compositions reflecting the enrichments observed in the host stars. Furthermore, several known host stars have been found to have C/O values above 0.8 (see Chapter 2). Systems with high C/O ratios will contain large amounts of C phases (such as SiC, TiC and graphite), resulting in any terrestrial planets within these systems being enriched in C and potentially having compositions and mineralogies unlike any body yet observed within our Solar System. Despite the likely chemical peculiarities and the early successes of terrestrial planet formation simulations, no studies of extrasolar terrestrial planet formation completed to date have considered both the dynamics of formation and the detailed chemical compositions of the final terrestrial planets produced. This present study addresses this issue by simulating late-stage in-situ terrestrial planet formation within nine extrasolar planetary systems while simultaneously determining the bulk elemental compositions of the planets produced. This is the first such study to consider both the dynamical and chemical nature of potential extrasolar terrestrial planets and it represents a significant step towards gaining a greater understanding of the full diversity of extrasolar terrestrial planets.

4.2 System Composition

Before considering detailed simulations of planetary formation and composition, we ﬁrst need to consider the geochemical trends observed in extrasolar planetary host stars.The two most important elemental ratios for determining the mineralogy of extrasolar terrestrial planets are C/O and Mg/Si. The ratio of C/O controls the distribution of Si among carbide and oxide species. If the C/O ratio is greater than 0.8, Si exists in solid form primarily as SiC. Additionally, signiﬁcant amounts of solid C are also present as planet building materials. For C/O values below 0.8, Si is present in rock-forming minerals as 97

SiO4, allowing for the formation of silicates. The silicate mineralogy is controlled by the

Mg/Si value. For Mg/Si values less than 1, Mg is in pyroxene (MgSiO3) and the excess Si is present as other silicate species such as feldspars. For Mg/Si values ranging from 1 to 2,

Mg is distributed between olivine (Mg2SiO4) and pyroxene. For Mg/Si values extending beyond 2, all available Si is consumed to form olivine with excess Mg available to bond with other elements as MgO. The photospheric C/O vs. Mg/Si values for known extrasolar planetary host stars are shown in Figure 4.1, based on stellar abundances taken from Gilli et al. (2006) (Si and Mg), Beirao˜ et al. (2005) (Mg), Ecuvillon et al. (2004) (C) and Ecuvillon et al. (2006b) (O). A conservative approach was taken in determining the errors shown in Figure 4.1 (bottom panel). The errors published for each elemental abundance were taken as being the 2 σ errors (based on the method used to determine them) and were used to determine the maximum and minimum abundance values possible with 2 σ confidence for each system. The elemental ratios produced by these extremum abundances were thus taken as the 2 σ range in ratio values and are shown as errors in Figure 4.1. The mean values of Mg/Si and C/O for all extrasolar planetary systems for which reliable abundances are available are 1.28 and 0.65 respectively, which are above solar values (Mg/SiJ = 1.00 and C/OJ = 0.54). This non-solar average and observed variation implies that a wide variety of terrestrial planet compositions are present within extrasolar planetary systems and that not all of them can be expected to be identical to that of Earth. Of the 62 systems shown, 13 have C/O values above 0.8, implying that carbide minerals are important planet building materials in over 20% of planetary systems. The idea of C- rich planets is not new (Kuchner and Seager, 2005) but the potential prevalence of these bodies has not been previously recognized, nor have specific systems been identified as likely C-rich planetary hosts. These data demonstrate that there are a significant number of systems in which terrestrial planets have compositions vastly different to any body observed in our Solar System. However, a high degree of uncertainty is associated with the values shown in Fig- 98

2.5

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Figure 4.1: Mg/Si vs. C/O for known planetary host stars with reliable stellar abundances. Stellar photospheric values were taken from Gilli et al. (2006) (Si, Mg), Beirao˜ et al. (2005) (Mg), Ecuvillon et al. (2004) (C) and Ecuvillon et al. (2006b) (O). Solar values are shown by the red circle and were taken from Asplund et al. (2005). The dashed line indicates a C/O value of 0.8 and marks the transitions between a silicate-dominated composition and a carbide-dominated composition. Top Panel: Stellar ratios with error bars removed for clarity. Bottom Panel: Stellar ratios with 2-σ error bars included. 99 ure 4.1. The primary source of this error is errors in the stellar elemental abundances themselves. Spectrally determined abundances are sensitive to continuum placement and stellar atmospheric parameters such as Teff , log g, metallicity and microturbulence. The uncertainty produced by each parameter is determined via sensitivity studies. Each parameter was varied in turn by a specified amount (±100 K for Teff , ±0.3 dex for log g, ±0.3 dex for metallicity and ±0.05 dex for microturbulence) and the resulting variation in the elemental abundance was determined. Thus the final error was obtained by summing in quadrature the sensitivity errors, continuum placement error (typically 0.05 dex) and standard deviation of each mean abundance (where the elemental abundance was determined from more than one spectral line):

2 2 2 2 2 2 2 σﬁnal = σstd + σcontinuum + σTeﬀ + σlogg + σ[Fe/H] + σmicro (4.1)

This results in an average error of ±0.04 for [Mg/H], ±0.07 for [Si/H], ±0.08 for [C/H] and ±0.09 for [O/H]. These errors result in considerable percentage uncertainties for the Mg/Si and C/O values of up to 124%. Although large, the errors will not decrease until we are able to improve the uncertainty on each individual stellar elemental abundance. Of the four elements considered here, O is the most troublesome and controversial. Three different spectral lines are available for determining the photospheric O abundance - the forbidden OI lines located at 6300 and 6363 A˚ the OI triplet located between 7771 A˚ and 7775 A˚ and the OH lines located near 3100 A˚ (Ecuvillon et al., 2006b). Previous studies have found discrepancies between abundances obtained from different spectral lines for the same star of up to 1 dex (Israelian et al., 2004). Each of these lines is subject to interferences from different stellar sources. The forbidden OI lines are weak and blended with Ni, the OI triplet is influenced by non-local thermodynamic equilibrium (non-LTE) effects and the OH lines are influenced by stellar surface features (Ecuvillon 100 et al., 2006b). Ecuvillon et al. (2006b) undertook a detailed examination of the correlation between these different O indicators for a sample of 96 host and 59 non-host stars. Their study showed that the discrepancies in abundances determined by the three different indicators was less than 0.2dex for the majority of stars examined. The forbidden OI and OH lines were found to be in excellent agreement with each other while abundances obtained from the O triplet lines (with the appropriate NLTE corrections applied to them) were systematically lower. However, all three indicators produced abundances in keeping with the galactic evolutionary trends observed for lower metallicity (and thus younger) stars (Ecuvillon et al., 2006b). The C/O ratios shown in Figure 4.1 are based on the O abundances from Ecuvillon et al. (2006b) obtained from the forbidden OI spectral line observed at 6300.3 A˚ as abundances from this line are in agreement with the abundances obtained from the OH line and produce a marginally better fit to stellar evolutionary models Given the errors associated with each individual elemental abundance and thus also ratio value, it is natural to consider the error associated with the dispersion seen in Figure 4.1. The observed dispersion is produced by both dispersion in the data and dispersion due to errors. Thus it is probable that the real range in elemental ratios is less than is shown in Figure 4.1 and fewer planetary systems have C/O values above 0.8. Based on the errors described above and shown in Figure 4.1, only 2 of the 62 planetary systems shown (3% of the sample) can be said with 2 σ confidence to have C/O values above 0.8. This increases to just 3 planetary systems (5% of the sample) when I reduce the confidence interval to 1 σ. It should be noted that such a drastically reduced C-rich population is a worst case scenario, based on the assumption that all potentially C-rich systems identified in Figure 4.1 have a true, error-free C/O ratio at the lower limit of their 2-sigma errors. This is not likely to be true for all systems but allows for a conservative estimation of the true C/O distribution. Similar values are observed for the abundance ratios of both host and non-host stars obtained in Chapter 2. Based on the values and errors listed in Tables A.1 and A.2 and 101 shown in Figure 4.2, only 2 of the 26 host stars (8% of the sample) and 2 of the 77 non-host stars (3% of the sample) can be said to have C/O values above 0.8 with 2 σ confidence. Thus it is probable that a reduced number of planetary systems will be C-rich while most will be Solar-like in terms of their bulk elemental ratios. This represents a significant reduction in spread from that shown in Figure 4.1. If the real dispersion is this small, then it implies that C-rich systems would be relatively rare and terrestrial extrasolar planets (if present) would be dominated by pyroxene and olivine in almost all known planetary systems. Although a diminished distribution in stellar C/O values would reduce the prevalence of the extremely C-rich condensation sequences of the present study, the fact remains that several planetary systems would still contain significant amounts of C and carbide phases as major planet building elements. Thus although potentially not common, the C-rich systems considered here still warrant detailed study and investigation. Although it is highly desirable to compare Mg/Si and C/O values for host and non- host star populations, sufficient data is currently unavailable for a large sample of non-host stars. Although Gilli et al. (2006); Beirao˜ et al. (2005); Ecuvillon et al. (2006b) do provide Mg, Si and O abundances for non-host stars, C abundances are provided by Ecuvil- lon et al. (2004) for just 3 non-host stars in common with these previous studies. Due to the potential of introducing systematic instrument and/or processing errors by combining stellar abundances from a variety of sources and the lack of available abundances, a comparison between host and non-host stars was not undertaken based on values published by these studies. Instead, the stellar abundances determined in Chapter 2 and shown in Figure 4.2 were compared. The C/O and Mg/Si distributions for both host and non-host stars are shown in Figure 4.3. Both the host and non-host star populations display the same distributions. The mean, median and standard deviation for both the host and non-host stars is shown in Table 4.1. Given the excellent agreement between the two populations, I conclude that known planetary host stars are not preferentially biased towards higher C/O or Mg/Si 102

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0.5

0.0 0.0 0.5 1.0 1.5 2.0 2.5 Mg/Si

Figure 4.2: Mg/Si vs. C/O for host and non-host stars based on abundances determined in Chapter 2. Open circles indicate stars not currently known to harbor a planetary companion. Filled circles indicate known planetary host stars. Solar values are shown by the red circle and were taken from Asplund et al. (2005). The dashed line indicates a C/O value of 0.8 and marks the transitions between a silicate-dominated composition and a carbide-dominated composition. Top Panel: Stellar ratios with error bars removed for clarity. Bottom Panel: Stellar ratios with 2-σ error bars included. 103

10 10 Chapter 2 Chapter 2 host stars host stars f stars f stars 5 5 r o r o Numbe Numbe

0 0 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 30 30 Chapter 2 Chapter 2

25 non-host stars 25 non-host stars

20 20 f stars f stars 15 15 r o r o

10 10 Numbe Numbe

5 5

0 0 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 C/O Ratio Mg/Si Ratio

Figure 4.3: C/O and Mg/Si distributions for host and non-host stars based on the abundances determined in Chapter 2 and shown in Table A.2. Left: C/O distributions for host (top) and non-host (bottom) stars. Right: Mg/Si distributions for host (top) and non-host (bottom) stars. values compared to stars not known to harbor a planetary companion. This in turn implies that the prevalence of C-rich planetary systems identiﬁed above is not statistically unusual (in terms of stellar composition). However, the Mg/Si and C/O values shown in Figure 4.3 should not be directly compared to values obtained from abundances provided other similar spectroscopic studies (such as those of Beirao˜ et al. (2005) and Gilli et al. (2006)). As was previously discussed in Chapter 2, systematic deviations from previously published abundances for the 29 host stars in common with the studies of Beirao˜ et al. (2005) and Gilli et al. (2006) were iden- tiﬁed in the [Mg/H] values. The observed difference is due to the use of a smaller number 104

Table 4.1: Statistical analysis of the host and non-host star distributions of Mg/Si and C/O. All values are based on the abundances determined in Chapter 2. The quoted uncertainty is the standard error in the mean. Mean Median Standard Deviation Mg/Si: Host Stars 0.83± 0.04 0.80 0.22 Non-Host Stars 0.80± 0.03 0.79 0.16

C/O: Host Stars 0.67± 0.03 0.68 0.23 Non-Host Stars 0.67± 0.03 0.69 0.23

of Mg spectral lines in the present study. This abundance shift acts to skew the Mg/Si distribution towards lower values and thus prohibiting direct numerical comparison. This can be seen in Figure 4.4 where the Mg/Si distributions for both host and non-host stars are noticeably offset from previously published values. Additionally, the O abundances of Chapter 2 were obtained from the O triplet located at 7771 A˚ 7774 A˚ and 7775 A˚ . As previously discussed, O abundances obtained from these spectral lines have been found by Ecuvillon et al. (2006b) to be lower than those from the OI line located at 6300.3 A˚ . The resulting shift in C/O values, however, appears to be negligible based on the distributions shown in Figure 4.4. Until the complete reasons for these offsets are understood and corrected for (where possible), direct comparisons will be restricted to comparing distributions only. 105

15 15 Previously published Previously published host stars host stars

10 10 f stars f stars r o r o

5 5 Numbe Numbe

0 0 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 15 15 Chapter 2 Chapter 2 host stars host stars

10 10 f stars f stars r o r o

5 5 Numbe Numbe

0 0 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 30 30 Chapter 2 Chapter 2 non-host stars non-host stars 25 25

20 20 f stars f stars 15 15 r o r o

10 10 Numbe Numbe

5 5

0 0 0.0 0.5 1.0 1.5 2.0 2.5 0.0 0.5 1.0 1.5 2.0 2.5 C/O Ratio Mg/Si Ratio

Figure 4.4: C/O and Mg/Si distributions for previously published host stars, along with values for host and non-host stars based on the abundances determined in Chapter 2 and shown in Table A.2. Left: C/O distributions for previously published host stars (top), host stars from Chapter 2 (middle) and non-host stars from Chapter 2 (bottom). Right: Mg/Si distributions for previously published host stars (top), host stars from Chapter 2 (middle) and non-host stars from Chapter 2 (bottom). 106

Planetary Systems

HD177830 HD4203 HD108874 HD19994 HD142415 HD72659 HD4208 GJ777A 55Cnc 0 1 2 3 4 5 6 7 Semimajor Axis (AU)

Figure 4.5: Location of known giant planets in the systems selected for study. The horizontal lines indicate the variation from periastron and apastron. The size of the circles scales with the planetary Msini value. All planets are assumed to have zero inclination. All values taken from the Butler et al. (2006) catalog.

4.3 Simulations

4.3.1 Extrasolar Planetary Systems

Nine known extrasolar planetary systems spanning the entire spectrum of observed planetary systems were selected for this study. By studying planetary systems with such a wide range of both chemical and dynamical properties, I am exploring the full diversity of possible extrasolar terrestrial planets. The dynamical and chemical details of each system are shown in Tables 4.2, (orbital parameters of known planetary companions), 4.3 (stellar abundances in logarithmic units) and 4.4 (stellar abundances normalized to 106Si atoms), while the known giant planet architecture is shown in Figure 4.5 and the C/O and Mg/Si values are shown in Figure 4.6. Note that the innermost planet of 55Cnc was neglected in our present simulations due to its location and low mass. 107

Table 4.2: Orbital parameters of known extrasolar planets for the systems studied. Values taken from the University of California catalog located at www.exoplanets.org

Planet M a e (MJupiter) (AU) 55Cnc-b 0.82 0.11 0.02 55Cnc-c 0.17 0.24 0.05 55Cnc-d 3.84 5.84 0.08 55Cnc-e 0.02 0.04 0.09 55Cnc-f 0.14 0.70 0.20

Gl777-b 1.55 4.02 0.35 Gl777-c 0.06 0.13 0.07

HD4203-b 2.07 1.16 0.52

HD4208-b 0.81 1.64 0.01

HD19994-b 1.69 1.43 0.30

HD72659-b 3.30 4.76 0.26

HD108874-b 1.30 1.05 0.21 HD108874-c 1.07 2.75 0.16

HD142415-b 1.69 1.07 0.50

HD177830-b 1.43 1.22 0.03 108 0.02 0.21 − − 0.07 0.13 − 0.10 0.38 − 0.11 0.21 0.48 0.23 0.14 0.31 0.12 0.030.01 0.01 0.15 0.13 0.18 − − − − 0.05 − 0.24 0.24 0.03 0.23 0.33 0.21 0.010.140.22 0.39 0.12 0.110.10 0.480.23 0.210.38 0.11 0.32 0.07 0.23 0.11 0.07 0.05 0.13 0.27 0.30 0.14 0.37 0.56 0.54 0.31 0.15 0.17 0.17 0.15 0.310.140.28 0.17 0.25 0.18 0.20 0.27 0.13 0.01 0.20 0.18 0.31 0.39 0.31 0.18 − − − − − − − − − − − − [S/H] 0.12 0.10 0.20 [C/H] 0.31 0.29 0.45 [O/H] 0.13 0.22 0.00 [Si/H] 0.29 0.24 0.44 [Ti/H] 0.36 0.32 0.41 [Fe/H] 0.33 0.24 0.40 [Al/H] 0.47 0.34 0.51 [Cr/H][Ni/H] 0.22 0.31 0.17 0.25 0.33 0.42 [Ca/H] 0.08 0.11 0.24 [Na/H] 0.26 0.26 0.42 [Mg/H] 0.48 0.33 0.48 Element 55Cnc Gl777A HD4203 HD4208 HD19994 HD72659 HD108874 HD177830 HD142415 Table 4.3: Target starreferences. elemental abundances in standard logarithmic units, normalized to H and Solar values. See text for 109 6 6 6 4 6 4 6 4 3 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 1.00 5.13 6.17 4.57 1.10 7.59 1.00 6.03 3.55 1.45 5.62 5 4 4 3 7 7 6 5 6 3 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 9.12 5.25 2.63 2.45 1.12 1.66 1.86 1.23 1.00 8.91 6.31 6 4 4 3 6 6 6 5 6 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 1.07 4.47 4.68 2.82 8.91 8.13 1.41 1.05 1.00 1.38 5.75 Si atoms. See text for references. 6 5 4 4 3 6 6 6 7 4 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 8.32 4.79 5.25 2.95 1.20 1.00 5.25 1.62 7.59 1.17 4.79 5 4 4 3 6 6 7 7 4 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 8.91 8.13 5.50 2.19 1.00 1.00 1.10 1.07 8.91 1.26 5.75 5 4 4 3 6 6 7 7 4 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 8.51 4.68 5.25 3.02 1.35 1.00 1.26 1.74 9.77 1.20 5.01 5 4 4 3 6 6 6 6 4 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 7.94 4.37 3.98 2.29 1.15 1.00 7.76 5.13 8.51 1.05 5.01 5 4 4 3 6 6 6 7 4 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × 8.71 4.79 4.68 2.95 1.29 1.00 8.51 1.35 9.12 1.15 5.37 5 3 4 4 6 6 6 6 5 4 4 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × C 7.94 O 9.77 Si 1.00 Ti 2.88 Fe 9.55 Al 1.10 CrNi 1.15 5.50 Ca 3.89 Na 4.27 Mg 1.62 Table 4.4: Target star elemental abundances as number of atoms and normalized to 10 Element 55Cnc Gl777A HD4203 HD4208 HD19994 HD72659 HD108874 HD177830 HD142415 110

2.5

2.0

1.5 C/O 1.0

0.5

0.0 0.0 0.5 1.0 1.5 2.0 2.5 Mg/Si

Figure 4.6: Mg/Si vs. C/O for planetary host stars studied. Filled circles indicate systems selected for this study. Stellar photospheric values were taken from Gilli et al. (2006) (Si, Mg), Beirao˜ et al. (2005) (Mg), Ecuvillon et al. (2004) (C) and Ecuvillon et al. (2006b) (O). Solar values are shown in red and were taken from Asplund et al. (2005). The dashed line indicates a C/O value of 0.8, the point at which the solid composition transitions from being silicate-dominated to being carbide-dominated.

4.3.2 Dynamical Simulations

In the current study, I build upon our recent success in simulating terrestrial planet formation within the Solar System (as discussed in Chapter 3) and apply the same methodology here. N-body simulations of terrestrial planet accretion in each of the selected extrasolar planetary systems are run using the SyMBA n-body integrator (Duncan et al., 1998). The orbital parameters of the giant planets in each system are taken from the catalog of Butler et al. (2006), and updated as additional data on these systems were obtained. Inclinations of each of the giant planets are assumed to be zero since no such measurements have been obtained for these systems. Due to the computational time required, current simulations only contain an initial population of embryos (i.e. no planetesimal swarm is included). This differs from the previous simulations of O’Brien et al. (2006) as used in Chapter 3. 111

As such, it is possible that differences between the current extrasolar dynamical simulations and those of the Solar System as used in Chapter 3 will occur, most likely with regards to the dynamical excitation of the system and accretion time. The inclusion of a large swarm of planetesimals (∼1000 in the simulations of O’Brien et al. (2006)) was found by O’Brien et al. (2006) to reduce the overall excitation of the system when compared to previous Solar System simulations such as that of Chambers and Wetherill (1998) and Chambers (2001). This is a result of the increased number of gravitational interactions occurring during the formation process between the planets and the large number of planetesimals. Each interaction acts to damp the eccentricity of the planet while simultaneously exciting the planetesimal. Thus the resultant degree of damping for the final planets is naturally increased in simulations containing a large number of planetesimals. As such, the exclusion of the planetesimal swarm from the present simulations may result in planets with a higher excitation. The degree of radial mixing for simulations with and without planetesimals is also of interest. O’Brien et al. (2006) found that all CJS simulations underwent a significant degree of radial mixing, indicating that planets were not being produced from material originally located adjacent to their final orbital positions. The simulations of Chambers (2001) with a smaller planetesimal population were found to produce a lower degree of radial mixing. However, the Chambers (2001) simulations only extended out to 2AU while the simulations of O’Brien et al. (2006) extended out to 4AU, thus permitting a larger degree of mixing to occur and not necessarily implying a connection between increased numbers of planetesimals and radial mixing. A better comparison of the effects of planetesimals on radial mixing within the dynamical simulations can be obtained by comparing the EJS simulations of O’Brien et al. (2006) with Model C from Chambers and Wetherill (1998) (with Jupiter and Saturn in their current orbital configurations and with planetesimals initially present to 4AU). Chambers and Wetherill (1998) find that 31% of embryos located between 2 and 3AU and 18% of embryos between 3 and 4AU are incorporated into the final planet. In contrast to this, O’Brien et al. (2006) finds that only trace 112 amounts of material from beyond 2.5AU are accreted by the final planet. No embryos from beyond 2.5AU are accreted and 5 of the 14 planets produced by the EJS simulations contain no planetesimals from this region either. The observed reduction in radial mixing is believed to be due to interactions with the planetesimals forcing embryos into reso- nances with the giant planets. From this point, it is relatively easy to remove the embryos from the system through ejection or accretion by the Sun. The same effect is believed to have produced the Kirkwood gaps within the asteroid belt within our Solar System. Thus it appears that the inclusion of a planetesimal swarm in simulations with eccentric orbits for Jupiter and Saturn reduces the degree of radial mixing occurring within the system. As this configuration is the most analogous to the majority of known extrasolar planetary systems, it is thought that a similar effect may be observed in simulations for extrasolar planetary systems. Of course, this is highly dependent on the exact orbital parameters of the planets present within the system. It is conceivable that planetesimals will be accreted from a larger semimajor axis range, but the increase in radial mixing produced may be offset by the embryos being accreted over a smaller range as gravitational interactions with planetesimals damp their orbits. To examine these effects in more detail, simulations were run for a hypothetical system with a Jupiter-mass planet located at 1AU. Two sets of four simulations were run - one with only embryos present and one with both embryos and ∼500 planetesimals present. Surface mass density profiles that vary as r−3/2, normalized to 10 gcm−3 at 1 AU, were assumed and each simulation was run for 1×108 years. The median number of planets formed in each set of simulations, 50% and 90% formation times and radial mixing parameter are shown in Table 4.5. As in O’Brien et al. (2006), the extent of radial mixing is characterized by σ where:

Σ m |a − a |/a σ = i i ﬁn,i init,i ﬁn,i (4.2) Σimi 113

where mi is the initial mass of each embryo, ainit,i is the initial semi-major axis of each embryo and afin,i is the final semi-major axis of each embryo (i.e. the semi-major axis of the planet produced) (Chambers, 2001). Based on the values shown in Table 4.5, it can be seen that the presence of a planetesimal swarm acts to reduce the number of terrestrial planets produced. Three of the four simulations run without planetesimals produced two terrestrial planets while each of the four simulations run with both embryos and planetesimals present produced only one terrestrial planet. Both planetary masses and formation timescales increased in the simulations including planetesimals. The increased timescale in the presence of planetesimals is in contrast to the result seen by Chambers and Wetherill (1998) and O’Brien et al. (2006) for Solar System simulations. The precise cause of this difference is not clear and requires further work. However, it is possible that in the planetesimal-free simulations the more highly excited embryos are simply unable to interact and collide with each other, thus resulting in a higher number of low mass terrestrial planets being produced in a relatively short timeframe. The presence of planetesimals, however, acts to dampen the excited embryos and produces a smaller number of larger planets. Additionally, the planetesimals may also be inducing some degree of radial migration within the embryo population, producing a pileup of material and thus a single, larger mass planet. Marginal differences can be seen in the degree of radial mixing encountered within each simulation. The embryo-only simulations experienced a smaller degree of mixing, primarily due to the lack of both planetesimal-induced migration and accretion of planetesimal material from a broader range of radii. Caution should be used when extrapolating the results of the test simulations described above to all extrasolar planetary systems considered here as variations in the precise structure of the system may influence the final results of any simulation. However, based on the test simulations with and without planetesimals, we do not expect to see drastic differences in the net results of the current study once planetesimals are included. The increase in radial mixing is not substantial enough to drastically alter the bulk planetary 114

Table 4.5: Statistical analysis of the embryo and planetesimal and embryo only extrasolar planetary system simulations. Number refers to the median number of terrestrial planets produced. σ is the median value of the radial mixing parameter (see text for deﬁnition). T50% and T90% is the median time required for the planet to accrete 50% and 90% of its ﬁnal mass respectively. Embryos and Embryos Planetesimals Only Number 1 2 σ 0.35 0.24 6 5 T50% 1.80×10 7.55×10 6 6 T90% 3.63×10 1.51×10

abundances produced by the current approach and the difference in planetary size and number is highly dependant on the exact system architecture. Thus I feel confident in the approach applied here where system-specific simulations are run with a population of embryos only. For each extrasolar planetary system modeled, planetary embryos are distributed in the zone between the star and the giant planets (or in the case where there are one or more giant planets close to the host star, in the region between the inner and outer giant planets) according to the relations between embryo mass, spacing, and orbital radius given by Kokubo and Ida (2000). No embryos are initially located interior to 0.3 AU. The timestep for the integration was set to at least 20× the orbital period of the innermost planet or planetary embryo, or the orbital period of a body at 0.1 AU, whichever is smaller (this corresponds to a 1 day timestep for an inner radius of 0.1 AU), and the simulations are run for 100-250 Myr. Surface mass density profiles that vary as r−3/2, normalized to 10 gcm−3 at 1 AU, were again assumed. For each system, a minimum of 4 accretion simulations were run. Migration of the giant planets is very likely to have occurred in all of these systems. However, if migration occurred very early, prior to planetesimal and embryo formation, then terrestrial planets could have potentially formed after migration, with the giant plan- 115 ets in their current configurations (e.g. Armitage 2003). Our simulations focus on this scenario (termed here “in-situ formation”). If giant planet migration occurred after planetesimals and embryos have formed, then our in-situ assumption does not apply and there is likely to be radial migration of planetary embryos, driven by giant planet migration (eg. Raymond et al. 2006; Mandell et al. 2007). However, as there is currently no clear consensus as to the most common timing of planetary migration, and no evidence for the specific systems that I propose to study, each of our simulations begin with the gas giants already fully formed and located in their current positions.

4.3.3 Chemical Simulations

Once again, the chemical composition of material within the disk is assumed to be determined by equilibrium condensation within the primordial stellar nebula. Equilibrium condensation sequences for an identical list of elements as used in Chapter 3 (H, He, C, N, O, Na, Mg, Al, Si, P, S, Ca, Ti, Cr, Fe and Ni) were obtained from the commercial software package HSC Chemistry (v. 5.1) using the same list of gaseous and solid species as previously listed in Table 3.2. Observed stellar photospheric abundances were adopted as a proxy for the composition of the stellar nebula. Given the deviations observed in the stellar elemental abundances of Chapter 2, elemental abundances for this study were taken from Gilli et al. (2006); Beirao˜ et al. (2005); Ecuvillon et al. (2004, 2006b). The input values used in HSC Chemistry for each system are shown in Table 4.6. All species are initially assumed to be in their elemental and gaseous form and no other species or elements are assumed to be present within the system. It should be noted that all abundances applied here were taken from the same research group and were obtained from the same spectra thus acting to limit any possible systematic differences in abundances due to instrument or method- ological differences between various studies. Similarly, stellar abundances were taken from the references speciﬁed instead of using the values determined in Chapter 2 due to the deviation in [Mg/H] values previously discussed. 116

The solar O abundance itself has experienced a ‘crisis’ in recent years (Ayres et al., 2006) with several studies suggesting that a downward revision of the solar O value is required (e.g. Ayres et al. 2006; Socas-Navarro and Norton 2007). This in turn would act to drastically alter the observed O abundances for extrasolar planetary host stars as all abundances are currently scaled against the solar value. However, realistic errors for the suggested new O abundance are approximately 0.1 dex (Socas-Navarro and Norton, 2007) and the abundances of the present study vary from [O/H] = 0.38 (HD 177830) to [O/H] = −0.21 (HD 142415), a range of 0.59 dex. As our present study range is almost six times larger than the errors of the revised O abundance, I feel confident in the compositional variations identified here as being caused by variations in the stellar O abundances of specific systems. Neither N or P abundances have been obtained for extrasolar planetary host stars, primarily due to the difficulty in finding unblended spectral lines to use within the visual spectral range (where most studies have been focussed). For this present study, I over- came this issue by obtaining approximate abundances for both elements based on the well known odd-even effect. Caused by the increased stability of even atomic number nuclei relative to odd-numbered nuclei, this effect produces the observed sawtooth pattern in the Solar elemental abundances. As both Na and P are odd-numbered nuclides, extrasolar abundances were obtained by fitting a linear trend through the solar abundances of odd nuclides and then applying this same fit to observed extrasolar host star abundances of Na and Al (odd nuclides). This approach assumes that extrasolar host stars will display the same atomic sawtooth pattern, a valid assumption as host stars do not appear to have undergone any form of systematic processing (such as pollution by a nearby supernova event) to cause a significant deviation (see Chapter 2). 117 8 8 8 6 7 6 7 5 7 6 5 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 2.34 0.58 2.82 2.09 5.01 3.47 4.57 3.39 1.82 2.75 1.62 6.61 4.57 2.57 1.00 8.51 8 8 8 6 7 6 7 5 7 6 5 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 7.41 1.82 3.47 8.13 6.61 7.95 2.82 1.74 1.62 5.89 6.03 4.17 1.00 8.51 10.96 12.30 8 8 8 6 7 6 7 5 7 6 5 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 3.98 0.98 3.63 2.00 6.31 4.68 4.47 4.57 1.91 2.09 1.26 6.17 4.79 2.57 1.00 8.51 8 8 8 6 7 6 7 5 7 5 6 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 1.91 0.47 5.89 1.74 4.37 2.75 3.63 2.69 0.81 1.07 1.91 4.27 3.02 1.74 1.00 8.51 8 8 8 6 7 6 7 5 7 5 6 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 6.03 1.48 5.89 4.47 5.50 4.90 5.50 4.79 1.23 1.20 3.02 6.92 4.90 3.16 1.00 8.51 8 8 8 6 7 6 7 5 7 5 6 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 2.40 0.59 3.31 0.89 2.57 1.86 1.91 1.82 0.58 0.58 1.00 2.29 1.62 0.95 1.00 8.51 8 8 8 6 7 6 7 5 7 5 6 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 6.92 1.70 4.57 3.89 7.59 8.13 7.42 2.19 2.04 3.55 9.33 7.08 4.47 1.00 8.51 10.23 8 8 8 6 7 6 7 5 7 5 6 5 7 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 5.25 1.29 7.59 2.69 7.24 5.13 5.62 5.01 1.66 1.74 6.46 4.90 3.02 2.63 1.00 8.51 8 8 8 6 7 6 7 5 5 7 5 7 6 6 12 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 10 × × × × × × × × × × × × × × × × 55Cnc Gl777A HD4203 HD4208 HD19994 HD72659 HD108874 HD177830 HD142415 P 6.76 S 1.82 C 5.01 H 1.00 NO 1.23 6.17 Si 6.31 Ti 1.82 Al 6.92 CrFeNi 7.24 6.03 3.47 Ca 2.45 He 8.51 Na 2.69 Mg 10.23 Element System Table 4.6: HSC Chemistry inputof values for HSC the Chemistry extrasolar planetary their systemslisted elemental studied. in and Tables All 4.3 gaseous inputs and are state. 4.4. entered into All the values simulations are in moles and are based on the stellar abundance values 118

As for the Solar System simulations of Chapter 3, “nominal” radial pressure and temperature proﬁles obtained from the Hersant et al. (2001) model were used to provide chemical compositions with a spatial location within the disk. The proﬁles were altered for the stellar mass of the host star obtained from the Simbad database1. Disk mass is assumed to vary linearly with stellar mass. From Hersant et al. (2001), disk mass depends on the mass accretion rate as roughly given by:

˙ 2/3 4/3 Mdisk ∝ M R (4.3)

˙ where Mdisk is the mass of the disk in solar masses, M is the mass accretion rate in solar masses per year and is the radius of the disk in AU. Adopting a uniform initial disk radius, the mass accretion rate for each of the extrasolar planetary systems can thus be scaled with the mass of the host star via:

˙ 3/2 M ∝ Mstar (4.4)

The resulting stellar accretion rates obtained are shown in Table 4.7. All other input parameters for the Hersant et al. (2001) models (α and initial disk radius) remained unchanged. It is important to note that the current approach does not include variations in the midplane conditions produced by different chemical compositions (which would alter parameters such as disk viscosity), nor does it include the effects of stellar luminosity (which aren’t incorporated into the Hersant et al. (2001) model). Furthermore, opacity changes due to varying disk composition are also likely to alter the midplane conditions. As such, the scaling applied here is a simplistic approach to a complex issue but is valid

1accessed at http://simbad.u-strasbg.fr/simbad/ 119 for the current aims of this study. For these reasons, the stellar masses considered were restricted to those close to solar (ranging from 0.95 MJ to 1.48 MJ) to prevent further complications due to widely varying stellar and disk masses. Please see Section 3.2.2 for a more detailed discussion of the Hersant et al. (2001) model. Midplane pressure and temperature values were determined with an average radial separation of 0.01AU throughout the study region. As in Chapter 3, an ensemble of planetary compositions were determined based on Hersant et al. (2001) disk conditions at seven different evolutionary times (t = 2.5×105yr, 5×105yr, 1×106yr, 1.5×106yr, 2×106yr, 2.5×106yr, 3×106yr). How- ever, as I previously found the best ﬁt to known planetary values in the Solar System to occur using disk conditions obtained for t = 5×105yr, our discussions will mostly focus on compositions produced by disk conditions at this time. The midplane temperature and pressure proﬁles for each system and “snapshot” time are displayed in Figures C.1 - C.9.

4.3.4 Combining Dynamics and Chemistry

The dynamical and chemical simulations were combined together as outlined in Chapter 3 whereby I assigned each embryo a composition based on its formation location and assumed that it then contributed that same composition to the ﬁnal terrestrial planet. Phase changes and outgassing were neglected and all collisions were assumed to be ideal (i.e. no mass loss occurred).

4.3.5 Stellar Pollution

The issue of stellar pollution produced by terrestrial planet formation is of great interest in extrasolar planetary systems. Pollution of the stellar photosphere via accretion of a large amount of solid mass during planet formation and migration has been suggested as a possible explanation for the observed metallicity trend for known host stars (Laughlin, 2000; Gonzalez et al., 2001; Murray et al., 2001). Thus I determined the amount of material accreted by the host star during the current terrestrial planet simulations and also determined the resulting change in spectroscopic photospheric abundance. As for 120

Table 4.7: Stellar accretion rates for the extrasolar host stars studied. Stellar masses were obtained from the Simbad database. Solar values are the nominal model determined by Hersant et al. (2001). See text for details on the scaling relations applied.

System Stellar Mass M˙ (MJ) (MJ/year) Solar 1.00 5.00×10−6

55Cnc 1.03 5.23×10−6

Gl777 1.04 5.30×10−6

HD4203 1.06 5.46×10−6

HD4208 0.93 4.48×10−6

HD19994 1.35 7.84×10−6

HD72659 0.95 4.63×10−6

HD108874 1.00 5.00×10−6

HD177830 1.48 9.00×10−6

HD142415 1.09 5.69×10−6 121

Table 4.8: Convective zone masses for each of the target stars. Teff values taken from Santos et al. (2004). See text for details on the determination of MCZ .

System Teff MCZ (K) (MJ) 55Cnc 5279 0.0398

Gl777 5584 0.0316

HD4203 5636 0.0288

HD4208 5626 0.0251

HD19994 6190 0.0045

HD72659 5995 0.0112

HD108874 5596 0.0321

HD142415 6045 0.0089

HD177830 4804 0.0501

the Solar System simulations, any solid material migrating to within 0.1AU from the host star is assumed to have accreted onto the stellar photosphere. This material is then assumed to have been uniformly mixed throughout the stellar photosphere and convective zone. Decreasing convective zone mass with time, granulation within the photosphere and gravitational settling and turbulence within the convective zone are again neglected, resulting in the values determined here being the maximum expected enrichments. The mass of each element accreted was determined in the same way as for terrestrial planets. As a reminder, the resulting photospheric elemental abundance is given by:   f  X  [X/H] = log J (4.5) fX, 122

where [X/H] is the resulting abundance of element X after accretion of terrestrial planet material, fX is the mass abundance of element X in the stellar photosphere after accretion J and fX, is the mass abundance of element X in the Sun (from Murray et al. (2001)). J Note that [X/H] is still dependant on fX, as by deﬁnition it is taken relative to the Solar abundance.

fX values for the extrasolar planetary host stars were obtained via the stellar abundances of Gilli et al. (2006); Beirao˜ et al. (2005); Ecuvillon et al. (2004, 2006b), as these papers represent a comprehensive, internally consistent catalogue of photospheric abundances for a large number of known planetary host stars. The mass of the convective zone of a star is known to vary with its mass, effective temperature (Teff ) and, to some extent, its metallicity. Values for the masses of the convective zone for each of the target stars was thus obtained from Pinsonneault et al. (2001) using the Teff values from Santos et al. J (2004). The convective zone masses are shown in Table 4.8. fX, values were obtained by utilizing the solar abundances of Asplund et al. (2005) and a current solar convective zone mass of 0.03MJ (Murray et al., 2001).

4.4 Results

4.4.1 Dynamical

Terrestrial planets were found to form in all simulations. 17 of the 40 simulations produced two or more terrestrial planets within one system.The general architecture of the resulting systems is shown in Figures 4.7 - 4.11. Several of these planets (e.g. those in the simulation for HD4203) can be seen to simply be embryos that have survived for the duration of the simulation but have not accreted any additional material. The median number of planets produced, along with their median masses, semimajor axes, eccentricities and inclinations are summarized in Table 4.9. 123

Of the nine planetary systems examined, only one (HD72659) is found to produce terrestrial planets with a median mass comparable to 1 ML with six of the 11 planets produced having masses equal to or greater than 1 ML. All other median planetary masses are less than 1 ML and only four out of 51 simulated planets have masses equal to or greater than 1 ML (excluding those of HD72659). Similarly, with the exception of 55Cancri, all simulated planets have median semimajor axes less than 1AU. This is believed to be a selection effect as I am currently only simulating terrestrial planet formation interior to the known giant planets (i.e. between the host star and the most distant known giant planet). Given that the majority of systems studied are known to contain giant planets orbiting within 2AU of their host star, it is thus expected that the terrestrial planets produced would be located in small orbits. 55Cancri is unusual in that its outermost giant planet has a periapse larger than 5AU, a signiﬁcant increase over the other systems selected for study. This in turn dictated that the embryos in the current simulations initially be located between 1 AU and 5AU, hence the larger median semi-major axis. Thus it can be seen that the present simulations are in general producing small terrestrial planets orbiting close to their host star. It is also interesting to note that no radial trends in planetary mass or orbital parameters can be seen. Similarly, no such trends with any of the orbital parameters and stellar metallicity can be seen. However, as the dynamical simulations are all currently adopting the same mass and surface density distribution, the stellar metallicity is not yet fully incorporated into the current simulations. As such, drawing any general conclusions based on the lack of radial trends in the current simulations would be premature. The planets themselves tended to accrete the vast majority of their mass from their immediately surrounding area with only small amounts of radial mixing occurring. Mean and median σ values are shown in Table 4.10. Only Gl777 and HD72659 display appreciable amounts of radial mixing with all other systems producing σ values well below the values observed for the Solar System simulations of O’Brien et al. (2006) and discussed in Chapter 3. As such, the terrestrial planets in the current simulations are forming pri- 124

Final Planetary Systems - 55 Cnc

0.66

0.50

0.64 0.25

0.47 0.8 0.2 0.1 MJ 3.84 MJ

0 1 2 3 4 5 6 7 Semimajor Axis (AU)

Final Planetary Systems - Gl 777

0.80

0.94 0.40

0.47 1.10

1.03 0.06 MJ 1.55 MJ

0 1 2 3 4 5 Semimajor Axis (AU)

Figure 4.7: Schematic of the results of the dynamical simulations for 55Cancri (top panel) and Gl777 (bottom panel). Known giant planets are also shown with their masses in Jupiter masses (MJ ). The horizontal lines indicate the range in distance from apastron to periastron. The vertical lines indicate variation in height above the midplane due to orbital inclination. Numerical values represent the mass of the planet in Earth masses. 125

Final Planetary Systems - HD4203

0.17

0.04

0.04 2.1 MJ

0 0.5 1 1.5 2 Semimajor Axis (AU)

Final Planetary Systems - HD4208

0.76 0.72

0.34 0.63 0.54

0.35 1.18 0.49

1.58 0.10 0.13 0.8 MJ

0 0.5 1 1.5 2 Semimajor Axis (AU)

Figure 4.8: Schematic of the results of the dynamical simulations for HD4203 (top panel) and HD4208 (bottom panel). Known giant planets are also shown with their masses in Jupiter masses (MJ ). The horizontal lines indicate the range in distance from apastron to periastron. The vertical lines indicate variation in height above the midplane due to orbital inclination. Numerical values represent the mass of the planet in Earth masses. 126

Final Planetary Systems - HD19994

0.57

0.62

0.35 0.10 0.06

0.28 0.46 1.7 MJ

0 0.5 1 1.5 2 Semimajor Axis (AU)

Final Planetary Systems - HD72659

1.53 1.35

0.60 1.10 1.03

1.28 0.99 0.26

0.441.32 0.71 3.3 MJ

0 1 2 3 4 5 Semimajor Axis (AU)

Figure 4.9: Schematic of the results of the dynamical simulations for HD19994 (top panel) and HD72659 (bottom panel). Known giant planets are also shown with their masses in Jupiter masses (MJ ). The horizontal lines indicate the range in distance from apastron to periastron. The vertical lines indicate variation in height above the midplane due to orbital inclination. Numerical values represent the mass of the planet in Earth masses. 127

Final Planetary Systems - HD108874

0.34

0.46

0.18

0.40 0.13 1.3 M J 1.07 MJ

0 0.5 1 1.5 2 2.5 3 Semimajor Axis (AU)

Final Planetary Systems - HD142415

0.14

0.04

0.13

0.05 1.7 MJ

0 0.2 0.4 0.6 0.8 1 1.2 Semimajor Axis (AU)

Figure 4.10: Schematic of the results of the dynamical simulations for HD108874 (top panel) and HD142415 (bottom panel). Known giant planets are also shown with their masses in Jupiter masses (MJ ). The horizontal lines indicate the range in distance from apastron to periastron. The vertical lines indicate variation in height above the midplane due to orbital inclination. Numerical values represent the mass of the planet in Earth masses. 128

Final Planetary Systems - HD177830

0.78 0.24

0.35 0.61 0.14

1.22 0.06

0.360.24 0.34 1.4 MJ

0 0.25 0.5 0.75 1 1.25 1.5 Semimajor Axis (AU)

Figure 4.11: Schematic of the results of the dynamical simulations for HD177830. Known giant planets are also shown with their masses in Jupiter masses (MJ ). The horizontal lines indicate the range in distance from apastron to periastron. The vertical lines indicate variation in height above the midplane due to orbital inclination. Numerical values represent the mass of the planet in Earth masses. 129

Table 4.9: Properties of the simulated extrasolar terrestrial planets for each system. Me- dian values are provided for the number of planets produced (N), planetary mass (M), semi-major axis (a), orbital eccentricity (e) and orbital inclination (i). The range in values for each system is also provided in parentheses.

System N M a e i (ML) (AU) (◦) 55Cancri 1 0.51 2.00 0.15 6.28 (0.67 − 0.25) (3.59 − 1.65) (0.30 − 0.11) (8.19 − 0.50)

Gl777 1.5 0.87 0.59 0.12 4.58 (1.11 − 0.40) (0.89 − 0.45) (0.29 − 0.07) (16.18 − 0.33)

HD4203 1 0.04 0.32 0.31 1.22 (0.17 − 0.04) (0.36 − 0.28) (0.39 − 0.27) (2.43 − 0.29)

HD4208 3 0.54 0.45 0.12 3.71 (1.58 − 0.10) (1.19 − 0.27) (0.28 − 0.05) (8.91 − 1.59)

HD19994 1.5 0.35 0.37 0.12 2.02 (0.63 − 0.07) (0.70 − 0.31) (0.19 − 0.10) (5.38 − 1.07)

HD72659 3 1.03 0.49 0.11 3.95 (1.53 − 0.26) (0.97 − 0.29) (0.22 − 0.07) (8.87 − 1.46)

HD108874 1 0.34 0.35 0.16 2.19 (0.47 − 0.13) (0.48 − 0.33) (0.19 − 0.13) (3.08 − 0.77)

HD142415 1 0.04 0.30 0.32 0.97 (0.14 − 0.04) (0.33 − 0.25) (0.41 − 0.23) (1.49 − 0.41)

HD177830 2.5 0.35 0.45 0.06 1.72 (1.22 − 0.06) (0.65 − 0.33) (0.17 − 0.02) (15.63 − 0.66) 130

Table 4.10: Degree of radial mixing for simulated systems. σ is the weighted measure of radial migration deﬁned in equation 4.2. The error on the mean is the rms error. Values for the Solar System simulations of O’Brien et al. (2006) are provided for comparison.

System σ Median Mean 55Cnc 0.16 0.19 ± 0.22

Gl777 0.43 0.49 ± 0.52

HD4203 0.01 0.04 ± 0.08

HD4208 0.32 0.32 ± 0.32

HD19994 0.15 0.15 ± 0.15

HD72659 0.40 0.41 ± 0.41

HD108874 0.09 0.09 ± 0.09

HD142415 0.04 0.09 ± 0.13

HD177830 0.18 0.17 ± 0.20

Solar System - CJS 0.56 0.59 ± 0.59

Solar System - EJS 0.45 0.48 ± 0.48

marily from material located in the region immediately surrounding them and are thus expected to have compositions reflecting any radial trends within the disk. These values, however, are likely to increase significantly once the full effects of a planetesimal swarm and orbital migration are incorporated into the planet formation simulations. Terrestrial planets produced in the current simulations attain their final masses relatively quickly. Median times to accrete 50% and 100% of the final planetary mass are shown in Table 4.11. As one would intuitively expect, systems with the highest median planetary masses (HD72659, Gl777) also had the highest median growth times (approxi- 131

Table 4.11: Median time required to accreted 50% and 100% of the ﬁnal terrestrial planet mass. N is the number of planets considered in determining the median values. Only planets that accreted two or more embryos are included.

System N Time (years) M50% M100% 55Cnc 4 3.70 ×105 3.34 ×106

Gl777 6 2.18 ×106 1.11 ×107

HD4203 1 3.03 ×102 1.63 ×104

HD4208 9 4.79 ×105 2.36 ×106

HD19994 5 1.52 ×105 6.92 ×105

HD72659 11 2.60 ×106 9.35 ×106

HD108874 4 1.02 ×105 4.05 ×105

HD142415 2 2.35 ×103 1.21 ×104

HD177830 8 3.94 ×105 1.53 ×106

mately 10Myr for the final planetary mass). Systems with lower median masses reached their final planetary masses significantly faster (3Myr or less). Planetary formation times have implications for planetary processes (such as differentiation and the resulting the distribution of siderophile elements). The current growth times, however, are expected to change once planetesimals are incorporated into the simulations as discussed in Section 4.3.2. The above results need to be interpreted with caution as I am currently only considering late stage, in-situ terrestrial planet formation. That is, I am only considering formation that has occurred after the known giant planets have formed and migrated to their current orbits. As previously discussed, such an approach is currently valid as there is no con- 132 sensus on the timing or extent of migration within these systems and it provides us with a valid starting point to consider the possible chemical composition of the system. How- ever, migration could potentially remove sufficient mass from the planetary disk so as to inhibit any terrestrial planet formation from occurring. Alternatively, migration of the giant planets at a later time (i.e. after terrestrial planet formation has begun) may result in the ejection of terrestrial planets from the system and radial redistribution of material such that it deviates from the currently assumed surface density profile (Mandell et al., 2007). These effects are believed to be of most importance in systems with close in giant planets in very eccentric orbits, such as is the case for HD4203. Simulations addressing these issues are currently running.

4.4.2 Chemical

The condensation sequences and abundances of solid species (normalized to the abundance of the least abundant species) are shown in Figures 4.12 - 4.20 in order of increasing C/O value. The 50% condensation temperatures (i.e. temperature at which half of the species has condensed) for each of the systems studied is shown in Table 4.12. As they are not the main focus of the current work, the equilibrium gas compositions are shown in Figures D.1 - D.9 in Appendix D. Two very distinct general types of condensation sequence are produced for the systems studied here - those resembling the Solar condensation sequence (Gl777, HD4208, HD72659 and HD177830) and those in which C (and occasionally O) is drastically more refractory in nature (55Cnc, HD4203, HD19994, HD108874 and HD142415). This compositional difference can be seen in Figures 4.18 - 4.20 where the high temperature (innermost) region of the disk is dominated by SiC, TiC and C. Furthermore, the spatial region where solid material is present increases with increasing C/O value. For HD72659 (the system with lowest C/O value), solids are present at temperatures below ∼1600 K while for HD19994, solids are present below ∼1800 K. HD4203, the most C rich system studied, has solids present below ∼2300 K. The implications of these variations in the dis- 133 tribution of solid material are discussed in Section 4.5.9. The composition of terrestrial planets produced in each of these general classes reﬂects these drastic compositional differences and shall be discussed in turn. Unless otherwise stated, all compositions shown are produced by disk conditions at t = 5×105 years. Compositional changes with disk conditions will be discussed in Section 4.5.3 and compositions produced by disk conditions at alternative times are presented in Figures E.1 - E.32. 134 150 949 909 1021 923 1064 < 150 < 150 < , C/O and Mg/Si for extrasolar planetary systems studied. Solar values are also shown for compari- 150 < 150 condensation < 50% HD72659 Solar HD177830 Gl777A HD4208 55Cnc HD142415 HD19994 HD108874 HD4203 PS 1308 1039 618 658 1333 1338 682 1325 656 1372 602 1374 1070 1389 1052 1388 628 1403 1082 1091 C O 180 180 183 181 176 263 941 180 759 997 Si 1346 1329 1359 1321 1275 1137 1114 1430 1556 1637 Ti 1587 1580 1583 1577 1521 1539 1499 1771 1803 1824 CrFe 1295 1333 1301 1339 1321 1359 1313 1351 1272 1310 1320 1359 1310 1348 1313 1352 1304 1350 1323 1366 Ni 1345 1351 1371 1363 1321 1370 1359 1363 1358 1375 Al 1657 1639 1672 1640 1575 1594 1519 1298 1305 1323 Ca 1535 1527 1524 1522 1463 1482 1429 1253 1284 1295 Na 939 941 815 909 916 857 846 866 848 857 Mg 1355 1339 1362 1338 1286 1131 1112 1070 1061 1069 C/O 0.32 0.54 0.68 0.69 0.72 0.81 0.83 1.02 1.10 1.51 Mg/Si 1.20 1.05 1.86 1.29 1.35 1.62 1.10 1.00 1.41 1.26 Element System Table 4.12: T son. All values are in K. 135

HD72659 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.12: Schematic of the output obtained from HSC Chemistry for HD72659 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 136

HD177830 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.13: Schematic of the output obtained from HSC Chemistry for HD177830 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 137

Gl777 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.14: Schematic of the output obtained from HSC Chemistry for Gl777 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 138

HD4208 AlN C Ca2Al2SiO7 1e+5 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 1e+4 CaAl2Si2O8 CaTiO3 CaS FeCr2O4 1e+3 Cr Fe Fe2SiO4 Fe3C 1e+2 Fe3O4 Fe3P FeS FeSiO3 1e+1 H2O Mg2SiO4 Mg3Si2O5(OH)4 Normalized Abundance (mole) 1e+0 MgAl2O4 MgS MgSiO3 NaAlSi3O8 1e-1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.15: Schematic of the output obtained from HSC Chemistry for HD4208 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 139

55Cnc AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.16: Schematic of the output obtained from HSC Chemistry for 55Cnc at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 140

HD142415 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.17: Schematic of the output obtained from HSC Chemistry for HD142415 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 141

HD19994 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) Normalized MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.18: Schematic of the output obtained from HSC Chemistry for HD19994 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 142

HD108874 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 TiC T (K) TiN

Figure 4.19: Schematic of the output obtained from HSC Chemistry for HD108874 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 143

HD4203 AlN C Ca2Al2SiO7 10000 Ca3(PO4)2 CaMgSi2O6 CaAl12O19 CaAl2Si2O8 1000 CaTiO3 CaS FeCr2O4 Cr Fe 100 Fe2SiO4 Fe3C Fe3O4 Fe3P 10 FeS FeSiO3 H2O Mg2SiO4 Mg3Si2O5(OH)4 1 Normalized Abundance (mole) MgAl2O4 MgS MgSiO3 NaAlSi3O8 0.1 Ni SiC 200 400 600 800 1000 1200 1400 1600 1800 2000 2200 TiC T (K) TiN

Figure 4.20: Schematic of the output obtained from HSC Chemistry for HD4203 at a pressure of 10−4 bar. Only solid species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 144

Earth-like Planets

Four systems (Gl777, HD4208, HD72659 and HD177830) were found to produce condensation sequences (and thus also terrestrial planets) comparable to those of the Solar System. The final elemental abundances for all times studied are shown in Tables E.1 - E.4 while schematic representations of the abundances (for disk conditions at 5×105 years) are shown in Figures 4.21 − 4.26. From these it can be seen that for Gl777, HD72659 and HD4208 the terrestrial planets produced are grossly similar in composition to known terrestrial planets. Their compositions are dominated by O, Fe, Mg and Si with varying amounts of minor elements (such as Al, Ca and Cr). Upon closer examination, however, large and important differences emerge, primarily due to variations in the compositions of the host star and thus the initial system itself. Gl777: Gl777 produces the most Earth-like terrestrial planets of all of the systems simulated (Figure 4.21). The final elemental abundances of the refractory lithophile and siderophile elements are remarkably similar to those of Earth with the simulated planets displaying a marginal enrichment in Mg (∼3wt%), depletion in Si (∼1wt%), and Fe (∼2wt%) compared to Earth values. Deviations from Earth abundances also occur for the most volatile species (P, Na and S). This enrichment is likely an artifact of the fact that I do not consider volatile loss during accretion, which is expected to be significant as for the Solar System simulations. The geochemical ratios of Mg/Si and Al/Si are also in excellent agreement with those of the Earth, falling on the Earth fractionation trend and within the upper limit of values estimated for the Earth (Drake and Righter, 2002) (see Figure 4.22). This increase in the planetary Mg/Si values over the previous Solar System simulations of Chapter 3 is due to J the slight increase the Mg/Si value of Gl777 itself (Mg/Si = 1.05, Mg/SiGl777 = 1.29) . In turn, this produces a system containing nearly equal amounts of olivine and pyroxene (compared to the pyroxene dominated Solar disk) and thus results in Mg-enriched planets.

The Ca/Si values, however, are lower than those of Earth (Ca/SiGl777 = 0.06−0.07, 145

Ca/SiL = 0.11), primarily due the fact that there is relatively less Ca within the system. Ca and S are the two least enriched elements within Gl777 ([Ca/H] = 0.10 vs. [Al/H] = 0.34), resulting in a relative Ca and S depletion within the solid material. The variation in the abundances of the host star are reflected in the lower Ca/Si value of the final planets produced. This difference is certainly no larger than that observed for the Solar Sys- tem simulations previously discussed and I feel confident in claiming that the terrestrial planets of Gl777 are essentially Earth-like in their chemical composition. It is also interesting to note Gl777 represents the average extrasolar planetary host star values of Mg/Si and C/O (1.29 and 0.63 respectively). This result thus implies that the “average” extrasolar planetary system contains terrestrial planets with compositions extremely comparable to that of our own Solar System. HD72659 & HD4208: More pronounced compositional variations can be seen in the terrestrial planets of HD72659 and HD4208 (Figures 4.23 and 4.24). Although producing planets with two of the highest radial mixing parameters for the extrasolar planetary simulations, clear radial compositional trends are evident. Planets located within ∼0.7AU from the host star for HD72659 and within ∼0.5AU for HD4208 are primarily composed of Al, Ca and O, indicating that these planets formed from the high-temperature Al and Ca condensates (such as spinel and gehlenite) (see Figures 4.23 and 4.24). However, beyond ∼0.7AU and ∼0.5AU respectively, planets have compositions more closely correlating with that of the Earth, dominated by O, Fe, Mg and Si. As expected, this difference is reflected in the planetary geochemical ratios as the planets located within the inner region have Al/Si and Ca/Si ratios well above and Mg/Si values well below observed Solar System terrestrial planet values. However, for planets located beyond the compositional transitional point, the reverse is true. In this region, planetary Al/Si and Ca/Si values are below Earth values while Mg/Si is within the upper limits of current Earth approximations (for HD72659) and in agreement with terrestrial peridotites (for HD4208). A steady transition between these two regions is seen for both systems (see Figure 4.25). This trend lies well above the observed Earth fractionation line 146 and is a result of the condensation of Mg-silicate species further out in the disk, combined with relatively little radial mixing of material during the formation process. The terrestrial planets of HD72659 and HD4208 can thus be characterized as being essentially similar in composition to Ca- and Al-rich inclusions (CAI’s) (for the inner planets) and Earth (for the outer planets). HD177830: HD177830 has the highest Mg/Si (and Al/Si) ratio of any system simulated. This enrichment alters the compositions of major silicate species present within the disk. While the Solar System should have condensed both olivine and pyroxene between 0.35 and 2.5 AU, HD177830 is dominated by olivine beyond 0.3 AU and contains only a small region where pyroxene is predicted to coexist. This unusual composition is reflected in the final planetary abundances as the planets contain large portions of Mg (up to 23 wt%) (Figure 4.26) and have a mean Mg/Si value of 1.71, well above Earth values (Mg/SiL = 1.01). Al is also similarly enriched (up to 18.07 wt%), again because of the high Al abundance of the host star and thus the initial system itself. Other refractory and lithophile elemental abundances within the final planets are comparable to that of the Solar System. The planets of HD177830 can best be described as being Mg- and Al-rich Earths. Such a Mg dominated planetary composition would undoubtedly alter the interior structure and processes of the planets themselves. Such considerations will be discussed further in Section 4.5. 147

Final Composition - Gl777 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Figure 4.21: Schematic of the bulk elemental planetary composition for Gl777. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 148

1.4 Earth fractionation line

1.2

1.0

0.8 ratio) ght 0.6

0.4 Mg/Si (wei Mars 0.2 fractionation line

0.0 0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6

Al/Si (weight ratio)

Figure 4.22: Al/Si v. Mg/Si for planets of Gl777. Black circles indicate values for disk conditions at t = 2.5×105 years while red circles indicate values for disk conditions at t = 5×105 years. Values at all other times are concentrated at the 5×105 years values and omitted for clarity. Earth values are shown in green and are taken from Kargel and Lewis (1993) and McDonough and Sun (1995). Martian values are shown in pink and are taken from Lodders and Fegley (1997). Venus values are shown in light blue and are taken from Morgan and Anders (1980). 149

Final Composition - HD72659 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Figure 4.23: Schematic of the bulk elemental planetary composition for HD72659. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 150

Final Composition - HD4208 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca

Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Figure 4.24: Schematic of the bulk elemental planetary composition for HD4208. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 151

1.4 Earth fractionation line

1.2

1.0

0.8 ratio) ght 0.6

0.4 Mg/Si (wei Mars 0.2 fractionation line

0.0 012345

Al/Si (weight ratio)

Figure 4.25: Al/Si v. Mg/Si for the planets of HD4208 and HD72659. Black circles indicate values for the terrestrial planets of HD72659 while red circles indicate values for the terrestrial planets of HD4208. Values are for disk conditions at 5×105 years. Earth values are shown in green and are taken from Kargel and Lewis (1993) and McDonough and Sun (1995). Martian values are shown in pink and are taken from Lodders and Fegley (1997). Venus values are shown in light blue and are taken from Morgan and Anders (1980). 152

Final Composition - HD177830 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Figure 4.26: Schematic of the bulk elemental planetary composition for HD177830. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 153

C-rich Planets

Far more profound mineralogical variations from Earth-like compositions occur in systems with C/O values above 0.8. The inner part of these systems are dominated by refractory carbon species such as graphite, SiC and TiC (e.g. see Figures 4.18 - 4.20). Five such systems were selected for the current study (55Cnc, HD142415, HD19994, HD108874 and HD4203). The final elemental abundances for all times studied are shown in Tables E.5 - E.9 while schematic representations of the abundances are shown in Figures 4.27 − 4.31. These planets clearly represent a completely different type of terrestrial body unseen in our Solar System. 55Cnc: For disk conditions at t = 5×105 years, 55Cancri produced terrestrial planets similar to those of HD177830 - essentially Earth-like in terms of major elements present but with Mg/Si and Ca/Si values well above those of Earth and Ca/Si values well below. (see Figure 4.27) This high planetary Mg abundance is caused by the fact that 55Cancri is highly enriched in Mg ([Mg/H] = 0.48), resulting in olivine becoming the major silicate species present within the disk and thus producing the high Mg/Si value observed. Although the disk of 55Cancri is predicted to contain a C-rich zone, only one predicted planetary composition for disk conditions at t = 5×105 years contains a significant amount of C (see Figure 4.16). For disk conditions at later times, none of the simulated planets are predicted to contain any C. This apparent C depletion is due to the location of the planets within the disk. All of the terrestrial planets in 55Cnc are located between 1.5 and 4AU while the C-rich zone is located between ∼1100 and 750K, corresponding to a radial distance of 0.6 and 1.35AU (for disk conditions in the Hersant et al. (2001) model at 5×105 years). Thus the primary feeding zones for each of the planets are located beyond the C zone, resulting in Mg-rich Earth-like planets being produced. For disk conditions at 2.5×105 years, however, the C zone extends from 0.92 to 2.11AU, producing planets that contain up to 9.41 wt% C. The location of the C zone also implies that the four inner known giant planets of the 55Cnc system (located at 0.038AU, 0.115AU, 0.24AU 154 and 0.781AU) should contain significant amounts of C, both in their solid cores and in their atmospheres (depending on the exact time of their formation). Given the variation in the location of the C rich zone, it is expected that C rich terrestrial planets would also be produced in the current simulations at later times if temporal variations in solid composition were incorporated. HD142415: Although the disk of HD142415 does contain a C-rich zone, for the disk conditions at t = 5×105 years the simulated terrestrial planets consist almost entirely of refractory siderophile and lithophile species (Ti, Al, Ca and O) (see Figure 4.28). As for 55Cnc, this is primarily due to their radial location within the system. All of the planets of HD142415 are located within 0.35AU from their host star, well inside the C zone which extends from 0.6 to 1.3AU (for disk conditions at 5×105 years). Thus their location, combined with the small radial mixing observed for the HD142415 simulation, results in planets consisting entirely of refractory species, resembling the CAI’s of the Solar System. At later times, however, the planetary composition changes to become more Earth- like, with planets dominated by O, Fe, Mg and Si and a significant amount of C. Up to 12.44 wt% C is predicted to exist in the planets for the disk conditions at 3×106 years. These planets are essentially C-enriched Earths, containing the same major elements in the same geochemical ratios as Earth, but also an enhanced inventory of C, primarily accreted as solid graphite. As for 55Cnc, it is expected that were I to incorporate giant planet migration and temporal variations in composition into our models that I would see C occurring in the terrestrial planets for all simulation times. HD19994, HD108874 & HD4203: HD19994, HD108874 and HD4203 all have C/O values above 1.0 (1.02, 1.10 and 1.51 respectively). In all three systems, the inner regions of the disk are completely dominated by refractory species composed of C, SiC and TiC, as opposed to the Ca and Al-rich inclusions characteristic of the earliest solids within our Solar System. Significant amounts of metallic Fe are also present within these systems. As all three systems produced terrestrial planets located within 0.7AU from their host star, 155 these unusual inner disk compositions produced terrestrial planets primarily composed of C, Si and Fe. HD19994 produced terrestrial planets composed almost entirely of SiC and metallic Fe and containing up to 60 wt% Si and between 16 and 21 wt% C, over 100 times more C than is estimated for Earth (see Figure 4.29). The outermost terrestrial planet for HD19994 does contain significant amounts of O and Mg, primarily as its feeding zone, although still undoubtedly dominated by C, is also rich in pyroxene. This presence of a Mg silicate species produces a slightly more varied composition for a single planet. More extreme deviations occur when I consider the planets formed for HD108874 and HD4203. Both of these systems have considerably wider graphite dominated regions, extending from 1.5AU to within 0.1AU (for disk conditions at 5×105 years). As a result, terrestrial planets are found to be composed almost entirely of C, Si and Fe. HD108874 produced terrestrial planets containing between 9.58 and 29.30 wt% C, 13.29 and 51.36 wt% Si and 18.03 and 63.62 wt% Fe (see Figure 4.30). HD4203 was even more extreme, producing terrestrial planets composed almost entirely of SiC and containing more than 50 wt% C (for midplane conditions at 5×105 years) (see Figure 4.31). At later times, Mg and O are also present, again due to the incorporation of pyroxene and olivine into the planetary feeding zones. It must be noted though that the terrestrial planets formed in the HD4203 simulations are single embryos that survived for the duration of the simulation but did not accrete any other solid material. As such, they are presumably more C-rich than terrestrial planets forming for other systems as they have not been combined with material drawn from any other region within the disk. Terrestrial planets within these systems are unlikely to have compositions resembling that of any body we have previously observed. The possible implications of these types of planetary compositions will be discussed in Section 4.5.

4.4.3 Stellar Pollution

The average change in stellar photospheric abundances produced by accretion for disk conditions at 5×105 years are shown in Table 4.13. The majority of systems experienced 156

Final Composition - 55 Cnc (0.5 Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Figure 4.27: Schematic of the bulk elemental planetary composition for 55Cnc. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 157

Final Composition - HD142415 (0.50Myr)

O Sim. 1 Fe M g

Sim. 2 Si C S Sim. 3 Al Ca Ti Sim. 4 O ther

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure 4.28: Schematic of the bulk elemental planetary composition for HD142415. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 158

Final Composition - HD19994 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 Semimajor Axis (AU)

Figure 4.29: Schematic of the bulk elemental planetary composition for HD19994. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 159

Final Composition - HD108874 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C

Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure 4.30: Schematic of the bulk elemental planetary composition for HD108874. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 160

Final Composition - HD4203 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Figure 4.31: Schematic of the bulk elemental planetary composition for HD4203. All values are wt% of the final simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. 161 minimal increases in photospheric abundances as a result of accretion during terrestrial planet formation. The largest elemental enrichments occurred for the most refractory elements (Al, Ca, Ti, Ni and Cr), primarily as a result of accretion of refractory material initially located closest to the star itself, as one would intuitively expect. Only the simulations for HD142415 produced drastic increases in the predicted observed elemental abundances, with simulations for HD19994 also producing significant enrichments for Ti and Si. This is primarily due to the low estimated mass of the convective zone masses for J J the host stars (MHD142415 = 0.0089M ,MHD19994 = 0.0045M ). However, as previously discussed, mixing within the stellar radiative zone are not incorporated into the current approach. As such, our current values are upper limits for those stars with low mass convective zones and large radiative zones as is the case for HD142415 and HD19994. It is also interesting to note that the two systems with the highest degree of radial mixing (Gl777 and HD72659) both accreted the largest amount of solid material onto their host stars. With the exception of HD142415 and HD19994, all predicted abundance changes are below the errors of current spectroscopic surveys (e.g. ±0.03 for Fischer and Valenti (2005)), meaning that definitively observable elemental enrichments are not necessarily produced by the current terrestrial planet formation simulations. Of course, inclusion of planetesimals within the formation simulations and migration of the giant planets is expected to increase the amount of material accreted by the host star and thus also the predicted stellar abundances. However, these increases are expected to be no more than a factor of two and would thus still result in only marginal increases in the observed elemental abundances. 162 ) Mg O S Fe Al Ca Na Ni Cr P Ti Si C L (M Table 4.13: Change in host star photospheric abundances produced by terrestrial planet formation. 42034208 0.63 0.28 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.02 0.00 0.02 0.00 0.00 0.00 0.00 0.00 0.00 0.01 0.00 0.01 0.02 0.00 0.01 0.00 Gl777 2.101999472659 0.01 0.90 0.00 2.04 0.01 0.01 0.02 0.01 0.00 0.02 0.01 0.00 0.01 0.01 0.00 0.01 0.01 0.03 0.02 0.01 0.02 0.01 0.00 0.00 0.02 0.01 0.04 0.01 0.01 0.04 0.02 0.00 0.01 0.02 0.07 0.01 0.05 0.03 0.01 0.02 0.00 55Cnc 0.14 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 System Mass Accreted Change in Abundance 142415177830 0.55 0.24 0.00 0.01 0.00 0.00 0.00 0.00 0.00 0.16 0.00 0.15 0.00 0.00 0.00 0.01 0.00 0.00 0.00 0.00 0.00 0.29 0.00 0.01 0.00 0.00 0.00 0.00 163

4.5 Implications and Discussion

4.5.1 Frequency of Terrestrial Planets

As terrestrial planets formed in all systems studied, it implies that terrestrial planets are ubiquitous within extrasolar planetary systems. This observation is not limited to those systems dynamically similar to our own (i.e. with the innermost giant planet located at 5AU). Systems containing multiple giant planets, along with those containing close-in giant planets, are found to be capable of both forming and retaining at least one terrestrial- sized planetary companion. Furthermore, many of the systems simulated produced multiple terrestrial planets. As Jovian-mass planets in close-in orbits are the most commonly known type of extrasolar planetary system, it is encouraging to note that the current simulations are forming several terrestrial planets in such systems. Assuming sufﬁcient mass is retained after the known giant planets have formed and migrated to their current positions, these current simulations indicate that terrestrial planets are likely to form in a wide variety of planetary architectures. These results provide further motivation for the development of future terrestrial planet searches and argue for the expansion of such searches to include a diverse range of dynamical structures beyond those most closely resembling the Solar System.

4.5.2 Planetary Types

A broad range of possible planetary compositions have been produced. Four distinct classes of planetary composition can be identiﬁed: Earth-like, Mg-rich Earth-like, refractory (compositions similar to CAI’s) and C-rich. These planetary types are primarily a result of the compositional variations of the host stars and thus the system as a whole. Based on their observed photospheric elemental abundances, the majority of known extrasolar planetary systems are expected to produce terrestrial planets with compositions similar to those within our own Solar System. Therefore, systems with elemental abundances and ratios similar to these (e.g. Gl777) are ideal places to focus future “Earth-like” 164 planet searches. On the other hand, this study also suggests that a large fraction of planetary systems are much more C-rich than anything in our own system. This result implies that a similar fraction of protoplanetary disks should contain high abundances of carbonaceous grains. Min et al. (2005) have shown that for comets and protoplanetary disks a mass fraction of 20% carbon grains successfully reproduces both the infrared spectrum and the polariza- tion of scattered light at optical wavelengths. Furthermore, infrared spectral features at 3.43 and 3.53 µm observed in protoplanetary disks have been identiﬁed as being produced by nano-diamonds (Acke and van den Ancker, 2006). Such high abundances of carbon- rich grains in nascent planetary systems is inconceivable if they have primary mineralogy similar to our Solar System, thus implying that C-rich planetary systems may be more common than previously thought. Future planetary formation and astrobiological studies should consider the implications of a C-dominated composition.

4.5.3 Timing of Formation

Specific planetary compositions have been found to be highly dependent on the time selected for the disk conditions. This is primarily due to the low degree of radial mixing encountered within the simulations. As a result, as conditions within the area immediately adjacent to the planet evolve, the composition of the solid material and thus the final planet itself drastically change. For disk conditions at later times, it can been from Tables E.1 - E.9 that the compositions evolve to more closely resemble those of the So- lar System. They become dominated by Mg silicate species and metallic Fe. Terrestrial planets in Solar-like systems attain more hydrous material while those in C-rich systems accrete more C. Temporal variations in composition are most noteworthy for those planets dominated by refractory compositions (such as the inner planets of HD72659 and HD142415). Under later disk conditions, these planets experience a complete shift in their composition, losing the majority of their refractory inventory to be composed primarily of Mg-silicates (olivine and pyroxene). Therefore if solid condensation and planet forma- 165 tion occurred significantly later, I would expect to observe predominantly Mg-silicate and metallic Fe planets with enrichments in other elements (such as C) depending on the exact composition of the system. Although disk conditions at 5×105 years provide the “best fit” for Solar System simulations and are thus utilized here, it remains to be seen whether or not disk conditions at this time provide an accurate description of accretion conditions in other planetary systems. Therefore, we require a more detailed understanding of the timing of condensation and planetesimal and embryo formation within protoplanetary disks to be able to further constrain the predicted elemental abundances. Similarly, as the disk evolves, the various condensation fronts migrate closer to the host star. For example, the water ice line for Gl777 migrates from 7.29 AU for midplane conditions at 2.5×105 years to 1.48AU for midplane conditions at 3×106 years. Similar degrees of migration also occur for other species. In effect, this migration alters the mass distribution within the disk, concentrating more mass in both the very closest regions of the disk (< 1AU) and in the outer most water rich regions (> 3AU) producing a bi-modal mass distribution. The full effects of this change will obviously require formation simulations to be run with alternative mass distributions but it is thought that such conditions will increase the efficiency of forming close-in terrestrial planets and/or the mass of the resulting planets. Additionally, it will also allow for efficient terrestrial planet growth in the outer regions, possibly to the extent of forming gas giant cores. It remains to be seen, however, if sufficient solid mass would be retained during the evolutionary process for Jovian-cores to develop.

4.5.4 Detection of Terrestrial Planets

The results of this study are of great importance for the design of terrestrial planet finding surveys. Our simulations provide not only predictions of the location of terrestrial planets but also constrain their mass and bulk composition, thus aiding in detection. Based on the present simulations, the masses of the terrestrial planets produced are certainly too low to be detected by current radial velocity surveys. However, it is believed that 166 many of the simulated planets are in orbits that would place them within the prime target space for detection by the Kepler mission. Designed to detect extrasolar planets via transit studies, Kepler is the first mission which will be capable of detecting Earth-mass (and lower) extrasolar planets located within the habitable zone of a planetary system. It will be able to detect an Earth mass body within 2AU from the host star and a Mars mass body (0.1ML) within 0.4AU. The vast majority of the terrestrial planets formed here (with the exception of the lowest mass, highest semimajor axis planets) are well within this range and thus should be detectable if they are indeed present within these systems. Only HD4203 and HD142415 produce no potentially detectable planets, based on their predicted masses. Thus it is likely that I will have an independent check of extrasolar terrestrial planet formation simulations within the next 5 years. Such information will be vital for further refinement of planetary formation models for both giant and terrestrial planets. Obtaining compositional checks, however, will be more difficult as the size and location of the predicted planets will prohibit direct spectroscopic studies. It is also unlikely that the terrestrial planets will contain atmospheres large enough to be detected with transiting surveys. As such, specific extrasolar planetary chemical compositions will remain unknown for the foreseeable future. In addition to detection via transit surveys, attempts are also being made to obtain direct images of extrasolar planetary systems. One such example is Darwin, a proposed ESA space based mission that would utilize nulling interferometry in the infrared to directly search for terrestrial extrasolar planets. The compositional variations outlined here are likely to influence our ability to successfully detect these planets. Carbon-rich asteroids are known to be highly non-reflective. For example, 624 Hektor (D-type asteroid) has a geometric albedo of 0.025 while 10 Hygiea (C-type asteroid) has a geometric albedo of 0.0717. As both of these asteroids are assumed to be carbon-rich, it is likely that the carbon-rich planets identified here are similarly dark. Thus it is expected that searches for these planets in the visible spectrum will be difficult due to the small amount of light reflected by these bodies. However, a lower albedo results in greater thermal emission 167 from a body, suggesting that the infrared signature from these planets is much larger than corresponding silicate planets and extends to shorter wavelengths, assuming that any planetary atmosphere present reflects the composition of the solid body. As a result, infrared searches (such as that of Darwin) are ideally suited to detect carbon-rich terrestrial planets and thus should be focused on stellar systems with compositions similar to that of the C-rich stars identified here to maximize results.

4.5.5 Hydrous Species

None of the simulated terrestrial planets directly accrete any hydrous species (water or serpentine) for disk conditions at 5×105 years. This is understandable as all of the planets are located relatively close to their host star and well interior to the hydrous species region of the disk. At later times, all planets formed for 55Cnc and only the outer most planets of Gl777, HD72659 and HD4208 contain hydrous species. The planets of 55Cnc are understandably more enriched in water as they are located further out in the system, thus producing a greater overlap between their feeding zones and the water-rich region of the disk. However, the majority of planets accrete dry and contain no primordial water for the disk conditions simulated here. This will obviously inﬂuence not only planetary and atmospheric processes but will also impact on the potential for life to develop on these planets. Cometary and asteroidal delivery of water has been widely suggested as the origin of Earth’s water. However, it is questionable how effective such processes would be in extrasolar planetary systems. Only 55Cnc has been found to be dynamically capable of hosting an asteroid belt analogous to the belt in our own Solar System. Gl777 and HD72659 both contain giant planets in orbits that would render any similar feature within these planetary systems unstable, thus making its long term survival unlikely. As described in Chapter 3, water can also be delivered to a planetary body via adsorption onto solid grains within the disk. As this process has not been considered in our current simulations, it is likely that there will be some water delivered during the forma- 168 tion process to the terrestrial planets produced in the Earth-like systems (Gl777, HD4208, HD72659 and HD177830) as the solid grains are bathed in water vapor over the entire span of the disk. This same process will likely not be as effective at delivering water to the C-rich systems (55Cnc, HD4203, HD19994, HD108874 and HD142415) as they only have water vapor present at temperatures below ∼800 K. The different distributions can be seen when comparing the equilibrium gas compositions for HD 72659 (C/O = 0.32, Figure 4.32) and HD 4203 (C/O = 1.51, Figure 4.33). This temperature range corresponds to beyond a radial distance of ∼1.2AU for Hersant et al. (2001) midplane conditions at 2.5×105 years and ∼0.2AU for Hersant et al. (2001) midplane conditions at 3×106 years. As few terrestrial planets accrete material from beyond 1.2AU, it is expected that C-rich planets forming early in the lifetime of the disk will remain dry without additional water being delivered to the planets via adsorption. Delivery of water by exogenous sources (such as comets and asteroids) would be required to produce an ocean-bearing planet within the C-rich systems. However, these systems are predicted to contain less water ice and serpentine than their Solar-like companions. This is due to the relative O depletions within the selected systems. [O/H] values for the C-rich systems range from −0.21 to 0.13, depleted when compared both to other elements within the same system and in comparison to the Solar-like systems which vary from [O/H] = −0.14 to 0.38. These reduced O abundances result in the production of smaller amounts of the water-bearing species and would thus make it increasingly difﬁ- cult to provide water to a terrestrial planet within these systems. This variation in water abundance can be seen most easily when comparing Figures 4.12 and 4.20. Thus it appears that terrestrial planets within Solar-like systems are likely to obtain some amount of water (through temporal variations in composition, adsorption and exogenous delivery) while those in C-rich systems are likely to remain dry. 169

T vs Al(g) T vs Al2O(g) T vs AlH(g) HD72659 T vs Ca(g) T vs CH4(g) T vs CN(g) T vs CO(g) 10000 T vs CO2(g) T vs Cr(g) T vs CS(g) T vs Fe(g) T vs H(g) T vs H2(g) T vs H2O(g) 1000 T vs H2S(g) T vs HCN(g) T vs He(g) T vs HS(g)

ce (mole)ce T vs Mg(g) T vs N2(g) T vs Na(g) 100 T vs NaOH(g) bundan T vs NH3(g) A T vs Ni(g) ed T vs O(g) T vs P(g) T vs PH(g)

Normaliz T vs PN(g) 10 T vs PO(g) T vs PS(g) T vs S(g) T vs S2(g) T vs Si(g) T vs SiH(g) T vs SiO(g) 1 T vs SiS(g) T vs SO(g) 200 400 600 800 1000 1200 1400 1600 1800 T vs SO2(g) T (K) T vs Ti(g) T vs TiO(g) T vs TiO2(g)

Figure 4.32: Schematic of the output obtained from HSC Chemistry for HD72659 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 170

Al(g) Al2O(g) HD4203 AlH(g) Ca(g) CH4(g) CN(g) 10000 CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) H2O(g) 1000 H2S(g) HCN(g) He(g) HS(g) Mg(g) ce (mole)ce N2(g) Na(g) 100 NaOH(g)

bundan NH3(g)

A Ni(g)

ed O(g) P(g) PH(g) PN(g) PO(g) Normaliz 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 2000 2200 Ti(g) TiO(g) T (K) TiO2(g)

Figure 4.33: Schematic of the output obtained from HSC Chemistry for HD4203 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 171

4.5.6 Planetary Interiors and Processes

Given the wide variety of predicted planetary compositions, a similarly diverse range of planetary interiors is also expected. To better quantify this, I examined four specific cases: a 1.1ML Earth-like planet (Gl777), a 0.61ML Mg-rich Earth-like planet (HD177830), a 1.1ML refractory planet (HD4208) and a 0.47ML C-rich planet (HD108874). Ap- proximate interior structures for each were calculated using equilibrium mineralogy for a global magma ocean with P = 27GPa and T = 2000◦C. Equilibrium compositions at these conditions have been found to produce the best agreement between predicted and observed siderophile abundances within the primitive upper mantle of the Earth (Drake, 2000). Elemental abundances were taken from the results discussed in Section 4.4.2. Re- sulting mineral assemblages were sorted by density to define the compositional layers. Approximate planetary radii were obtained from Sotin et al. (2007) based on planetary mass. These planetary radii are based on silicate planetary equations of state and as such are unlikely to completely describe the C-rich planets observed here. However, no studies have considered such assemblages, forcing us to assume a silicate based approximate radius. Density variations at high pressures were not considered in defining the depths of various layers. Large impacts (such as the moon forming impact) are also neglected. The resulting interior structures (shown to scale) can be seen in Figure 4.34. The interior mineralogy and structure of the planet orbiting Gl777 is similar to Earth. It contains a pyroxene dominated crust (∼450 km deep) overlying an olivine mantle (∼1500 km deep) with an F-Ni-S core (radius ∼ 4600 km). The crust is significantly thicker than seen on Earth as I am currently neglecting density and phase changes. Given its structure and comparable mineralogy, I would expect to observe planetary processes similar to those seen on Earth. Melting conditions and magma compositions are expected to be comparable and it is feasible that a liquid core would develop, resulting in the production of a magnetic dynamo. Although I cannot make any specific statements regarding the likelihood of plate tectonics without further detailed modelling, the thicker crust of 172 this planet will increase the magnitude of the shear stresses required for lithospheric de- formation, one of the key requirements of tectonic activity (Valencia et al., 2007b). Even if the crustal thickness does prohibit global tectonics, stagnant lid convection may occur as is believed to have been the case on Mars. In general, based on their mass and composition, the terrestrial planets of Gl777 are likely to have structures and mineral assemblages similar to those observed in our system. The simulated planet for HD177830 is depleted in Si, relative to the Earth, resulting in high spinel and olivine content in the mantle (resembling that of type I kimberlites) and a thinner pyroxene and feldspathic crust than observed for Gl777 (∼190 km deep). Given its small mass and size, unless significant amounts of radioactive material are accreted or tidal heating is significant, it will be difficult to attain temperatures high enough to produce large amounts of magmatic melt within the mantle. If produced, however, they would have compositions similar to komatiite (dominated by olivine with trace amounts of pyroxene and plagioclase). Volcanic eruptions would be comparable to basaltic flows observed on Earth due to the low silica content of the melt. Similarly, driving mantle convection also requires significant levels of melt production. Although more comparable in thickness to the Earth’s crust, it is still questionable whether sufficient melt could be produced to induce plate tectonics on this planet. The terrestrial planet orbiting HD4208 is highly refractory in composition. The crust is thick (∼800 km) and has a composition resembling continental crust (felsic upper, ul- tramafic lower). Production of large amounts of melt within the mantle would be limited by the highly refractory composition, even considering that the large core would produce a considerable amount of heat via potential energy release. Furthermore, given the thickness of the crust, extrusive volcanism and plate tectonics are unlikely to occur as high stress levels would be required to fracture the crust. Producing such stresses without a vigorously convecting mantle would be challenging. Therefore, it is questionable whether or not a planet with this composition and structure would be active. Given the similar composition and size of the core, a magnetic dynamo is still expected to be pro- 173 duced within the core. Finally, carbide planets are expected to form around HD108874. The resulting composition and structure is unlike any known planet. Its small size, refractory composition and possible lack of radioactive elements (due to the potential absence of phosphate species, common hosts for U and Th, and possible lack of carbonates, the common host of K) will inhibit long-term geologic activity due to the difficulty in melting the mantle. Only large amounts of heat due to core formation and/or tidal heating would be able to provide the required mantle heating. Once all the primordial heat has been removed, it is unlikely that the mantle would remain molten on geologic timescales. Until that time, given the buoyancy of molten carbon, volcanic eruptions would be expected to be highly enriched in C. The core is also expected to be molten, thus making it likely that a magnetic dynamo would be produced (Gaidos and Selsis, 2007). In essence, although initially molten and probably active, old carbide planets of this type would be geologically dead. Recent studies have shown that coreless planets may be produced for compositions with significant amounts of oxidized material (Elkins-Tanton and Seager, 2008). As previously discussed, none of the terrestrial planets directly accrete hydrous material for midplane conditions at t = 5×105 years, nor are they highly oxidized. However, several planets did accrete water and serpentine for later disk conditions. Additionally, adsorption is also believed to contribute significant amounts of water to terrestrial planets in systems with Solar-like compositions. Therefore, oxidized planetary compositions and thus coreless planets may be produced for the Solar-like systems. To examine the effects of such a water-rich compositions on the predicted structure, I determined the composition for the same planet orbiting Gl777 but for disk conditions at t = 3×106 years. At this time, it contains a sizeable amount of oxidized, hydrous material. The predicted interior structure and composition is shown in Figure 4.35. It is clear that for conditions at t = 5×105 years the composition is considerably more oxidized. It is now dominated by olivine and troilite with minor amounts of feldspar, diopside and water. No pyroxene is predicted. The crust is slightly thinner and now 174 le st ant & 74 ore cru arth) arth) AU r) m per) te Si c R(E M(E (up Ni- phi owe 0.38 0.38 HD1088 Fe- SiC C (l Gra 0.79 0.79 ! 0.47 0.47 Fe e

e livin eldspar eldspar rust arth) arth) tle AU cor ) c & F & O R(E Ni Ni M(E ide & Spinel ne ne man wer 0.30 0.30 HD4208 per) Fe- (lo oxe 0.92 0.92 0.77 0.77 (up Diops Pyr le t ant r) & & e ) m arth) arth) ne ne r crus cor AU uppe wer Ni Ni R(E oxe M(E dspa 0.48 0.48 (lo ine ( Fe- Pyr HD177830 Fel 0.86 0.86 0.61 0.61 Oliv Spinel le st re ant cru arth) arth) AU S co ne ne R(E Ni- ine m M(E Gl777 oxe 0.89 0.89 Fe- Oliv 1.1 1.1 1.03 1.03 Pyr Figure 4.34: Schematic of simulated interiorssizes. of four extrasolar terrestrial planets. Figures are to scale for planet and layer 175 contains water. The core is also more oxidized, containing a large amount of magnetite

(Fe3O4). Based on these simulations, I cannot definitively claim that core formation within these planets will occur. Based on equilibrium assumptions, I would expect to see some degree of core formation, albeit with an altered composition. However, based on the simulations of Elkins-Tanton and Seager (2008), if the Fe fragments mixed throughout the magma ocean are less than 1cm in size, oxidation of the Fe would occur before it could sink and I would not expect to see any significant core development. If that were to occur with the current simulated abundances, we would see a deep mantle composed of primarily of olivine and iron oxides. Such a composition would prohibit the formation of a magnetic dynamo and alter the rheology of the planet. Present models of planetary formation are unable to simulate accretionary collisions in such detail. The different rheology of a carbide magma ocean (as would be the case for the C-rich systems) may also alter the size requirements for Fe fragments needed to form a core. Similarly, incomplete mixing of material accreted at later times is likely to result in deviations from the equilibrium picture presented here. For example, accretion of oxidized and water-rich material late in the formation process may result in a stratified redox state and water-rich crust as observed for the Earth. Unfortunately, it is not possible to determine these effects with current models as it requires a level of understanding of the impact and accretion process (e.g. mantle mixing, fragmentation) on small planetary bodies that we currently do not have. These results are also key for super Earth studies such as that of Valencia et al. (2007a) and O’Neill et al. (2007). Previous simulations have assumed Earth-based compositions and structures. Based on the present simulations, a wide variety of both are possible and will need to be considered.

4.5.7 Planet Habitability

The habitable zone of a planetary system is deﬁned as being the range of orbital radii for which water may be present on the surface of a planet. For the stars considered here, that 176

Feldspar crust Pyroxene crust Water ocean

Olivine mantle Olivine mantle

Fe-Ni-S core Fe-S and Fe oxide core

Gl777 Gl777 t = 0.5Myr t = 3Myr

Figure 4.35: Schematic of simulated interiors of Gl777 under two different disk conditions. Left: Planetary interior for disk conditions at t = 5×105 years. Right: Planetary interior for disk conditions at t = 3×106 years. Figures are to scale for planet and layer sizes. 177 corresponds to radii from ∼0.7AU to ∼1.45AU. The vast majority of the planets produced by the current simulations orbit interior to this region (exterior in the case of 55Cnc) and thus are unlikely to be habitable. Only 6 planets are produced within the classical habitable zone, existing in orbits extending from 0.84AU to 1.19AU. All six of these planets are formed in Solar-like systems (Gl777, HD4208 and HD72659) and have compositions comparable to that of Earth. Neglecting possible variations in atmospheric compositions, I feel that these systems (and others similar to them) are the ideal place to focus future astrobiological searches as they not only contain planets with compositions similar to that of Earth and are likely to obtain some amount of water during their evolution but also exist in the biologically favorable region of the planetary system. Of the ﬁve C-rich systems, only one produced a planet close to the habitable zone. HD19994 formed a terrestrial planet at 0.70AU, just at the inner edge of the habitable zone. All other planets are located well outside the required radii. As such, under the current deﬁnition of habitable, I conclude that it is unlikely that any of the C-rich planets formed in the current simulations would be capable of supporting life.

4.5.8 Biologically Important Elements

In addition to water, complex life (as we know it) also requires several key elements to exist. The six essential elements are H, C, N, O, P and S. As was the case for the Solar System simulations discussed in Chapter 3, none of the planets accreted any N. Given that they also failed to accrete any water, the final planets are also laking in H. The terrestrial planets formed in the Solar-like systems all contained various amounts of O, P and S but, as for the Solar System simulations, were deficient in C. The most C-rich systems, on the other hand, were lacking in O, P and S. Thus it is clear that for life to develop on any of the terrestrial planets formed in the current simulations, significant amounts of several elements must be supplied from exogenous sources within the system. All elements may be supplied from the outer, cooler regions of the disk. Thus it is possible that migration, comets, meteorites and/or radial 178 mixing may produce planets with the necessary elements for life to develop. As for the Solar System simulations, all biologically required elements would be introduced in a form that could be utilized by early life. This is especially intriguing for those planets located within the habitable zone. On the other hand, alternative pathways could be de- veloped for the formation of an alternative biologic cycle without requiring the same six elements.

4.5.9 Mass Distribution

Radial midplane mass distributions based on the equilibrium condensation sequence were calculated for each system. As composition is correlated to a speciﬁc radial distance within the midplane (via the Hersant et al. (2001) model), the total mass of solid material present within a given radial distance within the disk is given by:

2 Mass of solid material = Σi2πri Msolid, i (4.6)

where Msolid, i is the mass of solid material determined by the chemical model to be located at ri. Based on this calculation, the most carbon-rich systems simulated have unexpected differences in their mass distributions. The combination of a broad zone of refractory carbon-bearing solids and the relatively small amount of water ice that condenses in these systems suggests that C-rich systems have more solid mass located in the inner regions of the disk than in the outer regions. This mass distribution is the opposite of that expected in Solar-like systems. This can be seen in Figure 4.36 which shows the distribution of solid mass within each system under the assumption of equilibrium and for disk conditions at t = 5×105 years. Radial mixing of material is neglected. It can be seen from Figure 4.36 (top panel) that the planetary systems with Solar-like compositions have mass distributions essentially identical to that of the Solar System. 179

Only HD72659 displays any appreciable deviation, with marginally lower percentages than the Solar System (<2% difference) over much of the disk. The increase in mass observed for all systems at ∼0.6AU is due to the appearance of silicate species within the disk while the increase at ∼4.8AU is due to the condensation of hydrous species. The variation in slope between the different systems is due to different stellar elemental ratios (especially Mg/Si) resulting in slightly different ratios of the dominate silicate species present within the disk. These compositional variations produce marginally different mass distributions. Significant differences appear when considering the C-rich systems as seen in Fig- ure 4.36 (bottom panel). It can be seen that all of the C-rich systems contain a greater proportion of their mass within the inner disk than the Solar System does. Additionally, due to the reduced amount of water ice present in these systems, the marked increase in mass at ∼4.8AU observed in the Solar-like systems is not seen here. This increased mass concentration suggests that protoplanets may form more easily in the inner regions of the C-rich systems. As I cannot determine the exact amount of solid material present within these disks, I am unable to say whether it would be possible to produce a Jovian-planetary core in this region. However, if sufficient mass were present, it is conceivable that a giant planet core composed of refractory C-rich species may be produced within several AU of the host star, allowing for giant planet formation to occur much closer to the host star than previously thought. Such a scenario would obviously alter the extent and nature of planetary migration required within these systems as we would no longer need to migrate a planet from 5AU into 1-2AU. Alternatively, if insufficient mass is available for Jovian core formation, production of terrestrial planets in this region may proceed faster and with greater ease, thus increasing the chance for planetary detection. The full implications of these scenarios need to be examined by using alternative mass distributions for extrasolar planetary formation simulations for both gas giant planets and smaller terrestrial planets. 180

100

80 Solar Gl777 60 HD4208 HD177830

40 HD72659 % of Solid Material

0 0 1 2 3 4 5 Semimajor Axis (AU)

100

80 Solar 55Cnc 60 HD19994 HD4203

40 HD108874 HD142415 % of Solid Material

0 0 1 2 3 4 5 Semimajor Axis (AU)

Figure 4.36: Solid mass distribution within the disk for known extrasolar planetary systems. The mass distribution for the Solar System is also shown for comparison. Top Panel: Distribution determined for Solar-like systems studied (Gl777, HD4208, HD72659, HD177830). Bottom Panel: Distribution determined for the C-rich systems studied (55Cnc, HD4203, HD19994, HD108874, HD142415). 181

4.5.10 Stellar Pollution

As previously discussed, stellar photospheric pollution has been suggested as a possible explanation for the observed high metallicity of extrasolar planetary host stars (Laughlin, 2000; Gonzalez et al., 2001; Murray et al., 2001). The current simulations, though, do not support this hypothesis. Enrichments are produced primarily in Al, Ca and Ti, not Fe as is required by the pollution theory. Furthermore, relatively small masses of solid material are accreted by the host stars during planet formation, suggesting that insufﬁcient material is accreted to produce the observed enrichments. Thus unless migration of the giant planets can systematically result in accretion of giant planets by the host star, I have to agree with previous authors (e.g Santos et al. 2001, 2003, 2005; Fischer and Valenti 2005) and conclude that the observed host star enrichment is primordial in origin. The current simulations also imply that enrichments due to stellar pollution are most likely to be observed for the refractory elements in high mass stars with low convective zone masses. This suggests that surveys for pollution effects caused by terrestrial planet formation should focus on Ti, Al and Ca abundances in A-type and high mass F-type stars as they are expected to have the lowest convective zone masses. However, more detailed simulations of the fate of material accreted into radiative zones need to be undertaken to support this hypothesis.

4.6 Summary

Terrestrial planet formation simulations have been undertaken for nine different extrasolar planetary systems. Terrestrial planets were found to be ubiquitous, forming in all cases examined. Almost half of the simulations produced multiple terrestrial planets. The simulated planets are expected to be detectable by Kepler, thus allowing for future independent veriﬁcation of formation simulations. The compositions of these planets are found to vary greatly, from those comparable to Earth and CAI’s to other planets highly enriched in carbide phases. These compositional 182 variations are produced by deviations in the abundances of the host star and thus the system as a whole. Based on this, it is expected that C-rich planets will comprise a sizeable portion of extrasolar terrestrial planets and need to be considered in signiﬁcantly more detail. These compositions are highly dependant on the disk conditions selected for study, requiring us to develop a more detailed understanding of the timing of planetary formation within these systems. Given the wide variety of compositions predicted, it is also likely that planetary mineralogies and processes within these planets will be altered from those of our own Solar System. Compositions range from planets dominated by Fe and Mg-silicate species to those composed almost entirely of Fe and C. These compositional variations are likely to generate differences in delectability with C-rich planets being easier to detect via infrared surveys due to their lower albedo. The most habitable planets are expected to be those forming in systems with compositions similar to Solar. Water delivery, composition and orbital location make these planets ideal site for future biological surveys and studies. C-rich planets are likely to be lacking water and located interior to the habitable zone, making such planets unfavorable for the development of life as we know it. Finally, pollution of the host star by the planetary formation process appears to be negligible for the majority of systems. Enrichments are produced only for those stars with the least massive convective zones and even then only in the most refractory elements (Ti, Al and Ca). Therefore, it is unlikely that pollution is a viable explanation for the currently observed host star metallicity trend. This also implies that pollution studies should be undertaken for A-type and massive F-type stars as they are more likely to display the preferential enrichment in Ti, Al and Ca that appears to be indicative of terrestrial planet formation. 183

CHAPTER 5

SUMMARY & CONCLUSIONS

A number of the questions posed in the introductory chapter of this dissertation have begun to be answered. Extrasolar host star enrichments are primordial, not produced by pollution or external processes or additions. Simulated extrasolar terrestrial planets have been found to be common amongst systems and, based on their primordial enrichments, are likely to have a wide diversity of compositions. In Chapter 2, extrasolar planetary host stars were been found to be systematically enriched over non-host stars in several r- and s- process elements. These enrichments, however, are in keeping with general galactic chemical evolution trends. Thus although enriched in a variety of elements, host stars (and presumably the rest of the planetary system) are believed to have not undergone any unusual processing or alteration. Therefore, the abundances we are observing today are primordial in nature. Furthermore, given the apparent preference of planets to form around such stars, planet formation appears to be a natural part of the stellar material evolution. Given the primordial nature of the stellar photospheric enrichment, it is thus natural to consider the composition of terrestrial planets forming within these systems. In Chapter 3, I determined the bulk elemental abundances of the simulated terrestrial planets produced by the Solar System simulations of O’Brien et al. (2006). These abundances are in excellent agreement with observed planetary values, implying that the combination of dynamical and chemical modeling is successfully reproducing the terrestrial planets of the Solar System (to first order). These planets were also found to form wet, acquiring sufficient water during their formation to not need significant delivery of material from external sources, and with little dependance on the orbital properties of Jupiter and Saturn (for the main rock forming elements only). 185

In Chapter 4, I expanded on these successes to finally consider terrestrial planet formation in extrasolar planetary systems. Terrestrial planets were found to be ubiquitous, forming in all simulations completed for each of the nine planetary systems considered. The simulated terrestrial planets are generally found to be small (< 1ML) and are located close to their host star. The compositions of these planets, however, are truly diverse, ranging from Earth-like to refractory dominated and (most intriguingly) C-rich, dominated by carbide species. As these compositions are a reflection of the host star elemental abundances, stars with Solar elemental ratios are the best place to focus future Earth-like planet searches as these systems are found to produce the most Earth-like terrestrial planets often located within the habitable zones of their systems and containing a significant amount of water. Finally, C should be a major planet building element in ∼20% of known extrasolar planetary systems based on their host star photospheric compositions. Therefore it is logical that carbide planets like those simulated here are likely to exist in a significant number of planetary systems. These planets would be unlike anything we have previously observed and would produce an entirely different suite of planetary processes and structures for future consideration.

APPENDIX A

STELLAR PHOTOSPHERIC ABUNDANCES

The full list of all stellar photospheric abundances determined in Chapter 2 for the target stars of the Anglo-Australian Planet Search (AAPS) are provided in the following tables. Internally consistent abundances for Fe, Si, C, O, Cr, Mg, Ba, Y, Zr, Eu and Nd are provided. Stellar abundances of O, Cr, Mg, Ba, Y, Zr, Eu and Nd were determined in Chapter 2 of this dissertation and the reader is referred to the text for more details. Abundances of Fe, Si and C for the same stellar spectra were previously published in Bond et al. (2006) and are reproduced here for ease of comparison discussion within the text. Table A.1 lists the stellar elemental photospheric abundances of all target stars examined in Chapter 2. Abundances are given in the astronomical form of [X/H] (= log²X - J log²X, ). Solar values were taken from Asplund et al. (2005). Table A.2 lists the stellar elemental photospheric abundances of all target stars examined in Chapter 2 in the cosmochemical form of number of atoms present. Values are normalized to 106 Si atoms. Solar values were again taken from Asplund et al. (2005). 187 0.16 0.20 0.03 0.05 0.10 0.01 0.10 0.02 0.11 0.25 0.10 0.04 0.04 0.05 0.04 0.05 0.04 0.04 0.02 0.02 0.06 0.10 0.06 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Nd/H] 0.17 0.28 0.06 0.21 0.05 0.04 0.03 0.01 0.02 0.23 0.07 0.23 0.11 0.31 0.05 0.15 0.11 0.27 0.11 0.18 0.12 0.13 0.02 − − − − − − − − − 0.01 0.01 0.06 0.03 0.06 0.04 0.04 0.02 0.03 0.01 0.03 0.07 0.08 0.04 0.07 0.04 0.02 0.06 0.06 0.02 0.03 0..04 Continued on next page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Eu/H] 0.06 0.30 0.43 0.10 0.32 0.18 0.11 0.05 0.08 0.07 0.23 0.27 0.16 0.14 0.32 0.12 0.00 0.22 0.04 0.00 0.00 0.05 − − − − − − − − − − − − − − − − − 0.03 0.04 0.05 0.09 0.05 0.03 0.04 0.04 0.04 0.05 0.03 0.04 0.04 0.05 0.04 0.03 0.03 0.01 0.04 0.08 0.02 0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Zr/H] 0.18 0.02 0.35 0.05 0.04 0.04 0.05 0.07 0.08 0.10 0.06 0.15 0.24 0.23 0.29 0.23 0.01 0.21 0.20 0.28 0.08 0.19 − − − − − − − − − − 0.08 0.12 0.10 0.08 0.10 0.13 0.11 0.10 0.08 0.10 0.15 0.09 0.12 0.11 0.07 0.12 0.06 0.08 0.08 0.10 0.08 0.07 0.10 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Y/H] 0.44 0.13 0.35 0.32 0.19 0.08 0.03 0.11 0.29 0.17 0.06 0.14 0.00 0.03 0.11 0.02 0.06 0.00 0.18 0.01 0.00 0.16 0.09 − − − − − − − − − − − 0.03 0.02 0.03 0.08 0.08 0.03 0.03 0.03 0.02 0.10 0.02 0.10 0.03 0.11 0.03 0.04 0.09 0.05 0.03 0.03 0.07 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Ba/H] 0.14 0.44 0.07 0.48 0.18 0.22 0.11 0.00 0.05 0.07 0.16 0.11 0.04 0.09 0.31 0.11 0.11 0.01 0.15 0.23 0.05 0.09 indicates that a value could not be obtained from the spectrum. Solar values were − − − − − − − − − 0.07 0.07 0.02 0.03 0.07 0.09 0.03 0.06 0.08 0.09 0.10 0.01 0.01 0.09 0.06 0.09 0.01 0.01 0.02 0.08 0.01 0.03 0.01 0.01 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Host Stars [Mg/H] 0.15 0.02 0.07 0.10 0.09 0.07 0.00 0.16 0.07 0.22 0.14 0.04 0.19 0.22 0.15 0.27 0.13 0.33 0.07 0.06 0.29 0.20 0.07 0.14 − − − − − − − − − − − 0.07 0.06 0.05 0.08 0.05 0.05 0.05 0.04 0.06 0.05 0.05 0.06 0.06 0.03 0.04 0.06 0.02 0.06 0.03 0.05 0.05 0.05 0.04 0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Cr/H] 0.25 0.08 0.12 0.30 0.15 0.21 0.12 0.17 0.29 0.22 0.05 0.30 0.20 0.07 0.14 0.20 0.27 0.14 0.18 0.16 0.07 0.21 0.20 0.01 − − − − − − − − 0.04 0.02 0.02 0.03 0.01 0.04 0.02 0.06 0.02 0.07 0.02 0.07 0.04 0.02 0.02 0.04 0.01 0.03 0.02 0.01 0.05 0.02 0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [O/H] 0.09 0.13 0.14 0.22 0.00 0.46 0.24 0.30 0.08 0.08 0.01 0.27 0.13 0.07 0.09 0.19 0.21 0.15 0.15 0.02 0.03 0.05 0.17 − − − − − − − − − − − 0.12 0.07 0.08 0.03 0.11 0.08 0.05 0.02 0.13 0.06 0.05 0.13 0.07 0.08 0.11 0.03 0.06 0.06 0.09 0.10 0.05 0.05 0.11 0.11 1 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Si/H] 0.2 0.28 0.09 0.43 0.07 0.19 0.26 0.23 0.50 0.29 0.02 0.32 0.31 0.11 0.10 0.05 0.26 0.15 0.22 0.32 0.08 0.03 0.02 0.28 − − − − − − − − 0.16 0.11 0.12 0.09 0.06 0.09 0.18 0.23 0.13 0.15 0.10 0.04 0.12 0.09 0.02 0.04 0.11 0.08 0.06 0.10 0.10 0.06 0.07 1 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [C/H] 0.22 0.09 0.01 0.27 0.03 0.22 0.12 0.30 0.33 0.01 0.29 0.01 0.31 0.14 0.15 0.08 0.20 0.16 0.14 0.17 0.07 0.05 0.28 − − − − − − Table A.1: Stellar abundances obtained for all AAPS target stars with viable template spctra. A taken from Asplund et al. (2005). 2039 4308 HD 10647 13445 17051 20782 23079 30177 39091 70642 73526 75289 83443 102117 108147 117618 134987 142415 154857 160691 164427 179949 187085 196050 188 0.04 0.20 0.20 0.27 0.08 0.06 0.14 0.02 0.03 0.02 0.06 0.09 0.18 0.05 0.15 0.08 0.05 0.06 0.10 0.08 0.05 0.01 0.08 0.03 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Nd/H] 0.01 0.19 0.21 0.25 0.19 0.15 0.20 0.04 0.02 0.03 0.09 0.31 0.12 0.23 0.09 0.22 0.07 0.03 0.08 0.16 0.03 0.00 0.25 0.09 − − − − − − − − − − − − − 0.07 0.03 0.09 0.04 0.04 0.11 0.07 0.05 0.04 0.10 0.06 0.07 0.11 0.11 0.07 0.01 0.04 0.05 Continued on next page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − − − − − − [Eu/H] 0.14 0.36 0.38 0.07 0.17 0.12 0.28 0.28 0.22 0.30 0.26 0.21 0.12 0.20 0.14 0.00 0.05 0.12 − − − − − − − − − − − − − − 0.05 0.12 0.04 0.08 0.04 0.07 0.05 0.10 0.08 0.09 0.04 0.05 0.03 0.03 0.05 0.03 0.08 0.10 0.05 0.06 0.10 0.10 0.03 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − [Zr/H] 0.10 0.14 0.09 0.04 0.15 0.21 0.20 0.27 0.08 0.13 0.09 0.24 0.03 0.09 0.29 0.02 0.13 0.02 0.13 0.02 0.04 0.01 0.03 − − − − − − − − − − − − 0.11 0.12 0.09 0.10 0.11 0.09 0.09 0.11 0.09 0.08 0.10 0.08 0.09 0.08 0.05 0.03 0.10 0.09 0.07 0.08 0.08 0.08 0.14 0.10 0.08 0.08 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Y/H] 0.06 0.01 0.28 0.31 0.03 0.09 0.11 0.12 0.33 0.35 0.24 0.40 0.34 0.28 0.10 0.13 0.15 0.34 0.23 0.20 0.05 0.06 0.19 0.03 0.10 0.21 − − − − − − − − − − − − − − − − − − − − − − 0.05 0.02 0.09 0.07 0.10 0.05 0.13 0.07 0.10 0.08 0.12 0.09 0.04 0.04 0.12 0.13 0.05 0.12 0.08 0.05 0.03 0.12 0.02 0.07 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Ba/H] 0.06 0.07 0.26 0.07 0.14 0.09 0.12 0.10 0.06 0.05 0.41 0.21 0.04 0.07 0.40 0.50 0.39 0.02 0.19 0.09 0.32 0.20 0.13 0.04 − − − − − − − − − − − − − − − − − − − 0.06 0.05 0.11 0.08 0.12 0.13 0.08 0.02 0.04 0.04 0.03 0.11 0.11 0.04 0.08 0.06 0.08 0.08 0.04 0.09 0.09 0.02 0.12 0.09 0.06 0.07 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 0.1 [Mg/H] Non-Host Stars 0.18 0.09 0.08 0.06 0.19 0.09 0.10 0.04 0.28 0.33 0.01 0.06 0.01 0.17 0.24 0.13 0.09 0.26 0.21 0.26 0.33 0.19 0.17 0.17 0.05 − − − − − − − − − − − − − − − − − − − − 0.05 0.04 0.04 0.05 0.05 0.04 0.05 0.04 0.04 0.05 0.05 0.05 0.07 0.03 0.04 0.03 0.05 0.04 0.05 0.04 0.04 0.04 0.03 0.04 0.04 0.04 Table A.1 – continued from previous page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Cr/H] 0.02 0.09 0.24 0.14 0.07 0.15 0.02 0.11 0.07 0.04 0.20 0.20 0.18 0.20 0.07 0.18 0.31 0.24 0.24 0.21 0.27 0.26 0.21 0.20 0.07 0.03 − − − − − − − − − − − 0.01 0.02 0.02 0.02 0.02 0.01 0.01 0.05 0.03 0.01 0.03 0.03 0.03 0.01 0.01 0.01 0.01 0.11 0.07 0.09 0.02 0.01 0.01 0.01 0.08 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [O/H] 0.09 0.00 0.21 0.12 0.30 0.22 0.01 0.29 0.04 0.02 0.09 0.27 0.07 0.13 0.28 0.19 0.27 0.31 0.20 0.01 0.24 0.12 0.10 0.10 0.06 0.00 − − − − − − − − − − − − − − 0.06 0.07 0.08 0.10 0.12 0.04 0.07 0.04 0.05 0.16 0.14 0.06 0.09 0.08 0.09 0.10 0.07 0.15 0.05 0.05 0.09 0.07 0.08 0.05 0.03 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Si/H] 0.14 0.21 0.25 0.18 0.28 0.06 0.20 0.05 0.09 0.32 0.52 0.30 0.01 0.32 0.13 0.09 0.14 0.22 0.08 0.18 0.16 0.16 0.18 0.19 0.15 0.01 − − − − − − − − − − − 0.02 0.08 0.09 0.10 0.14 0.19 0.11 0.09 0.12 0.11 0.05 0.13 0.09 0.10 0.08 0.09 0.09 0.09 0.12 0.11 0.09 0.12 0.06 0.03 0.06 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [C/H] 0.02 0.18 0.17 0.23 0.20 0.38 0.06 0.18 0.37 0.03 0.03 0.30 0.04 0.28 0.22 0.05 0.26 0.03 0.04 0.19 0.08 0.19 0.13 0.07 0.01 − − − − − − − − − − − − 1581 3823 7570 9280 HD 10180 11112 12387 18709 19632 20201 20766 20794 20807 30295 31827 33811 36108 38283 38382 38973 42902 43834 208487 213240 216435 216437 189 0.10 0.07 0.25 0.01 0.05 0.04 0.04 0.07 0.01 0.18 0.09 0.02 0.18 0.01 0.01 0.01 0.04 0.01 0.05 0.08 0.04 0.10 0.06 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − [Nd/H] 0.04 0.09 0.35 0.04 0.37 0.05 0.08 0.10 0.06 0.11 0.11 0.07 0.29 0.02 0.00 0.11 0.09 0.03 0.31 0.11 0.08 0.06 0.07 0.05 − − − − − − − − − − − − 0.11 0.03 0.06 0.04 0.06 0.03 0.07 0.07 0.02 0.03 0.09 0.04 0.06 0.08 0.08 0.08 0.06 0.07 0.05 0.06 0.09 0.04 Continued on next page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − − − [Eu/H] 0.18 0.28 0.12 0.12 0.14 0.06 0.35 0.33 0.15 0.22 0.26 0.09 0.05 0.22 0.08 0.42 0.28 0.14 0.10 0.24 0.05 0.07 − − − − − − − − − − − − − − − − − − − − 0.05 0.05 0.08 0.06 0.06 0.09 0.04 0.04 0.08 0.08 0.15 0.08 0.05 0.06 0.05 0.03 0.09 0.05 0.07 0.05 0.06 0.11 0.05 0.08 0.04 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Zr/H] 0.05 0.13 0.32 0.08 0.02 0.06 0.01 0.06 0.20 0.03 0.09 0.35 0.01 0.05 0.04 0.03 0.21 0.02 0.11 0.06 0.08 0.04 0.12 0.11 0.13 0.00 − − − − − − − − − 0.12 0.10 0.08 0.11 0.05 0.07 0.04 0.07 0.09 0.08 0.12 0.09 0.10 0.09 0.10 0.10 0.10 0.10 0.08 0.07 0.12 0.10 0.11 0.08 0.09 0.11 0.22 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Y/H] 0.10 0.04 0.16 0.21 0.24 0.37 0.28 0.10 0.30 0.04 0.01 0.16 0.03 0.32 0.20 0.15 0.27 0.17 0.10 0.36 0.06 0.06 0.10 0.17 0.25 0.16 0.06 − − − − − − − − − − − − − − − − − − − − − − − − 0.09 0.12 0.09 0.04 0.06 0.02 0.04 0.04 0.05 0.05 0.05 0.08 0.03 0.05 0.07 0.10 0.11 0.09 0.13 0.11 0.04 0.04 0.10 0.05 0.03 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Ba/H] 0.05 0.05 0.06 0.34 0.07 0.26 0.41 0.41 0.12 0.23 0.15 0.14 0.10 0.01 0.19 0.19 0.21 0.10 0.25 0.11 0.12 0.35 0.30 0.05 0.08 0.07 − − − − − − − − − − − − − − − − − − 0.06 0.09 0.04 0.12 0.01 0.12 0.10 0.02 0.16 0.14 0.13 0.09 0.11 0.09 0.06 0.08 0.12 0.04 0.06 0.07 0.12 0.12 0.11 0.19 0.08 0.03 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± 0.2 0.1 0.3 [Mg/H] 0.01 0.02 0.08 0.08 0.20 0.07 0.01 0.05 0.01 0.02 0.02 0.23 0.23 0.05 0.02 0.10 0.24 0.10 0.13 0.06 0.02 0.26 0.13 0.04 − − − − − − − − − − − − − − − − − − 0.06 0.06 0.04 0.04 0.04 0.05 0.06 0.05 0.04 0.04 0.05 0.04 0.04 0.05 0.15 0.05 0.02 0.04 0.02 0.04 0.04 0.04 0.04 0.03 0.04 0.03 0.04 Table A.1 – continued from previous page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Cr/H] 0.09 0.10 0.10 0.22 0.08 0.05 0.23 0.08 0.08 0.09 0.07 0.10 0.09 0.02 0.01 0.03 0.20 0.18 0.20 0.24 0.03 0.23 0.05 0.03 0.01 0.01 0.26 − − − − − − − − − − − − − − 0.01 0.02 0.02 0.02 0.02 0.02 0.02 0.02 0.05 0.03 0.02 0.01 0.01 0.04 0.10 0.03 0.04 0.01 0.01 0.03 0.02 0.02 0.03 0.02 0.01 0.01 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [O/H] 0.02 0.02 0.00 0.10 0.12 0.16 0.26 0.17 0.10 0.26 0.10 0.05 0.07 0.04 0.34 0.04 0.25 0.11 0.13 0.06 0.12 0.05 0.21 0.05 0.02 0.07 0.00 − − − − − − − − − − − − − − − − − − − 0.07 0.09 0.05 0.07 0.13 0.09 0.13 0.10 0.04 0.09 0.21 0.06 0.14 0.06 0.07 0.07 0.09 0.09 0.08 0.05 0.09 0.08 0.09 0.10 0.06 0.07 0.12 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Si/H] 0.2 0.09 0.15 0.04 0.18 0.24 0.08 0.34 0.05 0.12 0.08 0.44 0.09 0.07 0.01 0.12 0.16 0.07 0.12 0.13 0.12 0.15 0.12 0.14 0.01 0.04 0.19 − − − − − − − − 0.10 0.05 0.13 0.13 0.18 0.08 0.08 0.06 0.05 0.19 0.04 0.06 0.06 0.10 0.08 0.05 0.06 0.12 0.09 0.11 0.08 0.14 0.09 0.15 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − [C/H] 0.2 0.1 0.13 0.08 0.12 0.24 0.30 0.07 0.11 0.11 0.02 0.33 0.02 0.01 0.00 0.03 0.13 0.00 0.14 0.08 0.09 0.13 0.12 0.18 − − − − − − − HD 44120 44594 45289 45701 52447 53705 53706 55720 59468 69655 72769 73121 73524 78429 80635 82082 86819 88742 92987 93385 96423 102365 105328 106453 107692 108309 83529A 190 0.13 0.14 0.01 0.04 0.01 0.03 0.13 0.06 0.07 0.03 0.04 0.11 0.01 0.03 0.01 0.10 0.08 0.04 0.01 0.11 0.04 0.07 0.01 0.04 0.06 0.08 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Nd/H] 0.04 0.14 0.07 0.15 0.16 0.04 0.16 0.03 0.14 0.02 0.18 0.03 0.05 0.00 0.04 0.14 0.03 0.07 0.06 0.10 0.11 0.06 0.13 0.05 0.25 0.11 − − − − − − − − − − − − 0.04 0.10 0.04 0.10 0.09 0.12 0.04 0.07 0.05 0.09 0.08 0.03 0.04 0.04 0.04 0.08 0.07 0.07 0.03 0.02 0.02 0.04 0.03 0.02 Continued on next page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − [Eu/H] 0.17 0.37 0.15 0.09 0.02 0.27 0.01 0.36 0.27 0.03 0.15 0.19 0.34 0.35 0.27 0.27 0.27 0.19 0.23 0.12 0.03 0.11 0.03 0.05 − − − − − − − − − − − − − − − − − − − − 0.05 0.03 0.05 0.05 0.10 0.05 0.04 0.05 0.07 0.02 0.09 0.06 0.04 0.06 0.04 0.03 0.04 0.04 0.06 0.04 0.05 0.04 0.09 0.04 0.04 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − [Zr/H] 0.25 0.03 0.10 0.02 0.67 0.09 0.05 0.04 0.12 0.08 0.31 0.09 0.33 0.11 0.14 0.03 0.17 0.02 0.02 0.00 0.09 0.13 0.16 0.04 0.14 0.04 − − − − − − − − − − − − − − − − 0.10 0.07 0.10 0.10 0.09 0.07 0.08 0.10 0.12 0.06 0.09 0.08 0.08 0.17 0.10 0.11 0.07 0.08 0.10 0.09 0.09 0.10 0.11 0.10 0.07 0.07 0.08 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Y/H] 0.04 0.30 0.22 0.09 0.04 0.13 0.50 0.15 0.21 0.07 0.13 0.24 0.37 0.03 0.14 0.35 0.15 0.01 0.05 0.10 0.51 0.31 0.11 0.15 0.20 0.12 0.02 − − − − − − − − − − − − − − − − − − − − − − − − − − 0.02 0.06 0.02 0.08 0.03 0.04 0.01 0.18 0.06 0.04 0.03 0.05 0.02 0.05 0.10 0.04 0.04 0.10 0.03 0.03 0.05 0.17 0.03 0.05 0.06 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − [Ba/H] 0.24 0.15 0.02 0.04 0.02 0.54 0.06 0.12 0.00 0.25 0.23 0.27 0.09 0.10 0.21 0.17 0.07 0.28 0.13 0.48 0.29 0.11 0.16 0.24 0.03 − − − − − − − − − − − − − − − − − 0.02 0.14 0.06 0.09 0.11 0.11 0.10 0.06 0.03 0.07 0.04 0.08 0.06 0.05 0.01 0.07 0.03 0.04 0.11 0.08 0.09 0.06 0.02 0.03 0.12 0.09 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Mg/H] 0.01 0.02 0.02 0.13 0.01 0.13 0.06 0.04 0.28 0.20 0.12 0.04 0.27 0.10 0.25 0.12 0.27 0.21 0.24 0.12 0.03 0.02 0.43 0.36 0.14 0.07 0.16 − − − − − − − − − − − − − − − − − − − − 0.04 0.05 0.05 0.04 0.05 0.03 0.05 0.03 0.05 0.04 0.03 0.03 0.04 0.04 0.05 0.04 0.06 0.09 0.03 0.04 0.04 0.04 0.05 0.04 0.06 0.04 0.05 Table A.1 – continued from previous page ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Cr/H] 0.09 0.10 0.04 0.04 0.09 0.02 0.11 0.06 0.03 0.20 0.10 0.08 0.00 0.05 0.22 0.15 0.30 0.26 0.16 0.24 0.10 0.29 0.02 0.54 0.35 0.01 0.01 − − − − − − − − − − − − − 0.01 0.01 0.02 0.02 0.02 0.04 0.03 0.04 0.02 0.02 0.01 0.03 0.03 0.02 0.01 0.02 0.01 0.03 0.04 0.01 0.02 0.01 0.02 0.05 0.02 0.01 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [O/H] 0.03 0.28 0.15 0.15 0.01 0.05 0.03 0.27 0.03 0.14 0.05 0.05 0.08 0.25 0.05 0.11 0.29 0.11 0.05 0.03 0.02 0.41 0.34 0.16 0.33 0.13 0.07 − − − − − − − − − − − − − − − − − − − − 0.06 0.07 0.07 0.05 0.09 0.08 0.07 0.06 0.11 0.08 0.14 0.14 0.08 0.07 0.06 0.01 0.09 0.08 0.09 0.07 0.05 0.09 0.09 0.05 0.10 0.10 0.09 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Si/H] 0.19 0.22 0.10 0.07 0.11 0.07 0.17 0.07 0.03 0.01 0.31 0.15 0.15 0.03 0.03 0.15 0.21 0.09 0.26 0.15 0.04 0.16 0.05 0.33 0.44 0.24 0.04 − − − − − − − − − − − 0.06 0.05 0.08 0.09 0.08 0.15 0.09 0.21 0.13 0.13 0.11 0.12 0.09 0.12 0.13 0.11 0.08 0.13 0.11 0.19 0.15 0.10 0.08 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± − − − − [C/H] 0.1 0.05 0.03 0.05 0.07 0.09 0.06 0.16 0.36 0.21 0.07 0.13 0.17 0.08 0.04 0.18 0.01 0.33 0.07 0.33 0.41 0.15 0.03 − − − − − − − − − − − − HD 114613 114853 122862 128620 134060 134330 140901 143114 147722 155974 161612 177565 183877 189567 192865 193193 193307 194640 196068 196800 199190 199288 199509 202628 204385 205536 207700 191 0.01 0.01 0.14 0.05 0.11 0.10 0.02 0.11 0.03 0.13 0.01 0.01 0.11 0.06 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Nd/H] 0.04 0.15 0.10 0.13 0.05 0.04 0.20 0.23 0.02 0.30 0.02 0.03 0.21 0.01 − − − 0.04 0.03 0.07 0.07 0.04 0.01 0.06 0.03 0.03 0.06 0.01 0.05 0.03 ± ± ± ± ± ± ± ± ± ± ± ± ± − [Eu/H] 0.23 0.41 0.17 0.28 0.06 0.31 0.16 0.17 0.01 0.10 0.01 0.06 0.05 − − − − − − − − − 0.03 0.04 0.04 0.04 0.09 0.04 0.08 0.01 0.06 0.09 0.03 0.01 0.04 0.05 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Zr/H] 0.04 0.24 0.05 0.01 0.28 0.15 0.06 0.17 0.03 0.10 0.20 0.01 0.17 0.05 − − − − − − − 0.10 0.09 0.09 0.11 0.09 0.10 0.07 0.12 0.14 0.08 0.09 0.10 0.10 0.07 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Y/H] 0.23 0.41 0.23 0.06 0.18 0.09 0.03 0.08 0.12 0.34 0.24 0.01 0.10 0.02 − − − − − − − − − − − − − 0.01 0.11 0.01 0.08 0.04 0.02 0.04 0.08 0.07 0.02 0.11 0.13 0.09 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Ba/H] 0.17 0.39 0.16 0.02 0.18 0.06 0.01 0.11 0.06 0.14 0.29 0.23 0.01 0.04 − − − − − − − − 0.07 0.08 0.05 0.15 0.05 0.12 0.07 0.09 0.06 0.03 0.09 0.08 0.10 0.04 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Mg/H] 0.19 0.07 0.00 0.13 0.08 0.14 0.15 0.11 0.10 0.01 0.11 0.18 0.01 0.08 − − − − − − − − − 0.04 0.04 0.04 0.05 0.06 0.04 0.04 0.04 0.03 0.04 0.03 0.03 0.04 0.04 Table A.1 – continued from previous page ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Cr/H] 0.15 0.00 0.02 0.18 0.18 0.14 0.20 0.04 0.12 0.30 0.10 0.03 0.13 0.03 − − − − − − 0.03 0.01 0.01 0.01 0.01 0.01 0.02 0.03 0.01 0.06 0.02 0.03 0.02 0.02 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [O/H] 0.15 0.16 0.12 0.11 0.13 0.06 0.16 0.06 0.03 0.13 0.18 0.04 0.11 0.03 − − − − − − − 0.08 0.07 0.06 0.08 0.11 0.07 0.07 0.08 0.05 0.06 0.07 0.11 0.05 0.06 ± ± ± ± ± ± ± ± ± ± ± ± ± ± [Si/H] 0.25 0.00 0.03 0.28 0.26 0.02 0.32 0.04 0.33 0.09 0.08 0.16 0.06 0.09 − − − − 0.09 0.14 0.06 0.09 0.13 0.13 0.06 0.08 0.12 ± ± ± ± ± ± ± ± ± − − − − − [C/H] 0.25 0.27 0.09 0.15 0.19 0.08 0.08 0.02 0.19 − − − − C and Si abundances were previously determined for the same spectra and published in Bond et al. (2006). 1 HD 209653 208998 210918 211317 212618 212330 212708 214759 214953 217958 219077 220507 221420 223171 192 − − − Eu 0.10 0.05 0.13 0.09 0.08 0.07 0.08 0.07 0.06 0.09 0.07 0.11 0.08 0.05 0.10 − indicates that a value could Nd 0.87 0.78 0.83 0.76 1.57 0.83 0.93 0.55 0.68 0.62 1.05 0.31 0.63 0.85 1.07 0.78 1.77 − − − Ba 4.57 3.31 2.34 8.91 3.16 5.62 3.31 1.32 2.95 2.82 3.16 3.02 5.25 5.37 4.47 9.33 Continued on next page Si atoms. A − − Zr 7.24 4.07 9.33 3.39 9.33 6 12.02 11.22 18.88 13.80 12.88 13.49 11.22 13.80 10.47 10.79 11.75 Y − 5.01 3.63 2.57 6.17 4.68 4.07 3.89 2.00 4.17 2.29 3.09 1.82 2.57 4.90 3.72 3.63 8.13 4 4 4 4 4 4 4 4 4 4 3 4 3 4 4 4 4 4 Cr 1.35 x10 1.26 x10 1.02 x10 1.62 x10 1.86 x10 1.32 x10 1.12 x10 1.00 x10 1.62 x10 1.41 x10 9.77 x10 1.17 x10 8.32 x10 1.15 x10 1.58 x10 1.10 x10 1.29 x10 1.32 x10 6 5 5 6 6 5 5 5 6 5 5 5 5 5 5 5 5 5 Mg 1.05 x10 7.76 x10 8.91 x10 1.23 x10 1.17 x10 6.31 x10 5.89 x10 4.57 x10 1.12 x10 8.32 x10 6.76 x10 6.17 x10 4.79 x10 6.31 x10 9.33 x10 8.71 x10 8.32 x10 8.51 x10 Host Stars 7 6 7 7 7 7 6 7 7 6 6 6 6 7 7 7 7 O − 1.41 x10 9.12 x10 1.48 x10 1.70 x10 1.78 x10 1.15 x10 9.77 x10 1.05 x10 1.10 x10 9.33 x10 8.51 x10 8.32 x10 9.77 x10 1.26 x10 1.10 x10 1.05 x10 1.38 x10 6 6 6 6 7 6 6 6 6 6 6 6 6 6 6 6 7 1 − C 7.59 x10 6.61 x10 7.41 x10 9.12 x10 1.07 x10 7.59 x10 7.76 x10 5.25 x10 6.92 x10 6.92 x10 5.89 x10 4.79 x10 8.32 x10 7.08 x10 7.41 x10 7.08 x10 1.23 x10 5 6 5 5 5 5 6 5 6 6 5 5 5 5 5 6 5 6 Fe 8.71 x10 1.00 x10 8.32 x10 9.77 x10 8.32 x10 7.94 x10 1.07 x10 6.76 x10 1.05 x10 1.00 x10 6.17 x10 6.76 x10 4.68 x10 6.76 x10 8.71 x10 0.76 x10 6.17 x10 1.45 x10 2039 4308 Solar HD 10647 13445 17051 23079 30177 39091 70642 73526 75289 83443 102117 108147 117618 134987 142415 not be obtained from the spectrum. Solar values are based on the abundances in Asplund et al. (2005). Table A.2: Stellar abundances for all AAPS target stars with viable template spectra normalized to 10 193 − − − Eu 0.10 0.03 0.06 0.07 0.04 0.04 0.09 0.10 0.06 0.06 0.06 0.07 0.06 0.06 0.05 0.06 0.08 Nd 1.62 0.65 0.89 0.93 0.74 0.50 1.05 1.07 0.66 0.81 0.78 0.89 0.32 0.30 0.89 0.59 0.74 0.83 0.93 1.02 Ba 6.76 2.29 2.51 2.14 4.57 3.09 4.90 3.89 5.13 2.19 4.57 6.17 2.29 1.91 3.47 2.57 2.14 4.27 3.72 4.37 Continued on next page Zr 9.33 7.41 7.08 5.13 9.55 7.94 6.76 19.05 18.20 12.88 10.23 15.85 10.72 14.45 13.18 14.45 10.72 10.23 11.22 13.18 Y 4.90 2.45 2.63 5.75 3.89 3.39 5.37 3.55 4.79 3.02 3.63 4.07 3.09 2.14 3.39 2.40 2.82 3.39 2.57 5.13 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 3 4 4 4 Cr 1.62 x10 1.05 x10 1.23 x10 1.48 x10 1.17 x10 1.17 x10 1.74 x10 1.20 x10 1.45 x10 1.05 x10 1.23 x10 1.10 x10 1.05 x10 1.00 x10 1.23 x10 1.10 x10 9.33 x10 1.17 x10 1.41 x10 1.20 x10 6 5 5 5 5 5 6 5 5 5 5 5 5 5 5 5 5 5 5 5 Mg 1.26 x10 7.08 x10 6.92 x10 9.12 x10 6.76 x10 8.91 x10 1.02 x10 6.92 x10 9.77 x10 7.24 x10 7.59 x10 8.13 x10 6.76 x10 6.61 x10 7.94 x10 6.46 x10 8.51 x10 9.12 x10 6.92 x10 6.92 x10 Non-Host Stars 7 6 7 7 7 7 7 7 7 6 7 6 7 7 6 7 7 7 6 O − 1.91 x10 9.77 x10 1.05 x10 1.20 x10 1.12 x10 1.41 x10 1.02 x10 1.41 x10 1.05 x10 9.77 x10 1.41 x10 9.55 x10 1.45 x10 1.12 x10 9.33 x10 2.09 x10 1.15 x10 1.07 x10 8.51 x10 Table A.2 – continued from previous page 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 C 7.59 x10 7.59 x10 8.13 x10 8.51 x10 8.13 x10 6.61 x10 9.77 x10 8.32 x10 6.92 x10 7.24 x10 5.50 x10 7.59 x10 7.94 x10 9.55 x10 7.59 x10 7.24 x10 8.13 x10 6.31 x10 6.31 x10 5.62 x10 5 5 5 5 5 5 5 6 5 5 5 5 6 6 5 6 5 5 5 5 Fe 8.91 x10 6.46 x10 6.61 x10 9.12 x10 7.59 x10 5.89 x10 9.12 x10 0.78 x10 6.61 x10 6.61 x10 7.24 x10 7.24 x10 0.68 x10 0.62 x10 7.41 x10 0.65 x10 5.89 x10 7.76 x10 6.31 x10 7.41 x10 1581 3823 7570 9280 HD 10180 11112 12387 18709 19632 20201 154857 160691 164427 179949 187085 196050 208487 213240 216435 216437 194 − − − − − − Eu 0.05 0.09 0.08 0.09 0.09 0.08 0.06 0.10 0.05 0.08 0.05 0.05 0.10 0.10 0.12 − − − Nd 1.15 0.58 1.20 1.23 0.45 1.10 0.60 0.89 0.65 0.74 0.79 0.65 1.20 0.83 0.58 1.48 1.00 1.41 − − − Ba 2.63 2.09 2.82 2.09 0.89 5.75 3.09 7.41 3.31 3.72 4.17 3.63 3.63 6.61 3.09 3.31 2.51 2.34 Continued on next page − − − Zr 9.33 5.75 6.03 9.77 9.55 9.77 5.13 7.59 14.45 13.49 15.14 13.49 12.02 10.47 10.47 17.78 17.38 17.78 Y 2.88 2.57 3.31 3.98 1.91 1.12 1.78 3.55 4.17 3.24 4.37 3.89 3.24 3.24 3.24 3.16 5.89 1.78 3.80 3.02 3.47 4 4 4 4 4 3 4 4 4 4 4 4 4 4 3 4 4 3 4 4 4 Cr 1.20 x10 1.91 x10 1.05 x10 1.12 x10 1.02 x10 6.46 x10 1.02 x10 1.29 x10 1.20 x10 1.17 x10 1.23 x10 1.02 x10 1.17 x10 1.05 x10 9.33 x10 1.15 x10 1.10 x10 9.77 x10 1.12 x10 1.26 x10 1.12 x10 5 6 5 5 5 5 5 5 5 5 5 5 5 5 5 5 5 5 5 5 5 Mg 8.32 x10 1.17 x10 9.33 x10 8.71 x10 5.75 x10 4.90 x10 6.46 x10 7.59 x10 9.55 x10 7.24 x10 6.92 x10 4.47 x10 6.17 x10 8.71 x10 7.08 x10 9.12 x10 7.24 x10 7.24 x10 8.71 x10 8.71 x10 8.13 x10 7 6 7 7 7 6 6 7 7 7 7 6 7 7 7 7 7 7 7 7 7 O 1.07 x10 5.25 x10 1.32 x10 1.15 x10 1.38 x10 6.76 x10 7.24 x10 1.26 x10 1.51 x10 1.15 x10 1.10 x10 7.76 x10 1.05 x10 1.10 x10 1.05 x10 1.29 x10 1.17 x10 1.07 x10 2.69 x10 1.10 x10 1.26 x10 Table A.2 – continued from previous page 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 C − − 7.08 x10 9.12 x10 7.41 x10 5.37 x10 4.07 x10 8.71 x10 9.12 x10 8.32 x10 7.24 x10 7.24 x10 6.17 x10 8.32 x10 6.46 x10 9.12 x10 8.71 x10 8.71 x10 8.32 x10 8.71 x10 7.41 x10 5 5 5 5 6 6 6 5 5 5 5 6 6 5 5 5 6 6 5 5 5 Fe 7.59 x10 7.94 x10 6.03 x10 7.41 x10 0.59 x10 0.42 x10 0.62 x10 7.76 x10 7.08 x10 7.24 x10 7.24 x10 0.62 x10 0.74 x10 7.24 x10 6.76 x10 6.92 x10 0.68 x10 0.60 x10 6.76 x10 6.92 x10 6.17 x10 HD 20766 20782 20794 20807 30295 31827 33811 36108 38283 38382 38973 42902 43834 44120 44594 45289 45701 52447 53705 53706 55720 195 − − − − Eu 0.06 0.02 0.06 0.09 0.05 0.02 0.07 0.09 0.07 0.08 0.03 0.08 0.05 0.06 0.04 0.04 0.07 − − Nd 0.54 0.83 1.35 0.93 0.29 0.30 1.23 0.74 0.76 0.65 0.66 0.79 1.05 0.69 0.87 0.74 0.55 0.63 1.29 − − Ba 2.88 3.72 5.75 4.79 3.02 1.70 5.75 4.68 2.88 3.98 2.19 3.47 2.63 3.16 1.58 3.47 2.88 2.88 4.27 Continued on next page − Zr 8.32 5.62 8.91 18.20 11.48 14.79 10.47 11.22 18.20 10.72 15.85 10.00 10.23 10.00 11.75 10.96 13.18 12.30 11.48 10.96 Y 3.31 3.47 2.09 4.37 5.50 2.88 1.70 4.68 3.80 3.24 3.89 2.29 3.31 3.02 3.39 3.02 3.72 3.02 2.51 2.95 4.07 4 4 4 4 4 4 3 4 4 4 4 4 4 4 4 4 4 4 4 4 4 Cr 1.41 x10 1.07 x10 1.02 x10 1.12 x10 1.23 x10 1.26 x10 8.32 x10 1.32 x10 1.26 x10 1.23 x10 1.38 x10 1.12 x10 1.29 x10 1.23 x10 1.15 x10 1.15 x10 1.35 x10 1.29 x10 1.23 x10 1.07 x10 1.32 x10 5 5 5 5 5 5 5 6 5 5 5 5 5 5 5 5 5 5 5 5 5 Mg 6.92 x10 7.24 x10 5.75 x10 8.32 x10 7.59 x10 6.92 x10 6.03 x10 1.00 x10 9.55 x10 8.51 x10 8.51 x10 7.76 x10 9.77 x10 8.13 x10 8.91 x10 8.13 x10 6.61 x10 7.24 x10 7.94 x10 6.17 x10 8.91 x10 6 7 6 7 6 7 7 7 7 7 7 7 7 6 7 6 7 6 7 6 7 O 9.33 x10 1.07 x10 8.13 x10 1.12 x10 9.12 x10 1.07 x10 1.12 x10 1.05 x10 1.26 x10 1.12 x10 1.15 x10 1.05 x10 1.05 x10 9.55 x10 1.35 x10 8.71 x10 1.15 x10 9.12 x10 1.05 x10 9.77 x10 1.20 x10 Table A.2 – continued from previous page 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 C − − − − 7.41 x10 7.94 x10 8.71 x10 7.41 x10 6.61 x10 5.89 x10 6.46 x10 7.59 x10 7.94 x10 8.32 x10 7.94 x10 7.76 x10 7.08 x10 5.13 x10 5.75 x10 7.76 x10 8.32 x10 5 5 6 5 6 5 6 5 5 5 5 5 5 5 5 5 5 5 5 6 5 Fe 8.13 x10 7.24 x10 0.55 x10 7.41 x10 0.81 x10 7.41 x10 0.48 x10 6.61 x10 8.71 x10 7.24 x10 7.59 x10 6.17 x10 7.08 x10 6.76 x10 7.08 x10 6.76 x10 7.24 x10 6.03 x10 5.89 x10 0.65 x10 9.33 x10 HD 59468 69655 72769 73121 73524 78429 80635 82082 86819 88742 92987 93385 96423 102365 105328 106453 107692 108309 114613 114853 83529A 196 − − Eu 0.09 0.05 0.08 0.09 0.10 0.10 0.09 0.06 0.04 0.08 0.10 0.06 0.05 0.10 0.05 0.05 0.08 0.04 0.15 − Nd 0.78 0.45 0.69 0.81 0.93 1.12 0.79 0.41 0.87 1.10 1.15 0.56 0.91 1.35 1.07 0.55 0.71 0.59 1.58 1.41 − Ba 3.98 3.80 4.27 3.72 2.40 3.39 8.51 3.09 2.19 2.95 3.55 5.25 4.07 6.03 3.02 1.91 6.17 2.40 4.17 4.07 Continued on next page − Zr 5.75 8.91 4.68 7.24 9.33 9.55 7.94 9.33 13.80 10.00 10.72 13.80 10.00 14.45 11.22 16.60 12.59 12.30 15.49 16.22 Y 3.72 2.45 3.63 3.16 4.07 2.88 3.02 4.37 2.88 3.16 3.16 3.09 4.37 4.07 4.79 3.47 2.40 3.16 2.82 4.27 4.27 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 4 Cr 1.17 x10 1.02 x10 1.17 x10 1.26 x10 1.29 x10 1.23 x10 1.20 x10 1.05 x10 1.17 x10 1.32 x10 1.02 x10 1.12 x10 1.35 x10 1.20 x10 1.48 x10 1.26 x10 1.05 x10 1.20 x10 1.15 x10 1.07 x10 1.05 x10 5 5 5 5 5 6 5 5 5 5 6 5 6 5 6 5 5 5 5 6 5 Mg 8.13 x10 6.46 x10 8.71 x10 6.76 x10 7.41 x10 1.02 x10 7.08 x10 8.32 x10 7.41 x10 6.76 x10 1.55 x10 8.13 x10 1.00 x10 7.24 x10 1.29 x10 7.76 x10 6.92 x10 6.92 x10 7.08 x10 1.07 x10 7.94 x10 7 7 7 7 7 7 7 7 7 7 7 7 7 7 7 7 6 7 6 7 7 O 1.23 x10 1.20 x10 1.10 x10 1.07 x10 1.17 x10 1.38 x10 1.12 x10 1.45 x10 1.07 x10 1.07 x10 1.29 x10 1.15 x10 1.17 x10 1.23 x10 1.55 x10 1.07 x10 7.76 x10 1.07 x10 9.55 x10 1.51 x10 1.12 x10 Table A.2 – continued from previous page 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 6 C − − − 7.76 x10 6.76 x10 5.89 x10 6.31 x10 9.12 x10 7.24 x10 8.51 x10 6.03 x10 7.94 x10 8.13 x10 5.13 x10 8.13 x10 7.24 x10 7.59 x10 5.37 x10 8.71 x10 8.13 x10 12.30 x10 5 5 5 5 5 5 5 5 6 5 5 5 5 5 5 5 6 5 5 5 5 Fe 7.59 x10 5.62 x10 7.41 x10 7.08 x10 7.59 x10 6.31 x10 6.76 x10 7.76 x10 0.68 x10 6.31 x10 6.31 x10 7.08 x10 7.59 x10 7.94 x10 7.59 x10 0.63 x10 6.92 x10 6.76 x10 7.76 x10 10.00 x10 26.30 x10 HD 122862 128620 134060 134330 140901 143114 147722 155974 161612 177565 183877 189567 192865 193193 193307 194640 196068 196800 199190 199288 199509 197 − − Eu 0.07 0.06 0.05 0.07 0.06 0.08 0.07 0.05 0.08 0.05 0.03 0.11 0.09 0.06 0.05 0.08 Nd 1.29 0.91 1.45 0.79 0.95 0.89 1.35 0.55 0.95 1.29 0.78 0.50 0.66 0.83 1.12 0.85 0.66 0.72 Ba 6.46 2.95 2.45 3.02 3.72 2.69 3.63 2.69 3.02 4.90 2.45 3.24 5.01 1.58 2.88 2.45 2.09 4.07 Zr 8.13 7.94 6.46 3.02 9.55 8.32 18.20 12.30 13.18 10.00 12.30 10.00 12.88 14.13 10.00 11.75 10.47 10.96 Y 4.27 3.31 2.95 2.69 3.55 2.82 3.39 2.45 3.31 3.80 2.45 2.88 3.98 1.82 2.82 2.63 2.29 3.24 4 4 4 4 4 3 4 4 4 4 4 4 4 3 4 4 4 4 Cr 1.48 x10 1.23 x10 1.23 x10 1.07 x10 1.23 x10 9.77 x10 1.23 x10 1.07 x10 1.35 x10 1.32 x10 1.07 x10 1.12 x10 1.20 x10 8.91 x10 1.23 x10 1.15 x10 1.00 x10 1.20 x10 5 5 5 5 5 6 5 5 5 5 5 5 5 5 5 5 5 5 Mg 8.32 x10 8.32 x10 6.76 x10 8.51 x10 9.12 x10 2.34 x10 8.51 x10 6.92 x10 8.13 x10 7.76 x10 5.50 x10 5.62 x10 7.76 x10 6.76 x10 8.51 x10 9.33 x10 5.89 x10 7.08 x10 7 6 6 7 7 7 7 7 7 7 7 6 7 6 7 7 6 7 O 1.07 x10 6.17 x10 9.77 x10 1.17 x10 1.20 x10 1.41 x10 1.23 x10 1.02 x10 1.05 x10 1.15 x10 1.07 x10 8.91 x10 1.26 x10 9.12 x10 1.15 x10 1.41 x10 8.51 x10 1.23 x10 Table A.2 – continued from previous page 6 6 6 6 6 6 6 6 6 6 6 6 6 C − − − − − 5.89 x10 8.32 x10 6.61 x10 7.24 x10 7.59 x10 9.12 x10 8.32 x10 7.59 x10 4.90 x10 7.41 x10 8.91 x10 9.77 x10 9.55 x10 5 5 5 5 5 5 5 6 5 5 6 6 5 6 5 5 6 5 Fe 7.41 x10 7.24 x10 6.76 x10 6.61 x10 7.41 x10 5.75 x10 7.59 x10 0.62 x10 7.76 x10 8.13 x10 0.62 x10 0.63 x10 7.24 x10 0.51 x10 6.61 x10 6.61 x10 0.60 x10 7.24 x10 HD 202628 204385 205536 207700 209653 208998 210918 211317 212618 212330 212708 214759 214953 217958 219077 220507 221420 223171 C abundances were previously determined for the same spectra and published in Bond et al. (2006). 1 198

APPENDIX B

SOLAR SYSTEM TERRESTRIAL PLANET ABUNDANCES

Bulk elemental abundances were determined in Chapter 3 for all of the terrestrial planets produced in the simulations of O’Brien et al. (2006). Planetary abundances were obtained for seven different sets of midplane conditions at various disk evolutionary times (t = 2.5×105yr, 5×105yr, 1×106yr, 1.5×106yr, 2×106yr, 2.5×106yr and 3×106yr). Appendix B contains the graphical and numerical results of each of these simulations. Figures B.1 - B.7 show the normalized abundances for each planet produced by the simulations of O’Brien et al. (2006). Normalized abundances are shown for each of the seven sets of disk conditions examined. The terrestrial planet to which each simulated planet was normalized was determined based on the semi-major axis of the terrestrial planet and is shown in parentheses in the upper left of each plot. Figures B.1 and B.4 were previously shown as Figure 3.3 in the text and are included here again for completeness. The complete assemblage of all predicted bulk elemental abundances for each of the simulated planets and for each of the seven different sets of disk conditions examined are provided in Table B.1. Values are provided as bulk wt% of the final planet for each set of disk conditions. Note that planetary numbers start at 4 and increase with increasing distance from the Sun. Table B.2 lists the ensemble-averaged bulk elemental abundances for the terrestrial planets of the O’Brien et al. (2006) simulations. Abundances are averaged over all seven sets of disk conditions simulated (from t = 2.5×105 to t = 3×106 years). Table B.3 lists the difference in bulk elemental abundances between the first and last set of disk conditions examined. As such, this table provides the range in each elemental abundance for the terrestrial planets of O’Brien et al. (2006). The difference is defined as

AbundanceDisk at 3×106 yr − AbundanceDisk at 2.5×105 yr. 199

A ﬁrst order approximation of planetary abundances incorporating volatile loss during the accretion process was also undertaken (see Section 3.3.5). Figures B.8 - B.14 show the normalized abundances for each planet produced by the simulations of O’Brien et al. (2006) both with and without volatile loss incorporated. All abundances are determined for disk conditions at 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. Reference Solar System planetary abundances were taken from Morgan and Anders (1980) (Venus), Kargel and Lewis (1993)(Earth) and Lodders and Fegley (1997) (Mars). Figures B.8 and B.11 were previously shown as Figure 3.9 and are included here again for completeness. Table B.4 lists the bulk elemental abundances predicted for the terrestrial planets of the O’Brien et al. (2006) simulations after volatile loss during impact events has been incorporated. Abundances are listed as wt% for all planets for each of the seven disk conditions examined. 200

10 CJS1-4 (Venus)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS1-5 (Earth)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS1-6 (Mars)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.1: Normalized planetary abundances for CJS1 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 201

10 CJS2-4 (Venus) CJS2-6 (Earth) 10

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS2-5 (Venus) CJS2-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.2: Normalized planetary abundances for CJS2 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 202

10 10 CJS4-4 (Venus)

e CJS3-4 (Venus) e anc anc

1 1 Normalized Abund Normalized Abund Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS4-5 (Mars)

e CJS3-5 (Earth) e anc anc

1 1 Normalized Abund Normalized Abund Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS3-6 (Mars) Increasing volatility e anc

1 Normalized Abund Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.3: Normalized planetary abundances for CJS3 and CJS4 simulated terrestrial planets. Values are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. Left: CJS3 simulation results. Right: CJS4 simulation results. 203

10 EJS1-4 (Venus)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 EJS1-5 (Earth)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 EJS1-6 (Mars)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.4: Normalized planetary abundances for EJS1 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 204

10 EJS2-4 (Venus) EJS2-6 (Earth) 10

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS2-5 (Venus) EJS2-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.5: Normalized planetary abundances for EJS2 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 205

10 10 EJS3-4 (Venus) EJS3-6 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS3-5 (Earth) EJS3-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.6: Normalized planetary abundances for EJS3 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 206

10 10 EJS4-4 (Venus) EJS4-6 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS4-5 (Venus) EJS4-7 (Mars)

1 0.1 Normalized Abundance Normalized Abundance Normalized 0.01 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.7: Normalized planetary abundances for EJS4 simulated terrestrial planets. Val- ues are shown for each of seven time steps considered with the following color scheme: black = 2.5×105 years, red = 5×105 years, green = 1×106 years, pink = 1.5×106 years, light blue = 2×106 years, yellow = 2.5×106 years and dark blue = 3×106 years. 207 0.89 0.89 0.90 0.58 0.89 0.88 0.91 0.71 0.90 0.90 0.78 0.90 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.61 16.66 16.34 14.27 16.77 16.52 16.09 15.04 16.59 15.78 15.83 16.26 Ti 0.09 0.09 0.07 0.66 0.10 0.11 0.06 0.45 0.08 0.06 0.31 0.06 wt% P 0.09 0.11 0.13 0.03 0.09 0.10 0.14 0.05 0.11 0.13 0.07 0.13 wt% Cr 0.38 0.40 0.41 0.12 0.41 0.37 0.40 0.20 0.41 0.39 0.28 0.40 wt% Ni 1.79 1.79 1.78 0.60 1.80 1.74 1.76 1.04 1.80 1.73 1.34 1.78 wt% years Na 0.31 0.35 0.58 0.09 0.33 0.35 0.60 0.17 0.42 0.58 0.21 0.57 5 wt% 10 × Ca 2.39 2.41 1.84 2.58 2.81 1.63 2.07 1.71 7.70 1.73 wt% 16.12 11.07 t=2.5 Continued on next page Al 1.65 1.66 1.27 1.77 1.94 1.12 7.70 1.43 1.18 5.35 1.19 wt% 11.24 Fe 9.23 wt% 27.64 28.12 28.08 28.41 27.09 27.83 15.66 28.34 27.22 20.76 28.07 S 2.16 1.61 3.34 0.58 0.78 2.18 4.98 1.56 2.22 5.98 1.13 4.07 wt% O wt% 32.05 31.94 31.43 38.76 32.06 32.19 30.84 36.45 31.66 30.99 34.61 31.03 Mg 8.29 wt% 14.84 14.86 14.73 14.89 14.60 14.57 10.61 14.88 14.24 12.41 14.69 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Table B.1: Predicted bulk planetary abundancesﬁnal for planet. the CJS denotes terrestrial the planets simulations of of O’Brien theEJS et denotes O’Brien al. the et (2006) results with al. of Jupiter the (2006) and simulations simulations. Saturnincreasing with in distance Jupiter All the from and circular values the Saturn orbits are Sun. in predicted their wt% by current the of elliptical Nice the orbits. model while Planetary numbers start at 4 and increase with Simulation 208 0.86 0.77 0.91 0.64 0.81 0.90 0.91 0.71 0.90 0.87 0.90 0.71 0.86 0.91 0.90 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.77 15.64 16.34 13.71 16.67 16.95 15.68 14.87 17.05 16.81 15.35 15.41 16.62 16.97 16.56 Ti 0.15 0.34 0.06 0.55 0.25 0.08 0.06 0.44 0.08 0.14 0.07 0.46 0.16 0.06 0.07 wt% P 0.07 0.08 0.13 0.07 0.05 0.12 0.13 0.04 0.10 0.12 0.13 0.03 0.04 0.14 0.13 wt% Cr 0.35 0.28 0.41 0.22 0.28 0.42 0.39 0.23 0.41 0.36 0.37 0.20 0.38 0.42 0.41 wt% Ni 1.73 1.25 1.79 1.02 1.42 1.83 1.72 1.22 1.83 1.61 1.64 1.07 1.72 1.86 1.79 wt% Na 0.11 0.29 0.57 0.24 0.18 0.41 0.59 0.10 0.21 0.50 0.54 0.08 0.09 0.58 0.50 wt% Ca 3.90 8.47 1.65 6.20 2.17 1.58 2.28 3.44 1.82 4.05 1.72 1.95 wt% 13.59 10.85 11.27 Continued on next page Al 2.69 5.90 1.13 9.42 4.32 1.49 1.09 7.51 1.56 2.41 1.26 7.84 2.80 1.18 1.34 wt% Table B.1 – continued from previous page Fe wt% 26.53 19.64 28.25 15.95 21.61 28.91 27.16 18.54 28.78 25.26 25.85 16.04 26.97 29.36 28.29 S 0.31 1.22 4.18 0.65 0.82 0.36 7.57 0.25 0.22 1.81 6.89 0.14 0.13 0.48 2.72 wt% O wt% 32.91 34.92 30.70 35.84 34.63 32.07 29.84 35.36 32.20 32.95 32.30 36.60 32.71 31.87 31.31 Mg 8.74 wt% 14.48 11.97 14.79 13.57 15.20 14.19 10.60 15.27 14.60 13.79 10.86 14.33 15.36 14.91 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 209 0.90 0.90 0.91 0.89 0.91 0.90 0.91 0.90 0.91 0.91 0.90 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.34 16.44 15.56 16.88 16.28 16.15 15.27 16.57 15.99 15.12 16.51 15.47 Ti 0.09 0.08 0.08 0.11 0.08 0.08 0.07 0.10 0.08 0.07 0.09 0.07 wt% P 0.12 0.14 0.13 0.09 0.14 0.13 0.13 0.11 0.13 0.13 0.12 0.13 wt% Cr 0.39 0.41 0.39 0.39 0.41 0.40 0.38 0.39 0.40 0.38 0.40 0.39 wt% Ni 1.76 1.80 1.71 1.79 1.79 1.77 1.67 1.78 1.75 1.66 1.79 1.70 wt% years Na 0.42 0.60 0.58 0.20 0.59 0.57 0.57 0.34 0.59 0.56 0.44 0.58 wt% 5 10 × Ca 1.99 1.74 1.57 2.39 1.65 1.66 1.54 2.11 1.62 1.53 1.95 1.56 wt% t=5 Continued on next page Al 1.37 1.19 1.08 1.65 1.13 1.14 1.06 1.45 1.11 1.05 1.34 1.07 wt% Table B.1 – continued from previous page Fe wt% 27.73 28.42 26.94 28.07 28.19 27.90 26.42 27.92 27.70 26.16 28.18 26.78 S 2.90 2.97 7.03 1.11 3.97 4.07 7.28 2.12 4.96 7.20 2.61 6.73 wt% O wt% 32.20 31.36 30.86 32.28 31.06 31.53 31.71 32.26 31.21 32.35 31.72 31.47 Mg wt% 14.68 14.86 14.09 15.06 14.74 14.61 13.83 14.87 14.48 13.70 14.86 14.01 H 0.01 0.01 0.01 0.00 0.00 0.00 0.08 0.00 0.00 0.12 0.00 0.05 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 210 0.90 0.90 0.91 0.89 0.90 0.91 0.91 0.90 0.91 0.91 1.04 0.90 0.90 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.35 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.76 16.43 15.61 16.49 16.66 16.23 15.14 16.85 16.61 15.94 14.72 16.93 16.81 15.87 15.80 Ti 0.09 0.08 0.08 0.12 0.09 0.08 0.07 0.10 0.08 0.08 3.60 0.10 0.10 0.08 0.08 wt% P 0.14 0.12 0.13 0.09 0.13 0.14 0.13 0.11 0.14 0.13 0.09 0.11 0.13 0.13 0.13 wt% Cr 0.42 0.40 0.39 0.38 0.41 0.41 0.38 0.40 0.41 0.40 0.19 0.41 0.41 0.40 0.40 wt% Ni 1.82 1.79 1.71 1.75 1.81 1.78 1.66 1.80 1.82 1.75 0.04 1.82 1.81 1.74 1.73 wt% Na 0.59 0.41 0.58 0.33 0.46 0.60 0.57 0.40 0.61 0.52 0.62 0.35 0.57 0.59 0.58 wt% Ca 1.90 1.77 1.58 2.55 1.85 1.64 1.53 2.24 1.68 1.61 1.57 2.10 2.10 1.60 1.60 wt% Continued on next page Al 1.31 1.21 1.08 1.75 1.27 1.13 1.05 1.55 1.15 1.11 2.44 1.45 1.45 1.10 1.10 wt% Table B.1 – continued from previous page Fe wt% 28.82 28.22 27.00 27.48 28.56 28.12 26.19 28.36 28.74 27.61 18.26 28.66 28.58 27.49 27.36 S 1.34 3.40 6.47 2.84 2.12 4.55 7.36 1.11 2.45 5.79 0.95 0.91 1.01 6.64 5.89 wt% O wt% 31.73 31.34 31.27 31.61 31.62 30.66 32.22 32.01 31.30 30.65 40.93 31.99 31.94 30.02 31.04 Mg wt% 15.11 14.84 14.13 14.64 15.03 14.70 13.72 15.08 15.03 14.43 15.25 15.20 15.09 14.36 14.31 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 1.02 0.00 0.00 0.00 0.00 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 211 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.00 15.20 14.03 15.98 15.32 14.75 14.06 15.09 14.79 13.36 15.77 14.41 Ti 0.07 0.07 0.07 0.08 0.07 0.07 0.07 0.07 0.07 0.07 0.08 0.07 wt% P 0.13 0.13 0.12 0.13 0.13 0.13 0.12 0.13 0.13 0.11 0.13 0.12 wt% Cr 0.37 0.38 0.35 0.40 0.38 0.37 0.35 0.38 0.37 0.33 0.39 0.36 wt% Ni 1.64 1.67 1.54 1.75 1.68 1.62 1.54 1.65 1.62 1.46 1.73 1.58 wt% years Na 0.55 0.57 0.52 0.58 0.57 0.55 0.53 0.55 0.55 0.50 0.58 0.54 wt% 6 10 × Ca 1.52 1.54 1.42 1.62 1.55 1.49 1.42 1.52 1.49 1.35 1.59 1.46 wt% t=1 Continued on next page Al 1.04 1.05 0.97 1.11 1.06 1.02 0.98 1.05 1.03 0.93 1.09 1.00 wt% Table B.1 – continued from previous page Fe wt% 25.95 26.31 24.26 27.70 26.48 25.52 24.33 26.10 25.59 23.11 27.30 24.93 S 4.88 7.00 6.76 5.01 7.02 6.51 6.83 3.82 6.95 6.46 5.36 6.96 wt% O wt% 34.66 32.15 36.44 31.13 31.77 34.08 36.30 35.25 33.61 38.97 31.63 34.98 Mg wt% 13.58 13.76 12.70 14.47 13.87 13.35 12.74 13.66 13.39 12.10 14.28 13.05 H 0.63 0.19 0.83 0.07 0.12 0.56 0.76 0.75 0.42 1.27 0.08 0.55 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 212 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.87 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 8.72 wt% 15.82 15.73 15.17 16.07 15.89 15.47 14.14 15.98 15.62 16.81 16.23 15.87 15.17 15.26 Ti 0.08 0.08 0.07 0.08 0.08 0.07 0.07 0.08 0.08 0.16 0.05 0.08 0.08 0.07 0.07 wt% P 0.13 0.13 0.13 0.13 0.13 0.13 0.12 0.13 0.13 0.12 0.08 0.14 0.13 0.13 0.13 wt% Cr 0.40 0.39 0.38 0.40 0.40 0.39 0.35 0.40 0.39 0.36 0.22 0.41 0.40 0.38 0.38 wt% Ni 1.73 1.72 1.66 1.76 1.74 1.70 1.55 1.75 1.71 1.61 0.96 1.78 1.74 1.66 1.67 wt% Na 0.58 0.58 0.57 0.58 0.59 0.58 0.53 0.58 0.58 0.50 0.33 0.59 0.58 0.57 0.57 wt% Ca 1.60 1.59 1.53 1.62 1.60 1.56 1.43 1.61 1.58 3.44 0.88 1.64 1.60 1.53 1.54 wt% Continued on next page Al 1.10 1.09 1.05 1.11 1.10 1.07 0.98 1.11 1.08 2.41 0.61 1.13 1.10 1.05 1.06 wt% Table B.1 – continued from previous page Fe wt% 27.38 27.20 26.26 27.81 27.47 26.77 24.42 27.65 27.04 25.26 15.10 28.07 27.46 26.26 26.40 S 6.39 5.29 7.35 4.00 5.69 7.30 6.86 5.04 7.25 1.81 4.12 4.32 6.37 7.33 7.30 wt% O wt% 30.48 31.90 32.09 31.86 30.94 30.97 36.09 31.22 30.41 32.95 56.34 30.95 30.33 32.10 31.81 Mg 7.89 wt% 14.33 14.24 13.74 14.55 14.39 14.00 12.81 14.46 14.14 14.60 14.69 14.36 13.73 13.82 H 0.00 0.07 0.01 0.04 0.00 0.00 0.67 0.00 0.00 0.00 4.74 0.00 0.00 0.02 0.01 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 213 0.91 0.91 0.91 0.86 0.86 0.73 0.91 0.79 0.71 0.47 0.83 0.56 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.87 13.96 12.94 15.90 16.55 14.79 12.72 15.11 14.86 13.07 15.84 13.70 Ti 0.06 0.06 0.06 0.07 0.06 0.06 0.06 0.06 0.06 0.06 0.07 0.06 wt% P 0.14 0.11 0.13 0.14 0.14 0.14 0.13 0.14 0.14 0.12 0.14 0.13 wt% Cr 0.35 0.35 0.32 0.37 0.36 0.29 0.32 0.32 0.27 0.05 0.35 0.15 wt% Ni 1.52 1.53 1.42 1.62 1.67 1.27 1.39 1.41 1.23 0.46 1.56 0.77 wt% years Na 0.52 0.52 0.48 0.51 0.29 0.26 0.48 0.43 0.19 0.01 0.40 0.03 6 wt% 10 × Ca 1.40 1.41 1.31 3.65 4.32 1.28 7.14 5.48 wt% 10.00 10.58 19.92 16.24 t=1.5 Continued on next page Al 0.96 0.97 0.90 2.52 2.99 6.92 0.88 4.95 7.31 3.78 wt% 13.75 11.23 Table B.1 – continued from previous page Fe 6.19 wt% 23.99 24.15 22.39 25.57 26.16 20.02 22.00 22.26 19.25 24.38 11.75 S 5.84 6.73 6.29 4.22 0.22 0.44 6.18 2.94 0.09 0.03 1.85 0.03 wt% O wt% 37.60 36.65 40.49 31.63 33.01 34.73 41.35 33.08 35.16 39.63 32.89 37.70 Mg 6.10 7.71 wt% 12.56 12.64 11.72 13.75 14.18 10.79 11.52 11.98 10.56 13.16 H 1.23 0.94 1.58 0.00 0.00 0.00 1.72 0.00 0.00 0.00 0.00 0.00 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 214 0.91 0.91 0.91 0.91 0.91 0.91 0.90 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.88 8.22 wt% 15.38 14.67 14.60 15.71 15.41 15.16 15.70 15.16 14.40 15.70 15.36 14.89 14.75 Ti 0.07 0.07 0.06 0.07 0.07 0.06 0.06 0.07 0.07 0.15 0.04 0.07 0.07 0.06 0.06 wt% P 0.14 0.14 0.14 0.14 0.14 0.14 0.13 0.14 0.14 0.13 0.09 0.15 0.14 0.14 0.14 wt% Cr 0.38 0.37 0.37 0.39 0.39 0.38 0.20 0.39 0.38 0.36 0.21 0.39 0.38 0.37 0.37 wt% Ni 1.69 1.61 1.60 1.72 1.69 1.66 0.87 1.72 1.66 1.58 0.90 1.72 1.68 1.63 1.62 wt% Na 0.57 0.55 0.55 0.58 0.58 0.57 0.30 0.59 0.57 0.54 0.31 0.59 0.57 0.56 0.55 wt% Ca 1.55 1.48 1.48 1.59 1.56 1.53 0.80 1.59 1.53 1.45 0.83 1.59 1.55 1.50 1.49 wt% Continued on next page Al 1.07 1.02 1.01 1.09 1.07 1.05 0.55 1.09 1.05 1.00 0.57 1.09 1.07 1.03 1.02 wt% Table B.1 – continued from previous page Fe wt% 26.62 25.39 25.27 27.16 26.67 26.23 13.65 27.16 26.23 24.91 14.20 27.16 26.58 25.74 25.52 S 7.28 6.22 7.10 5.38 7.16 7.34 3.84 6.20 7.32 6.92 3.98 6.36 7.08 7.23 7.16 wt% O wt% 31.31 34.63 34.29 31.92 31.33 32.14 59.39 31.11 32.16 35.04 58.16 30.97 31.62 33.18 33.74 Mg 7.13 7.44 wt% 13.92 13.29 13.22 14.22 13.95 13.73 14.22 13.72 13.04 14.22 13.91 13.50 13.35 H 0.04 0.60 0.33 0.05 0.03 0.04 5.32 0.05 0.04 0.59 5.09 0.02 0.01 0.19 0.27 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 215 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.23 wt% 13.10 13.68 12.31 15.13 14.27 13.39 10.80 13.87 13.48 14.26 11.76 Ti 0.04 0.06 0.05 0.05 0.05 0.06 0.05 0.05 0.06 0.04 0.06 0.05 wt% P 0.11 0.11 0.10 0.13 0.12 0.10 0.08 0.12 0.11 0.07 0.11 0.09 wt% Cr 0.33 0.34 0.31 0.38 0.36 0.33 0.27 0.35 0.34 0.23 0.36 0.29 wt% Ni 1.44 1.50 1.35 1.66 1.56 1.47 1.18 1.52 1.48 1.01 1.56 1.29 wt% years Na 0.49 0.51 0.46 0.57 0.53 0.50 0.40 0.52 0.50 0.34 0.53 0.44 wt% 6 10 × Ca 1.32 1.38 1.24 1.53 1.44 1.35 1.09 1.40 1.36 0.93 1.44 1.19 wt% t=2 Continued on next page Al 0.91 0.95 0.85 1.05 0.99 0.93 0.75 0.96 0.93 0.64 0.99 0.82 wt% Table B.1 – continued from previous page Fe wt% 22.66 23.65 21.29 26.17 24.70 23.16 18.67 23.99 23.31 15.95 24.66 20.33 S 6.11 6.62 5.98 6.77 6.94 6.49 5.25 6.38 6.55 4.48 6.72 5.71 wt% O wt% 39.98 37.74 42.88 32.53 35.49 38.79 48.51 37.13 38.49 54.35 35.63 44.92 Mg 9.78 8.36 wt% 11.86 12.39 11.15 13.70 12.92 12.13 12.56 12.21 12.91 10.65 H 1.65 1.09 2.03 0.35 0.62 1.30 3.16 1.16 1.20 4.35 0.77 2.47 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 216 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.60 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.79 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.55 wt% 15.04 14.29 11.84 15.12 15.01 14.78 15.24 14.98 14.84 13.99 15.35 15.17 14.67 13.23 Ti 0.05 0.05 0.05 0.05 0.05 0.06 0.03 0.06 0.06 0.06 4.53 0.05 0.05 0.06 0.06 wt% P 0.13 0.12 0.09 0.13 0.13 0.12 0.06 0.13 0.13 0.12 0.12 0.13 0.13 0.11 0.10 wt% Cr 0.38 0.36 0.30 0.38 0.38 0.37 0.19 0.38 0.37 0.37 0.30 0.38 0.38 0.37 0.33 wt% Ni 1.65 1.57 1.30 1.66 1.65 1.62 0.83 1.67 1.64 1.63 0.01 1.68 1.66 1.61 1.45 wt% Na 0.56 0.53 0.44 0.57 0.56 0.55 0.28 0.57 0.56 0.55 0.97 0.57 0.57 0.55 0.49 wt% Ca 1.52 1.44 1.20 1.53 1.52 1.49 0.76 1.54 1.51 1.50 2.55 1.55 1.53 1.48 1.34 wt% Continued on next page Al 1.04 0.99 0.82 1.05 1.04 1.02 0.52 1.06 1.04 1.03 3.49 1.06 1.05 1.02 0.92 wt% Table B.1 – continued from previous page Fe wt% 26.01 24.72 20.48 26.15 25.97 25.56 13.05 26.37 25.90 25.67 32.96 26.56 26.24 25.36 22.89 S 7.22 6.73 5.75 6.67 7.22 7.18 3.67 7.00 7.27 7.21 1.49 7.12 7.20 7.13 6.43 wt% O wt% 32.67 35.50 44.61 32.75 32.70 33.62 60.65 32.06 32.86 33.39 30.35 31.57 32.24 34.10 39.40 Mg 6.84 8.37 wt% 13.62 12.94 10.73 13.69 13.60 13.39 13.80 13.57 13.43 13.90 13.74 13.29 11.98 H 0.13 0.75 2.40 0.27 0.19 0.24 5.57 0.13 0.13 0.20 0.06 0.07 0.04 0.25 1.37 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 217 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.67 9.07 wt% 12.72 13.32 11.28 14.75 13.91 12.60 13.63 12.39 13.95 10.48 Ti 0.05 0.06 0.05 0.06 0.06 0.05 0.04 0.06 0.05 0.04 0.06 0.04 wt% P 0.10 0.10 0.09 0.12 0.11 0.10 0.08 0.11 0.10 0.07 0.11 0.08 wt% Cr 0.32 0.33 0.28 0.37 0.35 0.32 0.24 0.34 0.31 0.23 0.35 0.26 wt% Ni 1.39 1.46 1.24 1.62 1.52 1.38 1.06 1.49 1.36 0.99 1.53 1.15 wt% years Na 0.48 0.50 0.42 0.55 0.52 0.47 0.36 0.51 0.46 0.34 0.52 0.39 6 wt% 10 × Ca 1.29 1.35 1.14 1.49 1.40 1.27 0.98 1.38 1.25 0.92 1.41 1.06 wt% t=2.5 Continued on next page Al 0.88 0.92 0.78 1.02 0.96 0.87 0.67 0.95 0.86 0.63 0.97 0.73 wt% Table B.1 – continued from previous page Fe wt% 22.01 23.04 19.50 25.52 24.05 21.80 16.71 23.59 21.43 15.68 24.13 18.12 S 6.11 6.46 5.48 7.03 6.76 6.12 4.70 6.54 6.02 4.40 6.73 5.09 wt% O wt% 41.25 39.08 46.72 33.67 36.88 41.73 52.73 37.83 42.54 54.96 36.69 49.71 Mg 8.75 8.21 9.49 wt% 11.52 12.06 10.21 13.35 12.60 11.41 12.34 11.22 12.63 H 1.88 1.31 2.82 0.47 0.88 1.86 4.01 1.24 2.00 4.47 0.93 3.40 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 218 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.55 8.05 wt% 14.74 13.90 10.32 14.86 14.07 14.37 15.00 14.81 12.82 15.08 15.01 13.92 11.63 Ti 0.06 0.06 0.04 0.06 0.06 0.06 0.03 0.06 0.06 0.05 0.03 0.06 0.06 0.06 0.05 wt% P 0.12 0.11 0.08 0.12 0.11 0.11 0.06 0.12 0.12 0.10 0.06 0.12 0.12 0.11 0.09 wt% Cr 0.37 0.35 0.26 0.37 0.35 0.36 0.19 0.38 0.37 0.32 0.20 0.38 0.38 0.35 0.29 wt% Ni 1.62 1.52 1.13 1.63 1.54 1.57 0.83 1.64 1.62 1.41 0.88 1.65 1.65 1.52 1.27 wt% Na 0.55 0.52 0.39 0.56 0.53 0.54 0.28 0.56 0.55 0.48 0.30 0.56 0.56 0.52 0.43 wt% Ca 1.49 1.40 1.04 1.50 1.42 1.45 0.76 1.52 1.50 1.30 0.81 1.52 1.52 1.41 1.17 wt% Continued on next page Al 1.02 0.96 0.72 1.03 0.98 1.00 0.52 1.04 1.03 0.89 0.56 1.05 1.04 0.96 0.81 wt% Table B.1 – continued from previous page Fe wt% 25.50 24.04 17.83 25.73 24.34 24.84 13.05 25.96 25.61 22.19 13.91 26.09 25.96 24.06 20.12 S 7.13 6.72 5.01 7.04 6.82 6.98 3.67 7.20 7.19 6.23 3.91 7.25 7.24 6.76 5.65 wt% O wt% 33.74 36.87 50.34 33.28 36.25 35.20 60.65 32.71 33.50 40.94 58.65 32.44 32.75 36.88 45.40 Mg 9.34 6.84 7.29 wt% 13.35 12.58 13.46 12.74 13.01 13.58 13.42 11.61 13.65 13.60 12.61 10.53 H 0.30 0.96 3.51 0.36 0.80 0.52 5.57 0.22 0.21 1.66 5.35 0.15 0.12 0.83 2.54 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 219 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.43 8.57 7.79 9.35 wt% 12.54 12.81 14.53 12.25 12.35 14.31 11.65 13.63 Ti 0.05 0.05 0.04 0.06 0.05 0.05 0.04 0.06 0.05 0.03 0.06 0.04 wt% P 0.10 0.10 0.07 0.11 0.10 0.10 0.07 0.11 0.09 0.06 0.11 0.07 wt% Cr 0.31 0.32 0.24 0.36 0.31 0.31 0.21 0.36 0.29 0.19 0.34 0.23 wt% Ni 1.37 1.40 1.03 1.59 1.34 1.35 0.94 1.57 1.28 0.85 1.49 1.02 wt% years Na 0.47 0.48 0.35 0.54 0.46 0.46 0.32 0.53 0.44 0.29 0.51 0.35 wt% 6 10 × Ca 1.27 1.29 0.95 1.47 1.24 1.25 0.87 1.45 1.18 0.79 1.38 0.94 wt% t=3 Continued on next page Al 0.87 0.89 0.65 1.01 0.85 0.86 0.59 0.99 0.81 0.54 0.95 0.65 wt% Table B.1 – continued from previous page Fe wt% 21.69 22.14 16.31 25.14 21.18 21.35 14.83 24.75 20.15 13.47 23.58 16.17 S 6.07 6.22 4.58 7.03 5.95 6.00 4.17 6.93 5.66 3.78 6.62 4.54 wt% O wt% 41.91 41.00 53.61 34.43 43.08 42.71 56.78 35.34 45.31 59.67 37.86 53.88 Mg 8.53 7.76 7.05 8.47 wt% 11.36 11.60 13.15 11.10 11.18 12.96 10.55 12.35 H 1.99 1.70 4.20 0.57 2.11 2.04 4.86 0.62 2.55 5.47 1.13 4.26 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 220 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 0.91 Mg/Si C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.55 8.02 wt% 14.44 12.93 12.32 14.78 14.57 13.28 14.65 14.39 11.18 14.88 14.66 13.35 10.73 Ti 0.06 0.05 0.05 0.06 0.06 0.06 0.03 0.06 0.06 0.05 0.03 0.06 0.06 0.06 0.04 wt% P 0.11 0.10 0.10 0.12 0.11 0.10 0.06 0.11 0.11 0.09 0.06 0.12 0.11 0.10 0.08 wt% Cr 0.36 0.32 0.31 0.37 0.36 0.33 0.19 0.37 0.36 0.28 0.20 0.37 0.37 0.33 0.27 wt% Ni 1.58 1.42 1.35 1.62 1.60 1.46 0.83 1.60 1.58 1.23 0.88 1.63 1.61 1.47 1.18 wt% Na 0.54 0.48 0.46 0.55 0.54 0.50 0.28 0.55 0.54 0.42 0.30 0.56 0.55 0.50 0.40 wt% Ca 1.46 1.31 1.24 1.49 1.47 1.34 0.76 1.48 1.45 1.13 0.81 1.50 1.48 1.35 1.08 wt% Al 1.00 0.90 0.85 1.02 1.01 0.92 0.52 1.02 1.00 0.78 0.56 1.03 1.02 0.93 0.74 wt% Table B.1 – continued from previous page Fe wt% 24.96 22.36 21.31 25.56 25.18 22.97 13.06 25.33 24.88 19.34 13.86 25.73 25.34 23.13 18.55 S 7.00 6.28 5.99 7.13 7.06 6.45 3.67 7.10 6.99 5.44 3.90 7.21 7.10 6.50 5.21 wt% O wt% 34.91 40.49 42.85 33.61 34.38 39.25 60.63 34.04 35.11 47.08 58.74 33.18 34.06 38.97 48.78 Mg 6.84 7.27 9.72 wt% 13.08 11.70 11.15 13.38 13.19 12.03 13.27 13.03 10.12 13.47 13.28 12.09 H 0.49 1.66 2.01 0.31 0.46 1.31 5.57 0.44 0.50 2.88 5.37 0.25 0.36 1.23 3.20 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 221 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.31 14.67 13.13 15.35 15.05 14.37 12.45 14.80 14.25 11.92 15.11 13.06 Ti 0.07 0.07 0.06 0.16 0.07 0.07 0.05 0.12 0.06 0.05 0.10 0.06 wt% P 0.11 0.12 0.11 0.11 0.12 0.11 0.11 0.11 0.11 0.10 0.11 0.11 wt% Cr 0.35 0.36 0.33 0.34 0.37 0.34 0.31 0.33 0.34 0.26 0.35 0.30 wt% Ni 1.56 1.60 1.44 1.52 1.62 1.51 1.37 1.49 1.50 1.17 1.57 1.33 wt% Na 0.46 0.51 0.49 0.43 0.47 0.45 0.46 0.44 0.45 0.38 0.46 0.41 wt% Ca 1.60 1.60 1.35 4.04 2.02 2.83 1.26 3.72 2.79 3.88 2.99 3.45 wt% Al Continued on next page 1.10 1.10 0.93 2.80 1.39 1.95 0.86 2.58 1.92 2.67 2.07 2.38 wt% Fe wt% 24.52 25.27 22.68 23.91 25.59 23.84 21.54 23.47 23.68 18.25 24.71 20.88 S 4.87 5.42 5.64 4.54 4.52 4.54 5.62 4.33 4.64 4.62 4.43 4.73 wt% O wt% 37.09 35.89 40.35 33.49 34.76 36.54 42.60 35.33 36.85 44.42 34.43 40.53 Mg 9.97 wt% 12.91 12.67 11.88 13.11 13.47 12.58 11.28 12.71 12.47 13.23 11.15 H 1.05 0.75 1.64 0.21 0.53 0.82 2.08 0.54 0.88 2.24 0.42 1.53 wt% 4 5 6 4 5 6 7 4 5 6 4 5 − − − − − − − − − − − − PLanet CJS1 CJS1 CJS1 CJS2 CJS2 CJS2 CJS2 CJS3 CJS3 CJS3 CJS4 CJS4 Table B.2: Ensemble-averaged bulk predicted planetaryconditions abundances. simulated. All values are in wt% and are averaged over all seven sets of midplane 222 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.16 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.56 14.80 13.74 15.50 15.48 14.99 10.88 15.30 15.46 14.71 11.06 15.58 15.70 14.96 13.82 Ti 0.08 0.10 0.06 0.07 0.08 0.11 0.05 0.14 0.09 0.09 1.19 0.12 0.07 0.07 0.06 wt% P 0.12 0.11 0.11 0.12 0.12 0.12 0.10 0.12 0.12 0.12 0.09 0.11 0.12 0.12 0.12 wt% Cr 0.38 0.35 0.34 0.38 0.38 0.36 0.27 0.36 0.37 0.36 0.24 0.37 0.39 0.37 0.34 wt% Ni 1.69 1.55 1.51 1.69 1.68 1.58 1.19 1.60 1.64 1.57 0.77 1.64 1.71 1.61 1.51 wt% Na 0.50 0.48 0.51 0.53 0.48 0.52 0.40 0.50 0.51 0.49 0.49 0.47 0.52 0.54 0.51 wt% Ca 1.92 2.50 1.39 1.71 1.90 2.50 1.10 3.37 2.21 1.80 1.29 2.96 1.72 1.76 1.43 wt% Al 1.32 1.72 0.95 1.18 1.31 1.73 0.75 2.33 1.52 1.24 1.33 2.04 1.18 1.22 0.99 wt% Fe wt% Table B.2 – continued from previous page 26.55 24.51 23.77 26.65 26.39 24.87 18.81 25.25 25.72 24.84 19.35 25.83 26.99 25.33 23.81 S 5.24 5.12 5.98 5.51 5.20 5.86 4.75 4.90 5.61 4.82 3.70 4.77 5.17 6.20 6.36 wt% O wt% 32.54 35.09 38.02 32.50 32.88 33.82 48.62 32.71 32.85 36.02 47.57 32.35 32.16 34.03 37.49 Mg 9.85 9.67 wt% 13.98 13.08 12.44 13.99 13.91 13.26 13.31 13.78 13.20 13.68 14.18 13.45 12.50 H 0.14 0.58 1.18 0.17 0.21 0.30 3.24 0.12 0.13 0.76 3.09 0.07 0.08 0.36 1.06 wt% 4 5 6 4 5 6 7 4 5 6 7 4 5 6 7 − − − − − − − − − − − − − − − Planet EJS1 EJS1 EJS1 EJS2 EJS2 EJS2 EJS2 EJS3 EJS3 EJS3 EJS3 EJS4 EJS4 EJS4 EJS4 223 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 . 5 10 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 × 5 . 2 4.07 3.86 6.92 4.52 4.18 7.52 0.72 4.94 7.99 2.20 6.91 Si 0.26 − − − − − − − − − − − Abundance − 0.04 0.04 0.03 0.60 0.05 0.06 0.02 0.39 0.03 0.03 0.25 0.02 Ti 6 − − − − − − − − − − − − 10 × 3 P 0.01 0.06 0.00 0.07 0.02 0.07 0.06 0.01 0.08 0.01 0.06 0.04 − − − − − − − 0.06 0.08 0.17 0.10 0.06 0.19 0.12 0.20 0.17 Cr 0.24 0.16 0.06 − − − − − − − − − 0.41 0.38 0.75 0.46 0.39 0.82 0.52 0.87 0.75 Ni 0.99 0.53 0.15 − − − − − − − − − 0.23 0.28 0.29 0.22 Na 0.16 0.12 0.45 0.13 0.12 0.36 0.02 0.30 − − − − yr. Difference is deﬁned as Abundance 6 10 × 1.13 1.12 0.89 1.34 1.57 0.76 9.63 0.90 0.93 6.32 0.79 Ca 14.65 − − − − − − − − − − − − Continued on next page yr and t = 3 0.78 0.77 0.61 0.93 1.08 0.52 6.71 0.62 0.64 4.40 0.54 Al 10.23 5 − − − − − − − − − − − − 10 × 5.95 5.98 7.23 5.74 8.19 Fe 11.78 13.00 13.75 11.90 9.09 2.82 15.91 − − − − − − − − − S 0.81 2.20 3.91 4.60 1.24 6.45 5.17 3.82 5.37 3.44 5.48 0.47 − − 4.33 1.10 O 9.86 9.06 3.26 22.18 11.02 10.52 25.93 13.65 28.68 22.85 − − 3.80 3.42 6.80 3.48 3.26 6.19 4.33 7.19 0.07 6.22 Mg 4.86 2.35 − − − − − − − − − − H 1.99 1.70 4.20 0.57 2.11 2.04 4.86 0.62 2.55 5.47 1.13 4.26 4 5 6 4 5 6 7 4 5 6 4 5 Table B.3: Difference in abundance between t = 2.5 − − − − − − − − − − − − Planet CJS1 CJS1 CJS1 CJS2 CJS2 CJS2 CJS2 CJS3 CJS3 CJS3 CJS4 CJS4 224 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 2.33 2.72 4.02 2.11 3.67 8.12 0.22 2.66 5.63 7.33 0.53 1.96 3.62 5.83 Si 1.07 − − − − − − − − − − − − − − 0.09 0.29 0.01 0.49 0.19 0.02 0.02 0.38 0.02 0.09 0.03 0.40 0.10 0.01 0.03 Ti − − − − − − − − − − − − − − − P 0.04 0.02 0.07 0.03 0.06 0.03 0.05 0.05 0.03 0.05 0.06 0.08 0.01 0.08 0.07 − − − − − − − 0.10 0.08 0.20 0.05 0.08 0.17 0.01 0.09 0.14 Cr 0.01 0.05 0.15 0.08 0.13 0.17 − − − − − − − − − 0.15 0.44 0.38 0.89 0.26 0.38 0.76 0.11 0.40 0.62 Ni 0.16 0.60 0.18 0.38 0.56 − − − − − − − − − − 0.11 0.30 0.08 0.24 0.08 0.10 Na 0.43 0.20 0.32 0.37 0.09 0.45 0.32 0.48 0.46 − − − − − − 2.44 7.17 0.41 4.72 0.83 0.82 9.37 0.82 2.32 1.01 9.77 2.57 0.37 0.87 Ca 12.10 − − − − − − − − − − − − − − − 1.69 5.00 0.28 8.39 3.31 0.57 0.56 6.49 0.57 1.63 0.70 6.81 1.78 0.25 0.60 Al − − − − − − − − − − − − − − − Table B.3 – continued from previous page 1.56 6.94 5.94 3.90 5.92 1.63 6.23 9.74 Fe 14.10 11.99 2.73 9.60 3.57 6.79 9.69 − − − − − − − − − − S 3.90 2.99 6.69 5.06 1.81 6.48 6.24 6.09 6.85 6.77 3.62 7.07 6.97 6.01 2.49 − − 2.23 0.26 1.32 3.42 O 2.00 5.57 7.19 2.92 1.34 7.11 12.16 30.79 14.12 26.44 17.47 − − − − 1.40 0.27 3.64 0.38 3.17 7.35 2.24 4.47 6.53 1.04 3.28 5.19 Mg 4.64 2.66 2.61 − − − − − − − − − − − − H 0.49 1.66 2.01 0.31 0.46 1.31 5.57 0.44 0.50 2.88 5.37 0.25 0.36 1.23 3.20 4 5 6 4 5 6 7 4 5 6 7 4 5 6 7 − − − − − − − − − − − − − − − Planet EJS1 EJS1 EJS1 EJS2 EJS2 EJS2 EJS2 EJS3 EJS3 EJS3 EJS3 EJS4 EJS4 EJS4 EJS4 225

10 CJS1-4 (Venus)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS1-5 (Earth)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS1-6 (Mars)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.8: Normalized planetary abundances for CJS1 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 226

10 10 CJS2-4 (Venus) CJS2-6 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS2-5 (Venus) CJS2-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.9: Normalized planetary abundances for CJS2 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 227

10 10 CJS3-4 (Venus) CJS4-4 (Venus)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 CJS3-5 (Earth) CJS4-5 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 CJS3-6 (Mars) Increasing volatility

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.10: Normalized planetary abundances for CJS3 and CJS4 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 228

10 EJS1-4 (Venus)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 EJS1-5 (Earth)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S 10 EJS1-6 (Mars)

1 Normalized Abundance Normalized 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility

Figure B.11: Normalized planetary abundances for EJS1 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 229

10 10 EJS2-4 (Venus) EJS2-6 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS2-5 (Venus) EJS2-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.12: Normalized planetary abundances for EJS2 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 230

10 10 EJS3-4 (Venus) EJS3-6 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS3-5 (Earth) EJS3-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.13: Normalized planetary abundances for EJS3 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 231

10 10 EJS4-4 (Venus) EJS4-6 (Earth)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S 10 10 EJS4-5 (Venus) EJS4-7 (Mars)

1 1 Normalized Abundance Normalized Abundance Normalized 0.1 0.1 Al Ti Ca Mg Si ONiFeCrPNa S Al Ti Ca Mg Si ONiFeCrPNa S Increasing volatility Increasing volatility

Figure B.14: Normalized planetary abundances for EJS4 simulated terrestrial planets. The solid line indicates the normalized abundances before volatile loss during impacts was considered while the dashed line indicates the normalized abundance once volatile loss during impacts has been incorporated. All abundances were determined for disk conditions at t = 5×105 years. The terrestrial planet each simulation is normalized to is shown in parentheses. 232 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.02 15.48 15.26 12.90 15.91 15.06 15.09 13.31 15.08 15.09 13.74 14.69 Ti 0.09 0.09 0.06 0.63 0.09 0.10 0.06 0.43 0.07 0.06 0.29 0.06 wt% P 0.06 0.08 0.10 0.02 0.07 0.07 0.10 0.03 0.08 0.11 0.04 0.09 wt% Cr 0.27 0.31 0.32 0.09 0.34 0.27 0.32 0.13 0.29 0.34 0.17 0.28 wt% Ni 1.35 1.45 1.47 0.45 1.56 1.34 1.47 0.74 1.37 1.52 0.90 1.34 wt% Na 0.14 0.20 0.34 0.04 0.22 0.16 0.35 0.06 0.19 0.40 0.06 0.25 wt% years 5 10 Ca 2.39 2.41 1.84 2.58 2.81 1.63 2.07 1.71 7.70 1.73 wt% 16.12 11.07 × t=2.5 Al 1.65 1.66 1.27 1.77 1.94 1.12 7.70 1.43 1.18 5.35 1.19 wt% 11.24 Continued on next page Fe 6.98 wt% 20.91 22.95 23.24 24.56 20.94 23.29 11.16 21.76 24.07 13.99 21.17 S 0.62 0.65 1.43 0.17 0.41 0.68 2.23 0.33 0.66 3.42 0.18 1.11 wt% O wt% 23.13 25.19 25.19 27.98 27.06 23.82 25.05 24.53 23.24 26.84 21.81 22.31 Mg 6.55 7.98 8.91 wt% 11.73 12.52 12.56 13.18 11.76 12.54 11.92 12.84 11.59 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Table B.4: Final bulk planetary abundances after volatile loss in impacts has been incorporated. All values are in wt%. CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 233 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.43 13.67 15.27 11.68 14.60 15.60 15.67 13.05 16.18 15.35 15.24 13.57 14.97 16.24 15.70 Ti 0.14 0.32 0.06 0.51 0.24 0.08 0.06 0.41 0.08 0.13 0.07 0.43 0.15 0.06 0.07 wt% P 0.04 0.04 0.10 0.03 0.03 0.09 0.13 0.02 0.08 0.08 0.12 0.02 0.03 0.12 0.11 wt% Cr 0.21 0.17 0.32 0.13 0.18 0.31 0.39 0.15 0.34 0.26 0.37 0.13 0.27 0.36 0.34 wt% Ni 1.14 0.85 1.48 0.65 0.98 1.45 1.72 0.85 1.58 1.24 1.61 0.74 1.28 1.65 1.54 wt% Na 0.03 0.09 0.32 0.06 0.06 0.21 0.58 0.03 0.14 0.24 0.51 0.03 0.04 0.41 0.32 wt% Ca 3.90 8.47 1.65 6.20 2.17 1.58 2.28 3.44 1.82 4.05 1.72 1.95 wt% 13.59 10.85 11.27 Al 2.69 5.90 1.13 9.42 4.32 1.49 1.09 7.51 1.56 2.41 1.26 7.84 2.80 1.18 1.34 wt% Continued on next page Fe wt% 17.50 13.47 23.40 10.24 14.98 22.99 27.12 12.93 24.87 19.65 25.38 11.27 20.23 25.99 24.38 Table B.4 – continued from previous page S 0.05 0.19 1.74 0.09 0.16 0.13 7.52 0.05 0.11 0.58 6.34 0.03 0.04 0.28 1.37 wt% O wt% 20.24 22.44 24.62 21.33 22.56 24.52 29.79 23.21 27.14 24.57 31.60 24.21 23.37 27.63 26.30 Mg 8.72 6.02 9.97 7.83 8.07 wt% 10.20 12.63 12.53 14.18 13.50 11.82 13.58 11.24 13.86 13.16 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 234 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.77 15.28 14.53 15.26 15.45 14.72 14.32 14.67 14.54 14.47 14.33 13.98 Ti 0.09 0.08 0.07 0.10 0.08 0.08 0.07 0.09 0.07 0.07 0.09 0.07 wt% P 0.08 0.10 0.10 0.06 0.11 0.09 0.10 0.07 0.09 0.11 0.07 0.09 wt% Cr 0.28 0.32 0.31 0.28 0.34 0.29 0.31 0.26 0.29 0.32 0.25 0.27 wt% Ni 1.33 1.46 1.41 1.35 1.54 1.36 1.40 1.26 1.34 1.46 1.20 1.27 wt% Na 0.19 0.33 0.33 0.09 0.39 0.27 0.34 0.13 0.27 0.39 0.14 0.25 wt% years 5 Ca 10 1.99 1.74 1.57 2.39 1.65 1.66 1.54 2.11 1.62 1.53 1.95 1.56 wt% × t=5 Al 1.37 1.19 1.08 1.65 1.13 1.14 1.06 1.45 1.11 1.05 1.34 1.07 wt% Continued on next page Fe wt% 20.98 23.20 22.30 21.24 24.38 21.57 22.12 19.90 21.26 23.13 19.00 20.20 Table B.4 – continued from previous page S 0.84 1.19 3.01 0.32 2.09 1.27 3.26 0.45 1.48 4.11 0.42 1.84 wt% O wt% 23.24 24.73 24.74 23.30 26.21 23.33 25.76 21.71 22.91 28.01 19.99 22.63 Mg wt% 11.61 12.53 12.01 11.91 13.04 11.77 11.91 11.18 11.59 12.35 10.67 11.05 H 0.00 0.00 0.00 0.00 0.00 0.00 0.02 0.00 0.00 0.04 0.00 0.00 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 235 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.42 14.36 14.58 14.04 14.59 14.94 15.14 14.79 15.76 14.56 15.07 14.90 15.15 15.18 14.98 Ti 0.08 0.08 0.07 0.11 0.08 0.08 0.07 0.10 0.08 0.07 0.06 0.09 0.09 0.07 0.07 wt% P 0.08 0.07 0.10 0.04 0.08 0.10 0.13 0.06 0.11 0.09 0.12 0.07 0.09 0.11 0.11 wt% Cr 0.25 0.25 0.31 0.22 0.26 0.30 0.38 0.25 0.35 0.29 0.37 0.26 0.29 0.34 0.33 wt% Ni 1.20 1.22 1.41 1.12 1.25 1.41 1.66 1.25 1.57 1.35 1.63 1.27 1.35 1.54 1.49 wt% Na 0.17 0.13 0.33 0.09 0.16 0.30 0.56 0.14 0.39 0.25 0.52 0.12 0.25 0.42 0.38 wt% Ca 1.90 1.77 1.58 2.55 1.85 1.64 1.53 2.24 1.68 1.61 1.53 2.10 2.10 1.60 1.60 wt% Al 1.31 1.21 1.08 1.75 1.27 1.13 1.05 1.55 1.15 1.11 1.05 1.45 1.45 1.10 1.10 wt% Continued on next page Fe wt% 19.01 19.36 22.36 17.63 19.79 22.35 26.15 19.79 24.83 21.48 25.78 20.13 21.43 24.34 23.58 Table B.4 – continued from previous page S 0.21 0.54 2.69 0.39 0.41 1.62 7.31 0.22 1.26 1.87 6.42 0.18 0.28 3.87 2.97 wt% O wt% 19.51 20.14 25.08 18.81 20.60 23.45 32.17 21.02 26.38 22.85 31.79 21.16 22.82 26.03 26.07 Mg wt% 10.64 10.82 12.06 10.08 11.04 12.12 13.70 11.14 13.30 11.68 13.52 11.29 11.85 12.96 12.62 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 236 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.56 14.12 13.10 14.44 14.53 13.44 13.19 13.35 13.45 12.78 13.69 13.02 Ti 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.06 0.07 0.07 wt% P 0.08 0.10 0.09 0.09 0.10 0.09 0.09 0.08 0.09 0.10 0.07 0.08 wt% Cr 0.26 0.29 0.28 0.28 0.32 0.27 0.28 0.25 0.27 0.29 0.24 0.25 wt% Ni 1.24 1.36 1.27 1.32 1.45 1.24 1.29 1.17 1.24 1.29 1.16 1.19 wt% Na 0.24 0.31 0.30 0.26 0.38 0.26 0.31 0.20 0.25 0.35 0.18 0.23 wt% years 6 Ca 10 1.52 1.54 1.42 1.62 1.55 1.49 1.42 1.52 1.49 1.35 1.59 1.46 wt% × t=1 Al 1.04 1.05 0.97 1.11 1.06 1.02 0.98 1.05 1.03 0.93 1.09 1.00 wt% Continued on next page Fe wt% 19.63 21.47 20.08 20.96 22.90 19.73 20.36 18.60 19.65 20.43 18.40 18.81 Table B.4 – continued from previous page S 1.41 2.81 2.89 1.44 3.69 2.02 3.05 0.81 2.08 3.69 0.86 1.90 wt% O wt% 25.02 25.35 29.21 22.47 26.81 25.22 29.48 23.72 24.67 33.74 19.94 25.15 Mg wt% 10.74 11.60 10.83 11.44 12.27 10.75 10.97 10.27 10.72 10.91 10.25 10.30 H 0.06 0.03 0.16 0.01 0.04 0.06 0.16 0.03 0.04 0.42 0.00 0.04 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 237 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 8.66 wt% 13.61 13.74 14.18 13.69 13.91 14.24 14.13 14.02 14.82 13.95 14.29 14.29 14.52 14.46 Ti 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.07 0.06 0.05 0.07 0.07 0.07 0.07 wt% P 0.07 0.08 0.10 0.07 0.08 0.09 0.12 0.08 0.11 0.08 0.08 0.08 0.09 0.11 0.10 wt% Cr 0.24 0.24 0.30 0.23 0.25 0.29 0.35 0.25 0.33 0.28 0.21 0.26 0.28 0.33 0.32 wt% Ni 1.14 1.17 1.37 1.12 1.20 1.34 1.55 1.21 1.48 1.30 0.94 1.24 1.30 1.47 1.44 wt% Na 0.17 0.19 0.32 0.16 0.20 0.30 0.53 0.20 0.38 0.27 0.31 0.21 0.25 0.40 0.37 wt% Ca 1.60 1.59 1.53 1.62 1.60 1.56 1.43 1.61 1.58 1.54 0.88 1.64 1.60 1.53 1.54 wt% Al 1.10 1.09 1.05 1.11 1.10 1.07 0.98 1.11 1.08 1.06 0.61 1.13 1.10 1.05 1.06 wt% Continued on next page Fe wt% 18.06 18.66 21.75 17.85 19.04 21.28 24.39 19.29 23.37 20.56 14.82 19.71 20.59 23.25 22.75 Table B.4 – continued from previous page S 1.00 0.84 3.06 0.54 1.11 2.61 6.82 1.01 3.72 2.15 3.80 0.87 1.79 4.28 3.68 wt% O wt% 18.74 20.50 25.73 18.96 20.15 23.68 36.03 20.50 25.63 24.07 55.12 20.47 21.67 27.84 26.72 Mg 7.77 wt% 10.09 10.38 11.73 10.02 10.56 11.54 12.79 10.68 12.51 11.20 10.92 11.27 12.39 12.19 H 0.00 0.00 0.00 0.00 0.00 0.00 0.66 0.00 0.00 0.01 4.06 0.00 0.00 0.01 0.00 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 238 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 12.54 12.97 12.09 15.26 15.45 14.72 14.32 14.67 14.54 14.47 14.33 13.98 Ti 0.06 0.06 0.06 0.10 0.07 0.07 0.06 0.08 0.06 0.06 0.08 0.06 wt% P 0.09 0.08 0.10 0.05 0.10 0.09 0.09 0.06 0.08 0.10 0.06 0.08 wt% Cr 0.24 0.27 0.26 0.28 0.34 0.29 0.31 0.26 0.29 0.32 0.25 0.27 wt% Ni 1.14 1.24 1.17 1.35 1.54 1.36 1.40 1.26 1.34 1.46 1.20 1.27 wt% Na 0.23 0.29 0.28 0.09 0.39 0.27 0.34 0.13 0.27 0.39 0.14 0.25 wt% years 6 10 Ca 1.40 1.41 1.31 2.39 1.65 1.66 1.54 2.11 1.62 1.53 1.95 1.56 wt% × t=1.5 Al 0.96 0.97 0.90 1.65 1.13 1.14 1.06 1.45 1.11 1.05 1.34 1.07 wt% Continued on next page Fe wt% 18.15 19.71 18.53 21.24 24.38 21.57 22.12 19.90 21.26 23.13 19.00 20.20 Table B.4 – continued from previous page S 1.68 2.70 2.69 0.32 2.09 1.27 3.26 0.45 1.48 4.11 0.42 1.84 wt% O wt% 27.14 28.90 32.45 23.30 26.21 23.33 25.76 21.71 22.91 28.01 19.99 22.63 Mg 9.93 wt% 10.65 10.00 11.91 13.04 11.77 11.91 11.18 11.59 12.35 10.67 11.05 H 0.11 0.16 0.31 0.00 0.00 0.00 0.02 0.00 0.00 0.04 0.00 0.00 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 239 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.88 8.16 wt% 13.23 12.82 13.65 13.38 13.49 13.95 13.79 14.38 13.15 13.82 13.84 14.25 13.98 Ti 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.04 0.06 0.06 0.06 0.06 wt% P 0.08 0.08 0.11 0.08 0.08 0.10 0.13 0.09 0.11 0.08 0.09 0.09 0.09 0.12 0.11 wt% Cr 0.23 0.23 0.29 0.23 0.24 0.28 0.20 0.25 0.32 0.26 0.20 0.25 0.27 0.32 0.31 wt% Ni 1.10 1.10 1.32 1.10 1.16 1.32 0.86 1.19 1.43 1.22 0.88 1.20 1.26 1.44 1.39 wt% Na 0.17 0.17 0.31 0.16 0.20 0.29 0.29 0.21 0.37 0.26 0.29 0.21 0.25 0.39 0.35 wt% Ca 1.55 1.48 1.48 1.59 1.56 1.53 0.80 1.59 1.53 1.45 0.83 1.59 1.55 1.50 1.49 wt% Al 1.07 1.02 1.01 1.09 1.07 1.05 0.55 1.09 1.05 1.00 0.57 1.09 1.07 1.03 1.02 wt% Continued on next page Fe wt% 17.56 17.41 20.93 17.43 18.48 20.85 13.63 18.95 22.67 19.38 13.94 19.08 19.93 22.78 21.99 Table B.4 – continued from previous page S 1.14 0.98 2.95 0.73 1.39 2.62 3.81 1.24 3.75 2.23 3.66 1.28 1.99 4.22 3.61 wt% O wt% 19.25 22.26 27.50 19.00 20.41 24.58 59.29 20.42 27.11 26.12 56.90 20.49 22.59 28.78 28.34 Mg 9.81 9.68 9.79 7.12 7.33 wt% 11.29 10.24 11.32 10.50 12.14 10.56 10.56 10.92 12.18 11.78 H 0.00 0.00 0.06 0.00 0.00 0.01 5.26 0.00 0.01 0.06 4.36 0.00 0.00 0.07 0.07 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 240 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 8.83 wt% 11.84 12.71 11.50 13.68 13.54 12.20 10.13 12.28 12.25 12.37 10.62 Ti 0.04 0.06 0.05 0.05 0.05 0.05 0.04 0.05 0.05 0.04 0.06 0.05 wt% P 0.07 0.08 0.07 0.09 0.10 0.07 0.07 0.07 0.07 0.06 0.06 0.06 wt% Cr 0.23 0.27 0.24 0.27 0.30 0.24 0.22 0.23 0.24 0.20 0.22 0.21 wt% Ni 1.08 1.22 1.11 1.25 1.35 1.13 0.99 1.08 1.13 0.89 1.05 0.97 wt% Na 0.22 0.28 0.26 0.25 0.35 0.23 0.24 0.19 0.23 0.24 0.17 0.19 wt% years 6 Ca 10 1.32 1.38 1.24 1.53 1.44 1.35 1.09 1.40 1.36 0.93 1.44 1.19 wt% × t=2 Al 0.91 0.95 0.85 1.05 0.99 0.93 0.75 0.96 0.93 0.64 0.99 0.82 wt% Continued on next page Fe wt% 17.15 19.30 17.62 19.80 21.36 17.90 15.63 17.10 17.89 14.11 16.62 15.34 Table B.4 – continued from previous page S 1.76 2.66 2.56 1.95 3.65 2.02 2.35 1.36 1.96 2.56 1.08 1.56 wt% O wt% 28.86 29.76 34.37 23.48 29.95 28.70 39.40 24.99 28.25 47.06 22.45 32.29 Mg 9.38 9.51 9.77 8.42 9.45 9.77 7.54 9.27 8.40 wt% 10.44 10.83 11.44 H 0.15 0.18 0.39 0.03 0.18 0.13 0.66 0.05 0.11 1.45 0.02 0.18 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 241 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.55 8.02 wt% 12.94 12.49 11.07 12.88 13.15 13.60 13.38 14.21 13.55 13.51 13.67 14.04 12.54 Ti 0.05 0.05 0.05 0.05 0.05 0.06 0.03 0.06 0.06 0.06 0.03 0.05 0.05 0.06 0.05 wt% P 0.07 0.07 0.07 0.07 0.08 0.08 0.06 0.08 0.10 0.08 0.06 0.08 0.09 0.10 0.08 wt% Cr 0.22 0.22 0.23 0.22 0.24 0.28 0.19 0.24 0.31 0.27 0.20 0.25 0.26 0.31 0.27 wt% Ni 1.08 1.07 1.07 1.06 1.13 1.28 0.83 1.16 1.41 1.26 0.87 1.17 1.24 1.42 1.25 wt% Na 0.17 0.17 0.25 0.15 0.19 0.28 0.28 0.20 0.36 0.27 0.29 0.20 0.25 0.39 0.32 wt% Ca 1.52 1.44 1.20 1.53 1.52 1.49 0.76 1.54 1.51 1.50 0.81 1.55 1.53 1.48 1.34 wt% Al 1.04 0.99 0.82 1.05 1.04 1.02 0.52 1.06 1.04 1.03 0.56 1.06 1.05 1.02 0.92 wt% Continued on next page Fe wt% 17.16 16.96 16.96 16.78 18.00 20.32 13.03 18.40 22.38 19.97 13.69 18.65 19.68 22.45 19.72 Table B.4 – continued from previous page S 1.13 1.06 2.39 0.91 1.40 2.56 3.64 1.40 3.73 2.32 3.61 1.44 2.02 4.16 3.24 wt% O wt% 20.09 22.81 35.77 19.49 21.30 25.71 60.55 21.05 27.70 24.89 57.43 20.88 23.04 29.58 33.10 Mg 9.59 9.43 9.16 9.42 9.99 6.83 7.20 wt% 11.04 10.19 12.00 10.88 10.33 10.78 11.99 10.57 H 0.00 0.01 0.40 0.01 0.01 0.03 5.51 0.01 0.03 0.02 4.45 0.00 0.00 0.09 0.35 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 242 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.07 8.67 9.47 wt% 11.50 12.38 10.53 13.33 13.20 11.49 12.07 11.27 12.10 Ti 0.05 0.05 0.05 0.06 0.06 0.05 0.04 0.05 0.05 0.04 0.05 0.04 wt% P 0.07 0.08 0.07 0.08 0.09 0.07 0.06 0.06 0.07 0.06 0.06 0.05 wt% Cr 0.22 0.26 0.22 0.26 0.29 0.23 0.19 0.22 0.22 0.19 0.21 0.18 wt% Ni 1.05 1.19 1.02 1.22 1.31 1.06 0.88 1.06 1.04 0.88 1.02 0.86 wt% Na 0.21 0.27 0.24 0.24 0.34 0.22 0.21 0.19 0.21 0.24 0.16 0.17 wt% years 6 10 Ca 1.29 1.35 1.14 1.49 1.40 1.27 0.98 1.38 1.25 0.92 1.41 1.06 wt% × t=2.5 Al 0.88 0.92 0.78 1.02 0.96 0.87 0.67 0.95 0.86 0.63 0.97 0.73 wt% Continued on next page Fe wt% 16.65 18.80 16.14 19.31 20.80 16.85 13.99 16.81 16.45 13.86 16.26 13.67 Table B.4 – continued from previous page S 1.76 2.59 2.35 2.03 3.56 1.90 2.10 1.39 1.80 2.52 1.08 1.39 wt% O wt% 29.78 30.82 37.45 24.30 31.12 30.88 42.83 25.46 31.23 47.59 23.12 35.74 Mg 9.11 8.71 9.19 7.54 9.28 8.99 7.40 9.07 7.48 wt% 10.17 10.56 11.15 H 0.17 0.22 0.55 0.04 0.26 0.19 0.84 0.05 0.18 1.49 0.02 0.25 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 243 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.64 7.55 7.99 wt% 12.68 12.15 12.66 12.32 13.22 13.17 14.06 11.71 13.27 13.53 13.32 11.03 Ti 0.06 0.05 0.04 0.06 0.06 0.06 0.03 0.06 0.06 0.05 0.03 0.06 0.06 0.06 0.05 wt% P 0.06 0.06 0.06 0.06 0.06 0.08 0.06 0.07 0.09 0.07 0.06 0.07 0.08 0.09 0.07 wt% Cr 0.22 0.22 0.20 0.21 0.22 0.27 0.19 0.24 0.31 0.23 0.20 0.24 0.26 0.30 0.24 wt% Ni 1.06 1.04 0.93 1.04 1.06 1.25 0.83 1.14 1.40 1.09 0.86 1.15 1.23 1.35 1.10 wt% Na 0.16 0.16 0.22 0.15 0.18 0.27 0.28 0.20 0.36 0.23 0.28 0.20 0.24 0.37 0.28 wt% Ca 1.49 1.40 1.04 1.50 1.42 1.45 0.76 1.52 1.50 1.30 0.81 1.52 1.52 1.41 1.17 wt% Al 1.02 0.96 0.72 1.03 0.98 1.00 0.52 1.04 1.03 0.89 0.56 1.05 1.04 0.96 0.81 wt% Continued on next page Fe wt% 16.83 16.49 14.77 16.51 16.87 19.75 13.03 18.11 22.13 17.26 13.65 18.33 19.47 21.30 17.33 Table B.4 – continued from previous page S 1.11 1.06 2.08 0.96 1.33 2.49 3.64 1.44 3.69 2.01 3.60 1.46 2.03 3.95 2.85 wt% O wt% 20.75 23.69 40.37 19.80 23.61 26.92 60.55 21.48 28.24 30.52 57.38 21.46 23.40 31.98 38.14 Mg 9.41 9.17 7.97 9.26 9.35 6.83 9.40 7.18 9.29 wt% 10.73 10.03 11.87 10.14 10.67 11.38 H 0.01 0.01 0.59 0.01 0.03 0.07 5.51 0.01 0.06 0.18 4.59 0.01 0.01 0.30 0.65 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 244 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 8.80 8.04 7.45 8.45 wt% 11.34 11.90 13.13 11.62 11.25 12.67 10.59 11.83 Ti 0.05 0.05 0.04 0.06 0.05 0.05 0.03 0.06 0.05 0.03 0.05 0.04 wt% P 0.07 0.07 0.06 0.08 0.08 0.07 0.05 0.07 0.06 0.05 0.06 0.05 wt% Cr 0.22 0.25 0.19 0.26 0.26 0.22 0.17 0.23 0.21 0.17 0.21 0.16 wt% Ni 1.04 1.14 0.85 1.20 1.16 1.04 0.78 1.11 0.98 0.75 1.00 0.77 wt% Na 0.21 0.26 0.20 0.24 0.30 0.22 0.19 0.20 0.20 0.20 0.16 0.15 wt% years 6 Ca 10 1.27 1.29 0.95 1.47 1.24 1.25 0.87 1.45 1.18 0.79 1.38 0.94 wt% × t=3 Al 0.87 0.89 0.65 1.01 0.85 0.86 0.59 0.99 0.81 0.54 0.95 0.65 wt% Continued on next page Fe wt% 16.41 18.07 13.50 19.02 18.32 16.50 12.41 17.64 15.47 11.91 15.90 12.20 Table B.4 – continued from previous page S 1.75 2.50 1.96 2.03 3.13 1.87 1.86 1.47 1.69 2.16 1.06 1.24 wt% O wt% 30.25 32.33 42.98 24.86 36.35 31.60 46.11 23.78 33.26 51.67 23.86 38.74 Mg 8.98 9.78 7.28 9.82 9.00 6.68 9.75 8.45 6.36 8.86 6.68 wt% 10.40 H 0.18 0.29 0.82 0.05 0.63 0.21 1.01 0.03 0.23 1.82 0.03 0.32 wt% CJS1-4 CJS1-5 CJS1-6 CJS2-4 CJS2-5 CJS2-6 CJS2-7 CJS3-4 CJS3-5 CJS3-6 CJS4-4 CJS4-5 Simulation 245 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.55 7.97 wt% 12.42 11.30 11.51 12.59 12.76 12.22 12.86 13.65 10.21 13.10 13.21 12.78 10.17 Ti 0.06 0.05 0.05 0.06 0.06 0.05 0.03 0.06 0.06 0.04 0.03 0.06 0.06 0.05 0.04 wt% P 0.06 0.06 0.07 0.06 0.07 0.07 0.06 0.07 0.09 0.06 0.06 0.07 0.08 0.09 0.07 wt% Cr 0.21 0.20 0.24 0.21 0.23 0.25 0.19 0.23 0.30 0.20 0.20 0.24 0.26 0.29 0.22 wt% Ni 1.04 0.97 1.12 1.03 1.10 1.15 0.83 1.11 1.36 0.95 0.86 1.14 1.20 1.29 1.01 wt% Na 0.16 0.15 0.26 0.15 0.19 0.25 0.28 0.19 0.35 0.20 0.28 0.20 0.24 0.35 0.26 wt% Ca 1.46 1.31 1.24 1.49 1.47 1.34 0.76 1.48 1.45 1.13 0.81 1.50 1.48 1.35 1.08 wt% Al 1.00 0.90 0.85 1.02 1.01 0.92 0.52 1.02 1.00 0.78 0.56 1.03 1.02 0.93 0.74 wt% Fe wt% 16.47 15.34 17.65 16.40 17.45 18.26 13.04 17.67 21.50 15.05 13.61 18.08 19.00 20.47 15.99 Table B.4 – continued from previous page S 1.09 0.99 2.49 0.97 1.37 2.30 3.64 1.42 3.59 1.75 3.59 1.45 1.99 3.79 2.63 wt% O wt% 21.47 26.02 34.37 20.00 22.39 30.02 60.53 22.35 29.60 35.10 57.47 21.95 24.33 33.80 40.98 Mg 9.21 8.53 9.52 9.21 9.69 9.91 6.84 9.80 8.19 7.15 8.57 wt% 11.52 10.01 10.42 10.91 H 0.01 0.01 0.34 0.01 0.02 0.18 5.51 0.02 0.13 0.32 4.60 0.01 0.03 0.44 0.82 wt% EJS1-4 EJS1-5 EJS1-6 EJS2-4 EJS2-5 EJS2-6 EJS2-7 EJS3-4 EJS3-5 EJS3-6 EJS3-7 EJS4-4 EJS4-5 EJS4-6 EJS4-7 Simulation 246

APPENDIX C

MIDPLANE TEMPERATURE AND PRESSURE PROFILES

In Chapter 4, midplane temperature and pressure proﬁles were used to obtain the radial compositional variation within each of the systems studied. The proﬁles were obtained from the Hersant et al. (2001) model and were scaled with host star (and thus also disk) mass. The scaling was achieved by altering the mass accretion rate (M˙ ) of the host star by the following relationship:

˙ 3/2 3/2 M ∝ Mdisk ∝ Mstar (C.1)

Appendix C displays the radial pressure and temperature proﬁles obtained for each of the planetary systems examined. Midplane conditions are shown for seven different evolutionary stages of disk evolution (t = 2.5×105yr, 5×105yr, 1×106yr, 1.5×106yr, 2×106yr, 2.5×106yr and 3×106yr) in Figures C.1 - C.9. All proﬁles were obtained from the Hersant et al. (2001) models for the “nominal” conditions (M˙ = 5×10−6MJyr−1,

Rinital = 17 AU and α = 0.009). Proﬁles are shown in order of increasing HD number. 247

55Cnc

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.1: Midplane temperature and pressure proﬁle for 55Cnc for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 248

Gl777

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.2: Midplane temperature and pressure proﬁle for Gl777 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 249

HD4203

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.3: Midplane temperature and pressure proﬁle for HD4203 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 250

HD4208

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.4: Midplane temperature and pressure proﬁle for HD4208 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 251

HD19994

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.5: Midplane temperature and pressure proﬁle for HD19994 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 252

HD72659

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.6: Midplane temperature and pressure proﬁle for HD72659 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 253

HD108874

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.7: Midplane temperature and pressure proﬁle for HD108874 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 254

HD177830

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.8: Midplane temperature and pressure proﬁle for HD177830 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 255

HD142415

1e-1

1e-2

1e-3 0.25 Myr 0.5 Myr 1 Myr 1e-4 1.5 Myr 2 Myr

P (atm) 1e-5 2.5 Myr 3 Myr 1e-6

1e-7

1e-8 012345 Radius (AU) 3000

2500 0.25 Myr 0.5 Myr 2000 1 Myr 1.5 Myr 2 Myr 1500

T (K) 2.5 Myr 3 Myr 1000

500

0 012345 Radius (AU)

Figure C.9: Midplane temperature and pressure proﬁle for HD142415 for all seven sets of disk conditions simulated. Proﬁles were obtained from the “nominal” model of Hersant et al. (2001). 256

APPENDIX D

HSC CHEMISTRY GAS ABUNDANCES

The equilibrium calculations of HSC Chemistry determine the equilibrium composition of both solid and gaseous species present within the disk. The solid species compositions obtained for each system were previously shown in Figures 4.12 - 4.20. Appendix D contains the same plots for the gaseous species only. Figures D.1 - D.9 show the equilibrium gaseous composition for each of the nine extrasolar planetary systems studied in order of increasing C/O value. The abundances are normalized to the least abundant species present and were obtained from the input molar abundances listed in Table 4.6. Values are shown for a pressure of 10−4 bar. Although pressure was varied in the simulations of Chapter 4 from 10−2 bar to 10−9 bar in accordance with the midplane models of Hersant et al. (2001), these variations do not alter the general structure and composition shown in the following ﬁgures. 257

ce (mole)ce T vs Mg(g) T vs N2(g) T vs Na(g) 100 T vs NaOH(g) bundan T vs NH3(g) A T vs Ni(g) ed T vs O(g) T vs P(g) T vs PH(g)

Figure D.1: Schematic of the output obtained from HSC Chemistry for HD72659 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 258

T vs Al(g) T vs Al2O(g) T vs AlH(g) T vs Ca(g) HD177830 T vs CH4(g) T vs CN(g) T vs CO(g) 10000 T vs CO2(g) T vs Cr(g) T vs CS(g) T vs Fe(g) T vs H(g) T vs H2(g) T vs H2O(g) 1000 T vs H2S(g) T vs HCN(g) T vs He(g) T vs HS(g) T vs Mg(g)

ce (mole)ce T vs N2(g) T vs Na(g) T vs NaOH(g) 100

bundan T vs NH3(g)

A T vs Ni(g)

ed T vs O(g) T vs P(g) T vs PH(g) T vs PN(g)

Normaliz T vs PO(g) 10 T vs PS(g) T vs S(g) T vs S2(g) T vs Si(g) T vs SiH(g) T vs SiO(g) T vs SiS(g) 1 T vs SO(g) 200 400 600 800 1000 1200 1400 1600 1800 T vs SO2(g) T vs Ti(g) T (K) T vs TiO(g) T vs TiO2(g)

Figure D.2: Schematic of the output obtained from HSC Chemistry for HD177830 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 259

Al(g) Gl 777 Al2O(g) AlH(g) Ca(g) CH4(g) 10000 CN(g) CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) 1000 H2O(g) H2S(g) HCN(g) He(g) HS(g)

ce (mole)ce Mg(g) N2(g) Na(g) 100 NaOH(g) bundan NH3(g) A Ni(g)

ed O(g) P(g) PH(g) PN(g)

Normaliz PO(g) 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.3: Schematic of the output obtained from HSC Chemistry for Gl777 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 260

T vs Al(g) T vs Al2O(g) T vs AlH(g) HD4208 T vs Ca(g) T vs CH4(g) T vs CN(g) T vs CO(g) 10000 T vs CO2(g) T vs Cr(g) T vs CS(g) T vs Fe(g) T vs H(g) T vs H2(g) T vs H2O(g) 1000 T vs H2S(g) T vs HCN(g) T vs He(g) T vs HS(g)

ce (mole)ce T vs Mg(g) T vs N2(g) T vs Na(g) 100 T vs NaOH(g) bundan T vs NH3(g) A T vs Ni(g) ed T vs O(g) T vs P(g) T vs PH(g)

Figure D.4: Schematic of the output obtained from HSC Chemistry for HD4208 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 261

Al(g) 55Cnc Al2O(g) AlH(g) Ca(g) CH4(g) 10000 CN(g) CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) 1000 H2O(g) H2S(g) HCN(g) He(g) HS(g)

ce (mole)ce Mg(g) N2(g) Na(g) 100 NaOH(g) bundan NH3(g) A Ni(g)

ed O(g) P(g) PH(g) PN(g)

Normaliz PO(g) 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.5: Schematic of the output obtained from HSC Chemistry for 55Cnc at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 262

Al(g) HD142415 Al2O(g) AlH(g) Ca(g) CH4(g) 10000 CN(g) CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) 1000 H2O(g) H2S(g) HCN(g) He(g) HS(g)

ce (mole)ce Mg(g) N2(g) Na(g) 100 NaOH(g) bundan NH3(g) A Ni(g)

ed O(g) P(g) PH(g) PN(g)

Normaliz PO(g) 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.6: Schematic of the output obtained from HSC Chemistry for HD142415 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 263

Al(g) HD19994 Al2O(g) AlH(g) Ca(g) CH4(g) 10000 CN(g) CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) 1000 H2O(g) H2S(g) HCN(g) He(g) HS(g)

ce (mole)ce Mg(g) N2(g) Na(g) 100 NaOH(g) bundan NH3(g) A Ni(g)

ed O(g) P(g) PH(g) PN(g)

Normaliz PO(g) 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.7: Schematic of the output obtained from HSC Chemistry for HD19994 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 264

Al(g) HD108874 Al2O(g) AlH(g) Ca(g) CH4(g) 10000 CN(g) CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) 1000 H2O(g) H2S(g) HCN(g) He(g) HS(g)

ce (mole)ce Mg(g) N2(g) Na(g) 100 NaOH(g) bundan NH3(g) A Ni(g)

ed O(g) P(g) PH(g) PN(g)

Normaliz PO(g) 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.8: Schematic of the output obtained from HSC Chemistry for HD108874 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 265

Al(g) Al2O(g) HD4203 AlH(g) Ca(g) CH4(g) CN(g) 10000 CO(g) CO2(g) Cr(g) CS(g) Fe(g) H(g) H2(g) H2O(g) 1000 H2S(g) HCN(g) He(g) HS(g) Mg(g) ce (mole)ce N2(g) Na(g) 100 NaOH(g)

bundan NH3(g)

A Ni(g)

ed O(g) P(g) PH(g) PN(g) PO(g) Normaliz 10 PS(g) S(g) S2(g) Si(g) SiH(g) SiO(g) SiS(g) SO(g) 1 SO2(g) 200 400 600 800 1000 1200 1400 1600 1800 2000 2200 Ti(g) TiO(g) T (K) TiO2(g)

Figure D.9: Schematic of the output obtained from HSC Chemistry for HD4203 at a pressure of 10−4 bar. Only gaseous species present within the system are shown. All abundances are normalized to the least abundant species present. Input elemental abundances are shown in Table 4.6. 266

APPENDIX E

EXTRASOLAR TERRESTRIAL PLANET ABUDNANCES

In Chapter 4, predicted bulk elemental abundances were determined for the terrestrial planets simulated to form in nine different extrasolar planetary systems. Abundances were determined for midplane conditions from the models of Hersant et al. (2001) at seven different disk evolutionary times (t = 2.5×105yr, 5×105yr, 1×106yr, 1.5×106yr, 2×106yr, 2.5×106yr and 3×106yr). Appendix E contains the graphical and numerical results of these simulations. Planetary systems are presented in order of increasing C/O value. Figures E.1 - E.32 provide a schematic depiction of the bulk planetary compositions (in wt%). The composition of planets produced in each of the four simulations completed for each system are shown for each set of disk conditions examined. Tables E.1 through E.9 list the bulk elemental abundances for the simulated terrestrial planets produced in each system. Abundances are listed as wt% of the ﬁnal predicted planet for all seven sets of disk conditions examined. 267

Final Composition - Gl777 (0.25Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU) Final Composition - Gl777 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Figure E.1: Schematic of the bulk elemental planetary composition for Gl777. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 268

Final Composition - Gl777 (1.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU) Final Composition - Gl777 (1.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Figure E.2: Schematic of the bulk elemental planetary composition for Gl777. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 269

Final Composition - Gl777 (2.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Final Composition - Gl777 (2.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Figure E.3: Schematic of the bulk elemental planetary composition for Gl777. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 270

Final Composition - Gl777 (3.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 Semimajor Axis (AU)

Figure E.4: Schematic of the bulk elemental planetary composition for Gl777. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 271

Final Composition - HD4208 (0.25Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU) Final Composition - HD4208 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca

Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Figure E.5: Schematic of the bulk elemental planetary composition for HD4208. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 272

Final Composition - HD4208 (1.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU) Final Composition - HD4208 (1.5Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Figure E.6: Schematic of the bulk elemental planetary composition for HD4208. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 273

Final Composition - HD4208 (2.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Final Composition - HD4208 (2.5Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Figure E.7: Schematic of the bulk elemental planetary composition for HD4208. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 274

Final Composition - HD4208 (3.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4 Semimajor Axis (AU)

Figure E.8: Schematic of the bulk elemental planetary composition for HD4208. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 275

Final Composition - HD72659 (0.25Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU) Final Composition - HD72659 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Figure E.9: Schematic of the bulk elemental planetary composition for HD72659. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 276

Final Composition - HD72659 (1.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Final Composition - HD72659 (1.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Figure E.10: Schematic of the bulk elemental planetary composition for HD72659. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 277

Final Composition - HD72659 (2.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU) Final Composition - HD72659 (2.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Figure E.11: Schematic of the bulk elemental planetary composition for HD72659. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 278

Final Composition - HD72659 (3.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 0.8 1.0 1.2 Semimajor Axis (AU)

Figure E.12: Schematic of the bulk elemental planetary composition for HD72659. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 279

Final Composition - HD177830 (0.25Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU) Final Composition - HD177830 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Al Ca Sim. 4 Other

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Figure E.13: Schematic of the bulk elemental planetary composition for HD177830. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 280

Final Composition - HD177830 (1.0Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Final Composition - HD177830 (1.5Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Figure E.14: Schematic of the bulk elemental planetary composition for HD177830. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 281

Final Composition - HD177830 (2.0Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU) Final Composition - HD177830 (2.5Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Figure E.15: Schematic of the bulk elemental planetary composition for HD177830. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 282

Final Composition - HD177830 (3.0Myr)

O Sim. 1 Fe Mg Si Sim. 2 C S

Sim. 3 Al Ca Other Sim. 4

0.0 0.2 0.4 0.6 Semimajor Axis (AU)

Figure E.16: Schematic of the bulk elemental planetary composition for HD177830. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 283

Final Composition - 55 Cnc (0.25 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU) Final Composition - 55 Cnc (0.5 Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Figure E.17: Schematic of the bulk elemental planetary composition for HD55Cnc. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 284

Final Composition - 55 Cnc (1.0 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Final Composition - 55 Cnc (1.5 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Figure E.18: Schematic of the bulk elemental planetary composition for HD55Cnc. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 285

Final Composition - 55nc C (2.0 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis U)(A

Final Composition - 55 Cnc (2.5 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Figure E.19: Schematic of the bulk elemental planetary composition for HD55Cnc. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 286

Final Composition - 55 Cnc (3.0 Myr)

O Sim. 1 Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

01234 Semimajor Axis (AU)

Figure E.20: Schematic of the bulk elemental planetary composition for HD55Cnc. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 287

Final Composition - HD142415 (0.25Myr)

O Sim. 1 Fe M g

Sim. 2 Si C S Sim. 3 Al Ca Ti Sim. 4 O ther

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU) Final Composition - HD142415 (0.50Myr)

O Sim. 1 Fe M g

Sim. 2 Si C S Sim. 3 Al Ca Ti Sim. 4 O ther

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.21: Schematic of the bulk elemental planetary composition for HD142415. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 288

Final Composition - HD142415 (1.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU) Final Composition - HD142415 (1.5Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.22: Schematic of the bulk elemental planetary composition for HD142415. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 289

Final Composition - HD142415 (2.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Final Composition - HD142415 (2.5Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.23: Schematic of the bulk elemental planetary composition for HD142415. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 290

Final Composition - HD142415 (3.0Myr)

O Sim. 1 Fe Mg

Sim. 2 Si C S Sim. 3 Al Ca Other Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.24: Schematic of the bulk elemental planetary composition for HD142415. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 291

Final Composition - HD108874 (0.25Myr)

Sim. 1 O Fe

Sim. 2 Mg Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU) Final Composition - HD108874 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C

Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.25: Schematic of the bulk elemental planetary composition for HD108874. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 292

Final Composition - HD108874 (1.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Final Composition - HD108874 (1.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.26: Schematic of the bulk elemental planetary composition for HD108874. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 293

Final Composition - HD108874 (2.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU) Final Composition - HD108874 (2.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.27: Schematic of the bulk elemental planetary composition for HD108874. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 294

Final Composition - HD108874 (3.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C Sim. 3 S Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 0.6 Semimajor Axis (AU)

Figure E.28: Schematic of the bulk elemental planetary composition for HD108874. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 295

Final Composition - HD4203 (0.25Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU) Final Composition - HD4203 (0.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Figure E.29: Schematic of the bulk elemental planetary composition for HD4203. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2.5×105 years. Bottom Panel: Compositions based on disk conditions at t = 5×105 years. 296

Final Composition - HD4203 (1.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU) Final Composition - HD4203 (1.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Figure E.30: Schematic of the bulk elemental planetary composition for HD4203. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 1×106 years. Bottom Panel: Compositions based on disk conditions at t = 1.5×106 years. 297

Final Composition - HD4203 (2.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Final Composition - HD4203 (2.5Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Figure E.31: Schematic of the bulk elemental planetary composition for HD4203. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Top Panel: Compositions based on disk conditions at t = 2×106 years. Bottom Panel: Compositions based on disk conditions at t = 2.5×106 years. 298

Final Composition - HD4203 (3.0Myr)

Sim. 1 O Fe Mg Sim. 2 Si C S Sim. 3 Other

Sim. 4

0.0 0.1 0.2 0.3 0.4 0.5 Semimajor Axis (AU)

Figure E.32: Schematic of the bulk elemental planetary composition for HD4203. All values are wt% of the ﬁnal simulated planet. Values are shown for the terrestrial planets produced in each of the four simulations run for the system. Size of bodies is not to scale. Compositions based on disk conditions at t = 3×106 years. 299 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.43 12.97 12.43 11.25 14.31 12.77 Ti 0.26 0.44 0.56 0.67 0.36 0.50 wt% P 0.09 0.00 0.04 0.03 0.05 0.04 wt% Cr 0.28 0.19 0.15 0.15 0.23 0.19 wt% Ni 1.48 1.28 1.00 0.77 1.49 1.06 wt% Na 0.21 0.00 0.11 0.00 0.13 0.09 wt% years 5 10 Ca 3.18 5.32 7.72 8.15 4.78 6.33 wt% × t=2.5 Al 6.31 6.36 wt% 11.26 12.16 16.13 11.21 Continued on next page Fe wt% 22.86 18.79 14.19 11.95 21.58 16.14 S 0.60 0.00 0.00 0.00 0.00 0.05 wt% O wt% 34.37 36.27 37.78 39.01 34.77 37.14 Mg wt% 15.93 13.46 13.85 11.88 15.92 14.50 Table E.1: Predicted bulk elemental planetary abundances for Gl777. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 300 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.89 16.46 16.38 16.68 15.92 16.09 Ti 0.08 0.08 0.08 0.08 0.08 0.08 wt% P 0.13 0.11 0.12 0.09 0.15 0.11 wt% Cr 0.33 0.33 0.33 0.33 0.34 0.34 wt% Ni 1.75 1.77 1.73 1.74 1.79 1.78 wt% Na 0.52 0.41 0.30 0.27 0.46 0.35 wt% years 5 Ca 10 1.04 1.06 1.03 1.04 1.06 1.06 wt% × t=5 Al 1.36 1.39 1.41 1.36 1.45 1.44 wt% Continued on next page Fe wt% 26.94 27.33 26.67 26.82 27.55 27.41 Table E.1 – continued from previous page S 1.89 0.00 1.13 0.23 1.32 0.74 wt% O wt% 32.45 32.87 32.71 33.06 32.13 32.70 Mg wt% 17.63 18.20 18.12 18.31 17.77 17.90 H 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 301 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.05 15.60 15.45 15.57 15.29 15.42 Ti 0.08 0.08 0.08 0.08 0.08 0.08 wt% P 0.15 0.15 0.15 0.15 0.15 0.15 wt% Cr 0.32 0.33 0.33 0.33 0.33 0.33 wt% Ni 1.69 1.75 1.73 1.75 1.72 1.73 wt% Na 0.59 0.61 0.61 0.61 0.60 0.60 wt% years 6 Ca 10 1.00 1.04 1.03 1.04 1.02 1.03 wt% × t=1 Al 1.32 1.37 1.35 1.36 1.34 1.35 wt% Continued on next page Fe wt% 26.07 27.01 26.77 26.97 26.48 26.72 Table E.1 – continued from previous page S 4.31 3.12 2.96 2.63 4.03 3.22 wt% O wt% 32.63 31.54 32.30 32.13 31.93 32.16 Mg wt% 16.79 17.39 17.23 17.36 17.05 17.20 H 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 302 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.78 15.11 15.01 15.21 14.84 15.00 Ti 0.07 0.08 0.08 0.08 0.07 0.08 wt% P 0.15 0.15 0.15 0.15 0.15 0.15 wt% Cr 0.31 0.32 0.32 0.32 0.32 0.32 wt% Ni 1.66 1.69 1.68 1.71 1.67 1.68 wt% Na 0.58 0.59 0.59 0.60 0.58 0.59 wt% years 6 10 Ca 0.99 1.01 1.00 1.01 0.99 1.00 wt% × t=1.5 Al 1.30 1.32 1.32 1.33 1.30 1.31 wt% Continued on next page Fe wt% 25.60 26.16 26.00 26.33 25.70 25.98 Table E.1 – continued from previous page S 4.86 4.47 4.87 4.39 5.08 4.81 wt% O wt% 33.22 32.26 32.25 31.92 32.76 32.36 Mg wt% 16.49 16.85 16.74 16.96 16.55 16.73 H 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 303 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.37 14.88 14.60 14.87 14.64 13.71 Ti 0.07 0.07 0.07 0.07 0.07 0.07 wt% P 0.14 0.15 0.14 0.15 0.14 0.13 wt% Cr 0.31 0.32 0.31 0.32 0.31 0.29 wt% Ni 1.61 1.67 1.64 1.67 1.64 1.54 wt% Na 0.56 0.58 0.57 0.58 0.57 0.54 wt% years 6 Ca 10 0.96 0.99 0.97 0.99 0.98 0.91 wt% × t=2 Al 1.26 1.30 1.28 1.30 1.28 1.20 wt% Continued on next page Fe wt% 24.88 25.77 25.27 25.75 25.36 23.73 Table E.1 – continued from previous page S 5.00 5.07 5.10 5.06 5.16 4.77 wt% O wt% 34.67 32.59 33.66 32.67 33.50 31.58 Mg wt% 16.02 16.60 16.28 16.58 16.33 21.42 H 0.15 0.00 0.10 0.00 0.00 0.09 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 304 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.13 14.61 14.47 14.65 14.44 14.44 Ti 0.07 0.07 0.07 0.07 0.07 0.07 wt% P 0.14 0.14 0.14 0.14 0.14 0.14 wt% Cr 0.30 0.31 0.31 0.31 0.31 0.31 wt% Ni 1.58 1.64 1.62 1.64 1.62 1.62 wt% Na 0.55 0.57 0.57 0.57 0.57 0.56 wt% years 6 10 Ca 0.94 0.97 0.97 0.98 0.96 0.96 wt% × t=2.5 Al 1.24 1.28 1.27 1.28 1.27 1.26 wt% Continued on next page Fe wt% 24.46 25.29 25.05 25.36 25.01 24.99 Table E.1 – continued from previous page S 4.98 5.13 5.10 5.14 5.10 5.08 wt% O wt% 35.59 33.70 34.16 33.51 34.33 34.35 Mg wt% 15.76 16.29 16.14 16.34 16.11 16.10 H 0.26 0.00 0.12 0.00 0.07 0.11 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 305 C 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.05 14.56 14.27 13.53 14.16 13.75 Ti 0.07 0.07 0.07 0.07 0.07 0.07 wt% P 0.13 0.14 0.14 0.13 0.14 0.14 wt% Cr 0.28 0.31 0.30 0.29 0.30 0.29 wt% Ni 1.46 1.63 1.60 1.52 1.59 1.54 wt% Na 0.51 0.57 0.56 0.53 0.55 0.54 wt% years 6 Ca 10 0.87 0.97 0.95 0.90 0.94 0.92 wt% × t=3 Al 1.14 1.28 1.25 1.19 1.24 1.20 wt% Fe wt% 22.59 25.21 24.71 23.42 24.52 23.81 Table E.1 – continued from previous page S 4.61 5.14 5.04 4.77 5.00 4.85 wt% O wt% 39.63 33.87 34.99 31.22 35.45 36.92 Mg wt% 14.56 16.24 15.92 22.44 15.80 15.34 H 1.10 0.00 0.20 0.00 0.23 0.62 wt% 1-4 2-4 2-5 3-4 3-5 4-4 Simulation 306 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 3.02 1.59 4.93 2.81 7.46 9.12 5.62 wt% 13.24 13.14 17.46 15.95 Ti 1.50 0.66 1.41 1.53 0.91 1.89 1.13 0.93 1.58 0.09 0.08 wt% P 0.00 0.05 0.00 0.00 0.00 0.00 0.00 0.01 0.00 0.18 0.17 wt% Cr 0.00 0.18 0.00 0.00 0.09 0.00 0.04 0.08 0.00 0.39 0.36 wt% Ni 0.00 1.13 0.00 0.00 0.71 0.00 0.37 0.80 0.00 1.83 1.67 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.58 wt% years 5 10 Ca 9.86 8.86 5.71 9.59 1.31 1.19 wt% 15.35 11.37 12.73 13.56 17.35 × t=2.5 Al 1.64 1.50 wt% 38.27 12.10 44.66 30.94 16.94 36.94 27.11 17.59 29.81 Continued on next page Fe 0.00 0.00 0.00 0.00 4.77 0.00 wt% 16.27 10.04 10.67 29.54 26.97 S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O wt% 45.15 36.99 46.64 43.44 39.99 44.85 42.05 38.10 43.10 32.25 32.92 Mg 2.20 0.00 3.82 6.82 3.92 4.33 9.16 2.54 wt% 10.52 15.32 18.62 Table E.2: Predicted bulk elemental planetary abundances for HD4208. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 307 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 7.89 8.43 9.82 6.63 wt% 16.19 13.82 13.50 15.01 13.17 15.94 15.52 Ti 1.28 0.08 1.25 1.01 0.41 1.50 0.60 0.24 0.72 0.08 0.08 wt% P 0.00 0.12 0.00 0.00 0.11 0.00 0.06 0.10 0.00 0.17 0.16 wt% Cr 0.05 0.36 0.08 0.12 0.26 0.00 0.21 0.31 0.17 0.35 0.35 wt% Ni 0.39 1.70 0.38 0.59 1.44 0.00 1.10 1.51 0.86 1.67 1.63 wt% Na 0.00 0.30 0.00 0.00 0.20 0.00 0.07 0.35 0.00 0.61 0.59 wt% years 5 Ca 10 1.21 5.63 7.61 3.27 1.19 1.16 wt% 16.20 14.98 13.26 20.17 10.01 × t=5 Al 1.52 7.46 4.36 1.50 1.46 wt% 23.55 23.05 18.53 27.51 11.12 13.16 Continued on next page Fe 5.35 6.15 9.48 0.00 wt% 27.39 21.96 17.09 24.23 13.78 26.96 26.27 Table E.2 – continued from previous page S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 2.60 wt% O wt% 40.74 32.79 40.75 39.21 34.59 42.57 36.82 33.91 37.96 32.92 32.06 Mg 4.55 4.92 7.98 1.61 wt% 18.32 14.12 11.82 16.72 10.17 18.61 18.12 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 308 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.95 15.58 16.15 15.62 15.98 17.18 15.45 15.62 15.85 15.25 14.74 Ti 0.44 0.08 0.61 0.22 0.08 0.33 0.23 0.08 0.08 0.08 0.08 wt% P 0.06 0.16 0.04 0.08 0.17 0.03 0.13 0.15 0.13 0.16 0.16 wt% Cr 0.23 0.35 0.16 0.31 0.36 0.26 0.30 0.35 0.34 0.34 0.33 wt% Ni 1.25 1.63 0.87 1.58 1.67 1.29 1.50 1.64 1.77 1.60 1.54 wt% Na 0.11 0.60 0.14 0.21 0.46 0.00 0.37 0.53 0.31 0.58 0.56 wt% years 6 Ca 10 6.27 1.17 8.69 3.21 1.20 4.82 3.35 1.17 1.30 1.14 1.10 wt% × t=1 Al 8.04 1.46 4.03 1.50 6.04 4.24 1.47 1.63 1.43 1.38 wt% 11.30 Continued on next page Fe wt% 19.15 26.35 13.74 24.40 27.05 20.69 23.60 26.43 27.71 25.76 24.93 Table E.2 – continued from previous page S 0.00 2.19 0.00 0.00 0.86 0.00 0.50 2.12 0.00 4.34 5.07 wt% O wt% 36.33 32.25 38.81 33.89 32.55 35.95 34.04 32.20 32.54 31.51 32.89 Mg 9.49 wt% 12.18 18.19 16.46 18.12 13.41 16.29 18.24 18.33 17.81 17.22 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 309 C 0.00 0.00 3.05 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 8.52 wt% 16.16 15.08 16.06 15.27 16.07 15.73 15.15 15.89 14.77 14.33 Ti 0.08 0.08 0.05 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.07 wt% P 0.12 0.16 0.82 0.15 0.16 0.12 0.15 0.16 0.16 0.16 0.15 wt% Cr 0.36 0.34 0.77 0.36 0.34 0.36 0.35 0.34 0.35 0.33 0.32 wt% Ni 1.69 1.58 0.00 1.68 1.60 1.68 1.65 1.59 1.66 1.55 1.50 wt% Na 0.23 0.58 0.01 0.37 0.58 0.19 0.46 0.57 0.47 0.57 0.55 wt% years 6 10 Ca 1.21 1.13 2.83 1.20 1.14 1.20 1.18 1.13 1.19 1.11 1.07 wt% × t=1.5 Al 1.52 1.41 0.96 1.51 1.43 1.51 1.48 1.42 1.49 1.39 1.35 wt% Continued on next page Fe wt% 27.32 25.50 64.28 27.18 25.85 27.17 26.61 25.65 26.89 25.01 24.27 Table E.2 – continued from previous page S 0.31 4.05 0.00 0.37 3.22 0.00 1.88 3.50 0.56 5.08 4.94 wt% O wt% 32.69 32.50 18.71 32.73 32.49 32.86 32.49 32.71 32.72 32.73 34.73 Mg 0.00 wt% 18.30 17.61 18.31 17.83 18.76 17.95 17.69 18.54 17.24 16.72 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 310 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.98 14.76 15.90 15.78 15.18 15.95 15.41 14.90 15.53 14.75 14.32 Ti 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.07 wt% P 0.17 0.16 0.17 0.17 0.16 0.17 0.16 0.16 0.16 0.16 0.15 wt% Cr 0.36 0.33 0.35 0.35 0.34 0.36 0.34 0.33 0.35 0.33 0.32 wt% Ni 1.67 1.55 1.66 1.65 1.59 1.67 1.61 1.56 1.63 1.54 1.50 wt% Na 0.45 0.57 0.35 0.55 0.58 0.57 0.50 0.57 0.59 0.56 0.55 wt% years 6 Ca 10 1.20 1.11 1.19 1.18 1.14 1.19 1.15 1.11 1.16 1.10 1.07 wt% × t=2 Al 1.50 1.39 1.49 1.48 1.42 1.50 1.45 1.40 1.46 1.38 1.35 wt% Continued on next page Fe wt% 27.03 24.98 26.89 26.70 25.68 26.98 26.07 25.20 26.28 24.92 24.26 Table E.2 – continued from previous page S 0.67 4.93 1.10 1.59 3.63 0.00 2.88 4.26 2.11 5.07 4.94 wt% O wt% 32.79 32.94 32.85 32.36 32.49 32.92 32.54 33.04 32.52 32.89 34.76 Mg wt% 18.11 17.23 17.97 18.09 17.72 18.62 17.81 17.39 18.13 17.22 16.71 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 311 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.71 14.58 15.68 15.53 14.88 15.90 15.13 14.76 15.32 14.33 14.32 Ti 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.07 0.07 wt% P 0.17 0.15 0.17 0.16 0.16 0.17 0.16 0.16 0.16 0.15 0.15 wt% Cr 0.35 0.33 0.35 0.35 0.33 0.35 0.34 0.33 0.34 0.32 0.32 wt% Ni 1.65 1.53 1.64 1.63 1.56 1.67 1.58 1.55 1.60 1.50 1.50 wt% Na 0.59 0.56 0.57 0.59 0.57 0.61 0.57 0.57 0.59 0.55 0.55 wt% years 6 10 Ca 1.18 1.09 1.17 1.16 1.11 1.19 1.13 1.11 1.15 1.07 1.07 wt% × t=2.5 Al 1.47 1.37 1.47 1.46 1.40 1.49 1.42 1.39 1.44 1.35 1.35 wt% Continued on next page Fe wt% 26.59 24.69 26.54 26.28 25.18 26.92 25.59 24.97 25.92 24.27 24.26 Table E.2 – continued from previous page S 0.99 5.01 1.14 2.15 4.23 0.22 3.61 4.47 3.06 4.94 4.94 wt% O wt% 32.88 33.59 32.87 32.49 33.13 32.84 32.72 33.42 32.46 34.73 34.76 Mg wt% 18.34 17.03 18.31 18.13 17.38 18.56 17.66 17.23 17.89 16.72 16.71 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 312 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 15.62 14.50 15.67 15.32 14.66 15.75 14.97 14.61 15.20 14.32 13.50 Ti 0.08 0.07 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.07 0.07 wt% P 0.16 0.15 0.17 0.16 0.15 0.17 0.16 0.15 0.16 0.15 0.14 wt% Cr 0.35 0.32 0.35 0.34 0.33 0.35 0.33 0.33 0.34 0.32 0.30 wt% Ni 1.64 1.52 1.64 1.61 1.54 1.65 1.57 1.53 1.59 1.50 1.42 wt% Na 0.60 0.56 0.60 0.59 0.56 0.60 0.57 0.56 0.58 0.55 0.52 wt% years 6 Ca 10 1.17 1.09 1.17 1.15 1.10 1.18 1.12 1.10 1.14 1.07 1.01 wt% × t=3 Al 1.47 1.36 1.47 1.44 1.38 1.48 1.40 1.37 1.43 1.34 1.27 wt% Fe wt% 26.42 24.54 26.53 25.92 24.79 26.64 25.33 24.76 25.72 24.26 22.86 Table E.2 – continued from previous page S 1.58 4.99 1.15 3.06 4.78 1.18 4.00 4.74 3.79 4.94 4.65 wt% O wt% 32.69 33.97 32.87 32.45 33.53 32.53 32.99 33.71 32.24 34.76 37.93 Mg wt% 18.23 16.92 18.30 17.89 17.11 18.39 17.48 17.05 17.73 16.71 15.76 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.56 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 313 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 5.76 6.54 7.71 7.63 6.88 wt% 16.01 11.70 15.78 13.79 14.94 17.08 Ti 2.06 0.21 1.90 1.20 0.19 1.63 0.79 0.44 1.64 2.11 0.09 wt% P 0.00 0.09 0.00 0.01 0.12 0.00 0.00 0.06 0.00 0.00 0.09 wt% Cr 0.00 0.32 0.00 0.06 0.34 0.00 0.11 0.28 0.00 0.00 0.32 wt% Ni 0.00 1.46 0.00 0.28 1.58 0.00 0.59 1.31 0.00 0.00 1.68 wt% Na 0.00 0.30 0.00 0.00 0.39 0.00 0.00 0.11 0.00 0.00 0.33 wt% years 5 10 Ca 3.20 2.79 6.66 1.30 wt% 18.18 20.24 14.93 23.37 11.70 23.17 15.24 × t=2.5 Al 3.12 2.72 6.50 1.27 wt% 30.25 27.73 17.41 23.79 11.44 24.01 31.05 Continued on next page Fe 0.00 0.00 4.62 0.00 9.43 0.00 0.00 wt% 24.07 26.11 21.62 26.81 S 0.00 0.00 0.00 0.00 0.67 0.00 0.00 0.00 0.00 0.00 0.00 wt% O wt% 43.30 34.11 42.70 41.35 32.93 41.76 39.47 34.76 41.81 44.23 33.25 Mg 0.45 0.89 8.43 1.75 1.73 0.49 wt% 17.10 16.39 12.69 13.31 17.78 Table E.3: Predicted bulk elemental planetary abundances for HD72659. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 314 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.64 9.53 9.81 wt% 16.51 10.09 15.86 16.21 17.40 16.77 12.95 16.46 Ti 1.46 0.08 1.34 0.41 0.08 1.38 0.09 0.08 1.08 1.18 0.08 wt% P 0.00 0.13 0.00 0.08 0.13 0.02 0.10 0.10 0.00 0.00 0.14 wt% Cr 0.00 0.36 0.00 0.26 0.35 0.06 0.33 0.36 0.05 0.11 0.36 wt% Ni 0.00 1.65 0.00 1.30 1.62 0.26 1.59 1.68 0.25 0.51 1.65 wt% Na 0.00 0.57 0.00 0.16 0.59 0.00 0.28 0.49 0.00 0.00 0.64 wt% years 5 Ca 10 1.23 4.19 1.21 1.37 1.25 1.23 wt% 21.67 19.83 15.62 16.09 17.16 × t=5 Al 1.20 6.04 1.18 1.33 1.22 1.20 wt% 21.27 19.47 19.95 15.70 17.05 Continued on next page Fe 0.00 0.00 4.24 4.05 8.38 wt% 27.27 21.02 26.82 25.97 27.69 27.25 Table E.3 – continued from previous page S 0.00 1.33 0.00 0.00 2.14 0.00 0.00 0.17 0.00 0.00 1.57 wt% O wt% 42.11 32.50 42.27 35.65 32.78 41.18 33.73 32.72 41.41 38.83 32.29 Mg 3.85 7.00 7.77 8.42 6.97 wt% 17.17 15.02 16.87 17.83 17.45 17.13 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 315 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 17.19 15.83 17.72 16.66 15.58 16.48 16.56 16.13 17.07 16.38 15.79 Ti 0.32 0.08 0.10 0.08 0.08 0.09 0.08 0.08 0.09 0.33 0.08 wt% P 0.03 0.13 0.04 0.13 0.13 0.09 0.14 0.13 0.08 0.04 0.13 wt% Cr 0.20 0.34 0.30 0.35 0.34 0.37 0.36 0.35 0.35 0.21 0.34 wt% Ni 1.18 1.59 1.57 1.67 1.56 1.71 1.66 1.61 1.70 1.40 1.58 wt% Na 0.00 0.62 0.04 0.53 0.61 0.13 0.64 0.63 0.17 0.19 0.62 wt% years 6 Ca 10 4.81 1.18 1.53 1.25 1.17 1.28 1.24 1.21 1.28 4.96 1.18 wt% × t=1 Al 4.68 1.15 1.49 1.22 1.14 1.25 1.21 1.18 1.25 4.83 1.15 wt% Continued on next page Fe wt% 18.29 26.20 25.05 27.39 25.78 28.26 27.40 26.66 27.70 19.82 26.12 Table E.3 – continued from previous page S 0.00 3.28 0.00 0.67 3.58 0.00 1.03 2.89 0.00 0.00 3.40 wt% O wt% 36.82 33.12 34.19 32.71 33.83 32.53 32.47 32.34 32.90 35.88 33.18 Mg wt% 16.49 16.47 17.97 17.34 16.21 17.80 17.22 16.79 17.41 15.96 16.42 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 316 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.93 15.40 16.81 16.31 15.08 16.68 16.20 15.62 16.74 16.88 15.45 Ti 0.09 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 0.08 wt% P 0.12 0.13 0.12 0.13 0.12 0.14 0.13 0.13 0.14 0.10 0.13 wt% Cr 0.37 0.33 0.37 0.35 0.33 0.36 0.35 0.34 0.36 0.37 0.34 wt% Ni 1.70 1.54 1.68 1.63 1.51 1.67 1.62 1.56 1.68 1.69 1.55 wt% Na 0.22 0.60 0.35 0.61 0.59 0.56 0.64 0.61 0.50 0.24 0.61 wt% years 6 10 Ca 1.27 1.15 1.26 1.22 1.13 1.25 1.21 1.17 1.25 1.26 1.16 wt% × t=1.5 Al 1.24 1.12 1.23 1.19 1.10 1.22 1.18 1.14 1.22 1.23 1.13 wt% Continued on next page Fe wt% 28.02 25.48 27.82 26.97 24.94 27.59 26.79 25.85 27.71 27.92 25.56 Table E.3 – continued from previous page S 0.00 3.64 0.00 1.99 3.78 0.44 2.65 3.20 0.15 0.29 3.93 wt% O wt% 32.77 34.50 32.80 32.52 35.52 32.66 32.29 34.07 32.76 32.65 34.00 Mg wt% 17.29 16.02 17.49 16.98 15.68 17.35 16.85 16.24 17.41 17.28 16.07 H 0.00 0.00 0.00 0.00 0.14 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 317 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.74 15.08 16.72 16.01 13.45 16.59 15.82 15.39 16.61 16.60 15.00 Ti 0.08 0.08 0.08 0.08 0.07 0.08 0.08 0.08 0.08 0.08 0.08 wt% P 0.14 0.12 0.14 0.13 0.11 0.14 0.13 0.13 0.14 0.14 0.12 wt% Cr 0.36 0.33 0.36 0.35 0.29 0.36 0.34 0.33 0.36 0.36 0.33 wt% Ni 1.68 1.51 1.67 1.60 1.35 1.66 1.58 1.54 1.66 1.66 1.50 wt% Na 0.60 0.59 0.59 0.63 0.53 0.65 0.62 0.60 0.65 0.49 0.59 wt% years 6 Ca 10 1.25 1.13 1.25 1.20 1.01 1.24 1.18 1.15 1.24 1.24 1.12 wt% × t=2 Al 1.22 1.10 1.22 1.17 0.98 1.21 1.15 1.12 1.21 1.21 1.09 wt% Continued on next page Fe wt% 27.72 24.95 27.67 26.49 22.25 27.44 26.16 25.47 27.48 27.48 24.81 Table E.3 – continued from previous page S 0.00 3.80 0.14 2.92 3.43 0.77 3.59 3.78 0.70 0.98 3.83 wt% O wt% 32.79 35.50 32.76 32.78 41.28 32.58 32.88 34.40 32.58 32.49 35.77 Mg wt% 17.41 15.69 17.40 16.65 14.00 17.27 16.46 16.00 17.28 17.27 15.61 H 0.00 0.13 0.00 0.00 1.25 0.00 0.00 0.00 0.00 0.00 0.14 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 318 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.68 14.83 16.63 15.78 13.25 16.36 15.55 15.30 16.48 16.52 15.32 Ti 0.08 0.07 0.08 0.08 0.07 0.08 0.08 0.08 0.08 0.08 0.08 wt% P 0.14 0.12 0.14 0.13 0.11 0.13 0.13 0.13 0.14 0.14 0.13 wt% Cr 0.36 0.32 0.36 0.34 0.29 0.36 0.34 0.33 0.36 0.36 0.33 wt% Ni 1.67 1.49 1.67 1.58 1.33 1.64 1.56 1.53 1.65 1.65 1.53 wt% Na 0.65 0.58 0.65 0.62 0.52 0.64 0.61 0.60 0.65 0.65 0.60 wt% years 6 10 Ca 1.25 1.11 1.24 1.18 0.99 1.22 1.16 1.14 1.23 1.23 1.14 wt% × t=2.5 Al 1.22 1.08 1.21 1.15 0.97 1.19 1.13 1.12 1.20 1.20 1.12 wt% Continued on next page Fe wt% 27.57 24.53 27.49 26.10 21.91 27.07 25.72 25.31 27.24 27.32 25.34 Table E.3 – continued from previous page S 0.31 3.78 0.61 3.47 3.38 1.64 3.88 3.90 1.44 1.13 3.91 wt% O wt% 32.71 36.42 32.61 33.16 42.06 32.64 33.67 34.65 32.40 32.53 38.24 Mg wt% 17.36 15.43 17.30 16.41 13.79 17.02 16.17 15.92 17.14 17.19 11.98 H 0.00 0.23 0.00 0.00 1.34 0.00 0.00 0.00 0.00 0.00 0.27 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 319 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 16.58 13.38 16.48 15.57 12.31 16.17 15.38 14.72 16.23 16.39 14.49 Ti 0.08 0.07 0.08 0.08 0.06 0.08 0.08 0.07 0.08 0.08 0.07 wt% P 0.14 0.11 0.14 0.13 0.10 0.13 0.13 0.12 0.13 0.13 0.12 wt% Cr 0.36 0.29 0.36 0.34 0.27 0.35 0.33 0.32 0.35 0.36 0.31 wt% Ni 1.66 1.34 1.65 1.56 1.23 1.62 1.54 1.47 1.63 1.64 1.45 wt% Na 0.65 0.52 0.65 0.61 0.48 0.63 0.60 0.58 0.64 0.64 0.57 wt% years 6 Ca 10 1.24 1.00 1.23 1.16 0.92 1.21 1.15 1.10 1.21 1.22 1.08 wt% × t=3 Al 1.21 0.98 1.20 1.14 0.90 1.18 1.12 1.07 1.18 1.19 1.06 wt% Fe wt% 27.42 22.13 27.26 25.75 20.35 26.75 25.45 24.33 26.86 27.07 23.95 Table E.3 – continued from previous page S 0.89 3.42 1.46 3.71 3.14 2.75 3.92 3.76 2.39 1.33 3.70 wt% O wt% 32.52 41.60 32.34 33.76 45.39 32.29 34.29 36.82 32.39 32.89 37.72 Mg wt% 17.25 13.92 17.15 16.20 12.81 16.83 16.00 15.32 16.89 17.05 15.07 H 0.00 1.25 0.00 0.00 2.03 0.00 0.00 0.31 0.00 0.00 0.40 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 3-5 4-3 4-4 4-5 Simulation 320 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 4.26 5.43 2.88 3.14 4.84 2.95 5.11 4.92 wt% 10.25 11.81 Ti 1.11 0.99 1.16 0.48 0.39 1.10 1.00 1.14 1.06 1.06 wt% P 0.00 0.00 0.00 0.00 0.09 0.00 0.00 0.00 0.00 0.00 wt% Cr 0.00 0.00 0.00 0.07 0.14 0.00 0.01 0.00 0.00 0.00 wt% Ni 0.00 0.15 0.00 1.16 1.10 0.00 0.12 0.00 0.00 0.00 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 5 Ca 10 8.78 8.23 9.12 4.10 3.10 9.88 8.54 9.25 8.68 8.95 wt% × t=2.5 Al wt% 31.38 29.59 34.85 13.89 11.09 31.15 29.43 34.59 31.80 30.11 Continued on next page Fe 0.00 1.33 0.00 0.00 1.59 0.00 0.00 0.00 wt% 14.05 15.09 S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O wt% 44.12 43.76 44.35 37.96 37.74 43.41 43.34 44.29 44.52 44.06 Mg 7.64 7.79 8.82 wt% 18.03 19.46 11.31 11.13 10.90 10.354 10.517 Table E.4: Predicted bulk elemental planetary abundances for HD177830. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 321 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.45 wt% 14.30 11.11 10.19 12.34 13.86 14.12 12.08 10.23 14.19 Ti 0.20 0.46 0.64 0.32 0.06 0.06 0.36 0.60 0.58 0.20 wt% P 0.00 0.02 0.00 0.02 0.11 0.15 0.03 0.00 0.03 0.00 wt% Cr 0.12 0.08 0.01 0.13 0.23 0.23 0.11 0.02 0.08 0.10 wt% Ni 1.28 0.98 0.49 1.28 1.83 1.86 1.17 1.07 0.60 1.34 wt% Na 0.00 0.03 0.00 0.04 0.19 0.00 0.02 0.00 0.00 0.00 wt% years 5 Ca 1.84 4.07 5.72 2.82 0.52 0.53 3.19 5.42 5.09 1.84 wt% 10 × t=5 Al 5.80 9.13 1.64 1.67 5.79 wt% 13.00 18.07 10.16 17.11 16.44 Continued on next page Fe 3.91 8.49 8.30 wt% 15.29 11.67 16.37 25.16 25.63 14.54 15.77 Table E.4 – continued from previous page S 0.00 0.00 0.00 0.00 1.60 0.00 0.00 0.00 0.00 0.00 wt% O wt% 37.75 39.05 42.53 37.07 32.48 32.99 37.86 40.17 40.62 37.50 Mg wt% 18.43 20.49 22.33 22.75 20.48 17.65 18.03 23.26 23.443 19.539 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 322 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.98 14.05 14.19 13.97 13.39 13.65 13.97 14.16 14.05 14.05 Ti 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 wt% P 0.16 0.12 0.02 0.17 0.18 0.18 0.14 0.04 0.12 0.16 wt% Cr 0.23 0.22 0.22 0.23 0.22 0.23 0.23 0.23 0.22 0.23 wt% Ni 1.84 1.85 1.87 1.84 1.76 1.80 1.84 1.87 1.85 1.85 wt% Na 0.27 0.14 0.00 0.23 0.32 0.31 0.19 0.00 0.10 0.28 wt% years 6 Ca 0.53 0.53 0.54 0.53 0.50 0.51 0.53 0.53 0.53 0.53 wt% 10 × t=1 Al 1.65 1.68 1.69 1.65 1.58 1.61 1.65 1.68 1.67 1.66 wt% Continued on next page Fe wt% 25.37 25.29 25.47 25.35 24.27 24.77 25.21 25.65 25.26 25.50 Table E.4 – continued from previous page S 0.51 0.44 0.00 0.65 4.60 2.71 0.88 0.00 0.55 0.00 wt% O wt% 32.89 32.95 33.04 32.85 31.56 32.18 32.79 32.96 32.93 33.06 Mg wt% 22.91 22.49 21.56 21.98 22.52 22.82 22.66 22.62 22.521 22.675 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 323 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.62 13.77 14.06 13.58 13.17 13.21 13.66 14.05 13.79 13.66 Ti 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 0.06 wt% P 0.18 0.18 0.19 0.18 0.17 0.18 0.18 0.19 0.18 0.18 wt% Cr 0.22 0.23 0.23 0.22 0.22 0.22 0.23 0.23 0.23 0.23 wt% Ni 1.80 1.82 1.85 1.79 1.74 1.74 1.80 1.85 1.82 1.80 wt% Na 0.31 0.26 0.19 0.31 0.31 0.30 0.28 0.22 0.25 0.31 wt% years 6 Ca 10 0.51 0.52 0.53 0.51 0.50 0.50 0.51 0.53 0.52 0.51 wt% × t=1.5 Al 1.61 1.63 1.66 1.60 1.56 1.56 1.61 1.66 1.63 1.61 wt% Continued on next page Fe wt% 24.72 24.99 25.51 24.65 23.89 23.97 24.78 25.50 25.03 24.79 Table E.4 – continued from previous page S 2.94 1.93 0.00 3.19 6.12 5.87 2.73 0.00 1.82 2.67 wt% O wt% 32.10 32.44 33.08 32.02 31.06 31.14 32.17 33.08 32.48 32.19 Mg wt% 21.93 22.18 22.64 21.87 21.21 21.27 21.99 22.63 22.21 21.99 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 324 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.36 13.53 13.91 13.30 12.79 13.14 13.40 13.85 13.54 13.31 Ti 0.06 0.06 0.06 0.06 0.05 0.05 0.06 0.06 0.06 0.06 wt% P 0.18 0.18 0.18 0.18 0.17 0.17 0.18 0.18 0.18 0.18 wt% Cr 0.22 0.22 0.23 0.22 0.21 0.22 0.22 0.23 0.22 0.22 wt% Ni 1.76 1.79 1.84 1.75 1.69 1.73 1.77 1.83 1.79 1.76 wt% Na 0.31 0.31 0.32 0.31 0.31 0.32 0.31 0.32 0.31 0.31 wt% years 6 Ca 0.50 0.51 0.52 0.50 0.48 0.49 0.50 0.52 0.51 0.50 wt% 10 × t=2 Al 1.58 1.60 1.64 1.57 1.51 1.55 1.58 1.64 1.60 1.57 wt% Continued on next page Fe wt% 24.24 24.56 25.24 24.14 23.17 23.82 24.31 25.15 24.57 24.16 Table E.4 – continued from previous page S 4.80 3.56 0.87 5.19 6.20 6.35 4.29 1.25 3.48 5.15 wt% O wt% 31.49 31.90 32.79 31.36 32.83 30.99 31.81 32.66 31.93 31.37 Mg wt% 21.51 21.79 22.40 21.42 20.60 21.16 21.57 22.31 21.81 21.43 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 325 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.23 13.30 13.58 13.12 12.78 13.13 13.25 13.46 13.42 13.18 Ti 0.06 0.06 0.06 0.05 0.05 0.05 0.06 0.06 0.06 0.06 wt% P 0.18 0.18 0.18 0.17 0.17 0.17 0.18 0.18 0.18 0.18 wt% Cr 0.22 0.22 0.22 0.22 0.21 0.22 0.22 0.22 0.22 0.22 wt% Ni 1.74 1.75 1.79 1.73 1.68 1.73 1.75 1.78 1.77 1.74 wt% Na 0.31 0.31 0.31 0.31 0.32 0.32 0.31 0.31 0.31 0.31 wt% years 6 Ca 10 0.50 0.50 0.51 0.49 0.48 0.49 0.50 0.51 0.50 0.50 wt% × t=2.5 Al 1.56 1.57 1.60 1.55 1.51 1.55 1.56 1.59 1.58 1.56 wt% Continued on next page Fe wt% 23.99 24.14 24.65 23.81 23.15 23.81 24.03 24.44 24.35 23.92 Table E.4 – continued from previous page S 5.73 4.87 3.20 5.99 6.22 6.39 5.30 4.06 4.37 6.04 wt% O wt% 31.18 31.69 32.02 31.41 32.84 30.97 31.51 31.73 31.63 31.08 Mg wt% 21.87 21.14 20.58 21.15 21.33 21.67 21.61 21.23 21.301 21.418 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 326 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.17 13.17 13.31 13.08 12.78 12.84 13.10 13.27 13.15 13.15 Ti 0.06 0.06 0.06 0.05 0.05 0.05 0.05 0.06 0.05 0.05 wt% P 0.17 0.17 0.18 0.17 0.17 0.17 0.17 0.18 0.17 0.17 wt% Cr 0.22 0.22 0.22 0.22 0.21 0.21 0.22 0.22 0.22 0.22 wt% Ni 1.74 1.74 1.76 1.72 1.68 1.69 1.73 1.75 1.73 1.73 wt% Na 0.32 0.31 0.31 0.31 0.32 0.31 0.31 0.31 0.31 0.31 wt% years 6 Ca 0.50 0.50 0.50 0.49 0.48 0.48 0.49 0.50 0.49 0.49 wt% 10 × t=3 Al 1.56 1.56 1.57 1.55 1.51 1.52 1.55 1.57 1.55 1.55 wt% Fe wt% 23.88 23.89 24.17 23.73 23.15 23.30 23.76 24.09 23.85 23.84 Table E.4 – continued from previous page S 6.14 5.74 5.09 6.24 6.22 6.25 5.93 5.41 5.67 6.29 wt% O wt% 31.05 31.44 31.39 31.36 32.84 32.49 31.59 31.28 31.63 31.00 Mg wt% 21.44 21.07 20.58 20.68 21.10 21.37 21.17 21.18 21.210 21.208 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 1-4 2-3 2-4 2-5 3-3 3-4 4-3 4-4 4-5 Simulation 327 C 9.05 9.41 8.12 0.00 5.48 0.00 0.00 0.00 0.00 4.70 wt% N 0.00 0.00 0.00 0.00 0.01 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 12.89 12.79 13.09 13.80 13.58 13.85 13.83 13.80 12.99 12.90 Ti 0.06 0.06 0.06 0.07 0.07 0.07 0.07 0.07 0.06 0.06 wt% P 0.15 0.15 0.15 0.16 0.16 0.16 0.16 0.16 0.15 0.15 wt% Cr 0.27 0.27 0.28 0.29 0.29 0.29 0.29 0.29 0.28 0.27 wt% Ni 1.48 1.47 1.51 1.59 1.56 1.59 1.59 1.59 1.49 1.48 wt% Na 0.45 0.45 0.46 0.48 0.23 0.48 0.48 0.48 0.45 0.45 wt% years years 5 5 10 Ca 10 0.72 0.71 0.73 0.77 0.75 0.77 0.77 0.77 0.72 0.72 wt% × × t=5 t=2.5 Al 1.36 1.35 1.38 1.45 1.43 1.46 1.46 1.45 1.37 1.36 wt% Continued on next page Fe wt% 24.50 24.29 24.88 26.23 25.79 26.31 26.28 26.22 24.68 24.52 S 2.59 2.95 2.22 4.55 2.51 4.54 4.55 4.49 4.28 3.64 wt% O wt% 28.44 28.20 28.81 30.90 29.47 30.88 30.89 31.25 35.35 31.68 Mg wt% 18.04 17.89 18.31 19.71 18.66 19.59 19.62 19.42 18.17 18.05 Table E.5: Predicted bulk elemental planetary abundances for 55Cnc. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 3-4 3-5 4-4 1-4 2-4 3-4 3-5 4-4 Simulation 328 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.24 wt% 12.98 12.96 12.96 12.18 13.27 12.86 12.84 12.75 11.15 Ti 0.06 0.06 0.06 0.06 0.07 0.06 0.06 0.06 0.05 0.05 wt% P 0.15 0.15 0.15 0.14 0.16 0.15 0.15 0.15 0.11 0.13 wt% Cr 0.28 0.28 0.28 0.26 0.28 0.27 0.27 0.27 0.20 0.24 wt% Ni 1.49 1.49 1.49 1.40 1.53 1.48 1.48 1.47 1.06 1.28 wt% Na 0.45 0.45 0.45 0.00 0.46 0.45 0.45 0.32 0.00 0.23 wt% years years 6 6 10 Ca 10 0.72 0.72 0.72 0.68 0.74 0.71 0.71 0.71 0.51 0.62 wt% × × t=1 t=1.5 Al 1.37 1.36 1.36 1.28 1.40 1.35 1.35 1.34 0.97 1.17 wt% Continued on next page Fe wt% 24.67 24.65 24.62 23.07 25.18 24.42 24.40 24.20 17.49 21.19 Table E.5 – continued from previous page S 4.28 4.27 4.27 4.00 4.36 4.23 4.23 4.19 3.03 3.67 wt% O wt% 35.31 35.36 35.23 39.12 33.74 35.90 35.96 36.49 51.14 42.88 Mg wt% 18.24 18.25 18.39 17.04 18.68 17.99 17.96 17.83 12.92 15.80 H 0.00 0.00 0.00 0.77 0.14 0.11 0.13 0.23 3.29 1.59 wt% 1-4 2-4 3-4 3-5 4-4 1-4 2-4 3-4 3-5 4-4 Simulation 329 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 9.21 9.21 wt% 11.75 12.37 11.76 11.15 10.13 10.18 11.43 11.15 Ti 0.06 0.06 0.06 0.05 0.05 0.05 0.05 0.06 0.05 0.05 wt% P 0.14 0.15 0.14 0.11 0.13 0.12 0.12 0.13 0.11 0.13 wt% Cr 0.25 0.26 0.25 0.20 0.24 0.21 0.22 0.24 0.20 0.24 wt% Ni 1.35 1.42 1.35 1.06 1.28 1.16 1.17 1.31 1.06 1.28 wt% Na 0.21 0.14 0.31 0.00 0.23 0.10 0.00 0.10 0.00 0.23 wt% years years 6 6 10 Ca 10 0.65 0.69 0.65 0.51 0.62 0.56 0.56 0.63 0.51 0.62 wt% × × t=2 t=2.5 Al 1.24 1.30 1.24 0.97 1.17 1.07 1.07 1.20 0.97 1.17 wt% Continued on next page Fe wt% 22.26 23.44 22.33 17.44 21.16 19.17 19.27 21.65 17.45 21.19 Table E.5 – continued from previous page S 3.86 4.06 3.87 3.02 3.67 3.32 3.34 3.75 3.02 3.67 wt% O wt% 40.70 38.22 40.52 51.25 43.01 47.43 47.30 42.11 51.24 42.88 Mg wt% 16.43 17.30 16.45 12.88 15.69 14.16 14.24 15.98 12.88 15.80 H 1.11 0.59 1.08 3.31 1.59 2.51 2.48 1.40 3.31 1.59 wt% 1-4 2-4 3-4 3-5 4-4 1-4 2-4 3-4 3-5 4-4 Simulation 330 C 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 wt% Si 9.88 9.24 9.99 9.21 wt% 11.11 Ti 0.05 0.05 0.05 0.05 0.05 wt% P 0.12 0.11 0.12 0.11 0.13 wt% Cr 0.21 0.20 0.21 0.20 0.24 wt% Ni 1.13 1.06 1.15 1.06 1.28 wt% Na 0.00 0.00 0.00 0.00 0.23 wt% years 6 Ca 10 0.55 0.51 0.55 0.51 0.62 wt% × t=3 Al 1.04 0.97 1.05 0.97 1.17 wt% Fe wt% 18.71 17.50 18.92 17.45 21.08 Table E.5 – continued from previous page S 3.24 3.03 3.28 3.02 3.65 wt% O wt% 48.52 51.12 48.05 51.24 43.26 Mg wt% 13.82 12.93 13.98 12.88 15.54 H 2.74 3.28 2.64 3.31 1.65 wt% 1-4 2-4 3-4 3-5 4-4 Simulation 331 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si 2.79 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Ti 2.44 2.62 3.02 2.59 2.62 2.60 2.62 2.60 wt% P 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Cr 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Ni 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 5 10 Ca 8.69 0.00 1.12 0.00 0.00 0.00 0.00 0.00 × wt% t=2.5 Al wt% Continued on next page 40.18 49.18 48.08 49.21 49.18 49.20 49.18 49.19 Fe 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O wt% 45.90 48.20 47.78 48.20 48.20 48.20 48.20 48.20 Mg 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Table E.6: Predicted bulk elemental planetary abundances for HD142415. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 332 C 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% N 1.94 6.54 2.12 0.00 6.54 3.23 6.54 0.00 wt% Si 3.32 0.00 5.22 7.75 0.00 4.25 0.00 6.95 wt% Ti 2.01 2.20 wt% 14.70 47.42 17.15 47.42 24.36 47.42 P 0.17 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Cr 0.14 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Ni 2.29 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 5 Ca 10 wt% 16.13 20.95 22.80 23.80 20.95 23.25 20.95 21.68 × t=5 Al 8.28 0.00 0.00 0.00 wt% Continued on next page 16.25 24.50 10.62 26.71 Fe 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 28.92 Table E.6 – continued from previous page S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O wt% 23.45 25.09 36.45 41.94 25.09 33.18 25.09 42.46 Mg 0.66 0.00 0.00 0.00 0.00 1.11 0.00 0.00 wt% H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 333 C 5.00 0.00 0.84 0.00 0.00 0.00 0.00 0.00 wt% N 0.03 0.00 0.00 0.06 0.00 0.03 0.00 0.00 wt% Si 8.59 4.42 6.43 5.67 4.42 4.69 4.42 1.28 wt% Ti 0.14 0.14 0.19 0.21 0.14 0.18 0.14 0.27 wt% P 0.22 0.30 0.28 0.28 0.30 0.29 0.30 0.35 wt% Cr 0.67 0.82 0.74 0.91 0.82 0.86 0.82 0.71 wt% Ni 3.18 4.25 4.04 4.04 4.25 4.17 4.25 5.01 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 6 Ca 10 2.34 3.15 3.04 2.95 3.15 3.08 3.15 3.89 wt% × t=1 Al 1.98 2.67 2.57 2.50 2.67 2.60 2.67 3.30 wt% Continued on next page Fe wt% 53.18 70.30 65.80 68.18 70.30 69.58 70.30 78.44 Table E.6 – continued from previous page S 0.40 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O 6.38 wt% 17.06 10.83 11.59 11.95 10.83 11.16 10.83 Mg 7.21 3.11 4.48 3.25 3.11 3.37 3.11 0.37 wt% H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 334 C wt% 10.28 11.54 10.51 12.46 11.54 11.79 11.54 11.54 N 0.00 0.00 0.01 0.03 0.00 0.00 0.00 0.00 wt% Si wt% 14.32 13.82 14.23 14.50 13.82 14.05 13.82 13.82 Ti 0.09 0.09 0.10 0.09 0.09 0.09 0.09 0.09 wt% P 0.12 0.13 0.13 0.12 0.13 0.13 0.13 0.13 wt% Cr 0.41 0.42 0.43 0.39 0.42 0.41 0.42 0.42 wt% Ni 1.78 1.84 1.87 1.72 1.84 1.81 1.84 1.84 wt% Na 0.22 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 6 10 Ca 1.30 1.35 1.37 1.26 1.35 1.32 1.35 1.35 × wt% t=1.5 Al 1.10 1.14 1.16 1.06 1.14 1.12 1.14 1.14 wt% Continued on next page Fe wt% 30.06 31.15 31.72 29.11 31.15 30.62 31.15 31.15 Table E.6 – continued from previous page S 0.99 1.08 1.10 1.00 1.08 1.06 1.08 1.08 wt% O wt% 26.27 24.97 23.92 25.72 24.97 25.18 24.97 24.97 Mg wt% 13.04 12.47 13.47 12.54 12.47 12.41 12.47 12.47 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 335 C wt% 11.34 12.16 11.71 11.92 12.16 12.03 12.16 12.33 N 0.02 0.03 0.02 0.00 0.03 0.03 0.03 0.03 wt% Si wt% 14.02 14.21 14.33 14.28 14.21 14.22 14.21 14.29 Ti 0.08 0.09 0.08 0.06 0.09 0.09 0.09 0.09 wt% P 0.11 0.12 0.12 0.12 0.12 0.12 0.12 0.12 wt% Cr 0.38 0.38 0.38 0.38 0.38 0.38 0.38 0.38 wt% Ni 1.65 1.67 1.68 1.68 1.67 1.67 1.67 1.68 wt% Na 0.49 0.43 0.39 0.53 0.43 0.47 0.43 0.17 wt% years 6 Ca 10 1.20 1.22 1.23 1.23 1.22 1.22 1.22 1.23 wt% × t=2 Al 1.02 1.04 1.05 1.04 1.04 1.04 1.04 1.04 wt% Continued on next page Fe wt% 27.86 28.28 28.49 28.42 28.28 28.28 28.28 28.39 Table E.6 – continued from previous page S 2.51 0.98 0.63 0.00 0.98 0.98 0.98 0.98 wt% O wt% 26.07 26.05 26.44 26.82 26.05 26.10 26.05 26.01 Mg wt% 13.25 13.35 13.46 13.54 13.35 13.39 13.35 13.29 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 336 C 5.45 9.68 6.86 9.68 9.10 9.68 wt% 10.85 10.53 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.84 14.63 14.26 14.79 14.63 14.67 14.63 14.51 Ti 0.09 0.09 0.08 0.09 0.09 0.09 0.09 0.06 wt% P 0.12 0.12 0.12 0.12 0.12 0.12 0.12 0.12 wt% Cr 0.40 0.39 0.38 0.40 0.39 0.39 0.39 0.39 wt% Ni 1.74 1.72 1.68 1.74 1.72 1.72 1.72 1.71 wt% Na 0.55 0.55 0.53 0.55 0.55 0.55 0.55 0.54 wt% years 6 10 Ca 1.27 1.26 1.22 1.27 1.26 1.26 1.26 1.25 × wt% t=2.5 Al 1.08 1.07 1.04 1.08 1.07 1.07 1.07 1.06 wt% Continued on next page Fe wt% 29.49 29.07 28.33 29.38 29.07 29.15 29.07 28.84 Table E.6 – continued from previous page S 2.73 0.00 1.08 1.66 0.00 0.25 0.00 0.00 wt% O wt% 28.14 27.57 26.90 28.04 27.57 27.71 27.57 27.25 Mg wt% 14.08 13.88 13.53 14.03 13.88 13.92 13.88 13.76 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 337 C 7.21 wt% 12.20 11.60 11.46 12.44 12.03 12.20 12.44 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.23 13.71 13.74 13.65 13.79 13.69 13.71 13.79 Ti 0.09 0.08 0.07 0.08 0.05 0.08 0.08 0.05 wt% P 0.12 0.11 0.11 0.11 0.11 0.11 0.11 0.11 wt% Cr 0.38 0.37 0.37 0.36 0.37 0.37 0.37 0.37 wt% Ni 1.67 1.61 1.61 1.60 1.62 1.61 1.61 1.62 wt% Na 0.53 0.51 0.51 0.51 0.52 0.51 0.51 0.52 wt% years 6 Ca 10 1.22 1.18 1.18 1.17 1.19 1.17 1.18 1.19 wt% × t=3 Al 1.04 1.00 1.00 0.99 1.01 1.00 1.00 1.01 wt% Fe wt% 28.27 27.22 27.30 27.15 27.47 27.20 27.22 27.47 Table E.6 – continued from previous page S 4.69 3.02 3.45 4.05 2.36 3.30 3.02 2.36 wt% O wt% 27.05 25.98 26.01 25.89 26.00 25.94 25.98 26.00 Mg wt% 13.51 13.02 13.04 12.96 13.08 12.99 13.02 13.08 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 5-3 6-3 7-3 8-3 Simulation 338 C wt% 28.91 29.89 28.29 29.89 14.42 29.89 28.66 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 69.42 69.58 69.31 69.59 33.54 69.57 69.38 Ti 0.51 0.53 0.50 0.52 0.24 0.54 0.50 wt% P 0.00 0.00 0.00 0.00 0.43 0.00 0.00 wt% Cr 0.00 0.00 0.01 0.00 0.00 0.00 0.00 wt% Ni 0.13 0.00 0.21 0.00 0.00 0.00 0.16 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years 5 10 Ca 0.00 0.00 0.00 0.00 0.00 0.00 0.00 × wt% t=2.5 Al 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Continued on next page Fe 1.03 0.00 1.69 0.00 0.00 1.30 wt% 51.37 S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Mg 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Table E.7: Predicted bulk elemental planetary abundances for HD19994. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 339 C wt% 18.65 16.81 20.59 16.41 19.77 17.39 16.57 N 0.00 0.21 0.00 0.00 0.02 0.00 0.00 wt% Si wt% 51.15 45.92 54.30 38.60 13.81 60.17 35.57 Ti 0.32 0.28 0.35 0.26 0.05 0.30 0.24 wt% P 0.19 0.23 0.19 0.62 0.13 0.00 0.28 wt% Cr 0.14 0.23 0.03 0.43 0.32 0.08 0.25 wt% Ni 1.81 2.30 0.75 2.20 1.66 2.10 1.74 wt% Na 0.00 0.00 0.00 0.00 0.19 0.00 0.00 wt% years 5 Ca 10 0.00 0.03 0.00 1.43 1.09 0.00 0.73 wt% × t=5 Al 0.00 0.41 0.00 1.57 1.18 0.00 0.82 wt% Continued on next page Fe wt% 27.74 32.72 23.79 35.49 24.54 19.96 36.71 Table E.7 – continued from previous page S 0.00 0.51 0.00 0.00 0.87 0.00 0.35 wt% O 0.00 0.00 0.00 2.72 0.00 4.18 wt% 24.60 Mg 0.00 0.37 0.00 0.27 0.00 2.55 wt% 11.77 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 340 C wt% 14.17 14.13 14.34 14.49 19.15 10.10 17.80 N 0.01 0.25 0.16 0.00 0.00 0.36 0.01 wt% Si wt% 17.27 16.63 16.73 17.38 13.53 18.67 14.62 Ti 0.08 0.08 0.08 0.07 0.05 0.11 0.05 wt% P 0.22 0.21 0.21 0.18 0.13 0.29 0.15 wt% Cr 0.48 0.48 0.49 0.44 0.32 0.63 0.36 wt% Ni 2.68 2.65 2.56 2.28 1.62 3.57 1.83 wt% Na 0.12 0.22 0.03 0.54 0.90 0.00 0.38 wt% years 6 Ca 10 1.18 1.23 1.52 1.49 1.06 1.02 1.19 wt% × t=1 Al 1.69 1.84 1.81 1.63 1.16 2.44 1.30 wt% Continued on next page Fe wt% 38.77 38.70 37.64 33.67 23.95 51.88 27.00 Table E.7 – continued from previous page S 0.72 1.35 1.35 0.69 1.29 1.35 0.64 wt% O 6.32 wt% 14.68 14.58 15.26 18.11 25.12 23.16 Mg 7.93 7.65 7.83 9.03 3.25 wt% 11.71 11.52 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 341 C 6.24 wt% 18.88 18.90 19.11 18.80 19.62 18.16 N 0.01 0.00 0.01 0.01 0.00 0.01 0.00 wt% Si wt% 13.84 13.75 13.79 13.61 15.28 13.66 13.78 Ti 0.05 0.05 0.05 0.05 0.06 0.05 0.05 wt% P 0.13 0.13 0.13 0.13 0.15 0.13 0.13 wt% Cr 0.33 0.33 0.32 0.32 0.36 0.33 0.32 wt% Ni 1.68 1.68 1.67 1.64 1.84 1.68 1.65 wt% Na 0.62 0.58 0.61 0.79 1.02 0.25 0.88 wt% years 6 10 Ca 1.10 1.09 1.09 1.07 1.20 1.10 1.08 × wt% t=1.5 Al 1.20 1.15 1.15 1.17 1.31 1.15 1.18 wt% Continued on next page Fe wt% 24.81 24.72 24.61 24.11 27.06 24.82 24.39 Table E.7 – continued from previous page S 0.53 0.58 0.41 1.61 3.71 0.74 1.07 wt% O wt% 24.78 25.07 25.09 24.94 28.56 24.57 25.38 Mg wt% 12.05 11.96 11.95 11.76 13.23 11.89 11.92 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 342 C 0.00 wt% 17.93 17.89 18.33 12.65 18.37 13.07 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 13.85 13.80 13.81 14.48 16.26 13.90 14.42 Ti 0.05 0.05 0.05 0.05 0.06 0.04 0.05 wt% P 0.13 0.13 0.13 0.14 0.16 0.13 0.14 wt% Cr 0.32 0.32 0.32 0.34 0.38 0.32 0.34 wt% Ni 1.66 1.66 1.66 1.74 1.95 1.67 1.73 wt% Na 0.88 0.90 0.91 0.96 1.08 0.90 0.96 wt% years 6 Ca 10 1.09 1.08 1.08 1.13 1.27 1.09 1.13 wt% × t=2 Al 1.19 1.18 1.18 1.24 1.39 1.19 1.23 wt% Continued on next page Fe wt% 24.52 24.44 24.46 25.67 28.84 24.62 25.53 Table E.7 – continued from previous page S 0.76 1.09 0.55 2.08 4.15 0.20 2.05 wt% O wt% 25.62 25.50 25.55 26.98 30.38 25.54 26.85 Mg wt% 11.99 11.94 11.95 12.54 14.07 12.02 12.49 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 343 C 8.05 0.00 wt% 15.52 13.83 17.41 17.35 10.46 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.07 14.32 13.76 15.13 16.27 13.98 14.69 Ti 0.05 0.05 0.05 0.06 0.06 0.05 0.05 wt% P 0.14 0.14 0.13 0.15 0.16 0.13 0.14 wt% Cr 0.33 0.33 0.32 0.35 0.38 0.33 0.34 wt% Ni 1.69 1.72 1.65 1.82 1.95 1.68 1.76 wt% Na 0.94 0.95 0.92 1.01 1.08 0.93 0.98 wt% years 6 10 Ca 1.10 1.12 1.08 1.18 1.27 1.09 1.15 × wt% t=2.5 Al 1.20 1.23 1.18 1.30 1.39 1.20 1.26 wt% Continued on next page Fe wt% 24.91 25.36 24.37 26.79 28.80 24.74 26.01 Table E.7 – continued from previous page S 1.73 1.87 1.64 2.88 4.15 0.50 3.05 wt% O wt% 26.16 26.66 25.57 28.20 30.40 25.93 27.38 Mg wt% 12.18 12.40 11.91 13.09 14.08 12.10 12.71 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 344 C 7.26 0.00 5.22 wt% 11.84 10.54 12.97 18.22 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 14.51 14.72 14.31 15.16 15.09 13.62 15.46 Ti 0.05 0.05 0.05 0.06 0.06 0.05 0.06 wt% P 0.14 0.14 0.14 0.15 0.14 0.13 0.15 wt% Cr 0.34 0.34 0.33 0.35 0.35 0.32 0.36 wt% Ni 1.74 1.77 1.72 1.82 1.81 1.64 1.86 wt% Na 0.97 0.98 0.95 1.01 1.00 0.91 1.03 wt% years 6 Ca 10 1.14 1.15 1.12 1.19 1.18 1.07 1.21 wt% × t=3 Al 1.24 1.26 1.22 1.30 1.29 1.17 1.32 wt% Fe wt% 25.70 26.08 25.33 26.83 26.72 24.11 27.38 Table E.7 – continued from previous page S 2.75 2.79 2.81 3.50 3.85 1.67 3.72 wt% O wt% 27.03 27.43 26.66 28.26 35.45 25.31 28.86 Mg wt% 12.56 12.75 12.38 13.12 13.06 11.79 13.38 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 3-4 3-5 4-3 4-4 Simulation 345 C 9.58 wt% 34.91 33.18 34.91 34.41 31.96 25.91 22.89 27.46 29.30 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 1.10 wt% Si wt% 64.53 61.18 64.53 63.59 57.36 45.45 38.23 48.85 51.36 13.29 Ti 0.56 0.52 0.56 0.55 0.49 0.38 0.32 0.41 0.43 0.10 wt% P 0.00 0.00 0.00 0.00 0.00 0.17 0.28 0.10 0.17 0.24 wt% Cr 0.00 0.00 0.00 0.00 0.00 0.06 0.08 0.03 0.03 0.48 wt% Ni 0.00 0.00 0.00 0.00 0.00 1.30 1.74 0.74 0.68 2.56 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years years 5 5 10 Ca 10 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 1.37 × wt% × t=5 t=2.5 Al 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 2.12 wt% Continued on next page Fe 0.00 5.11 0.00 1.45 wt% 10.20 26.72 36.45 22.41 18.03 63.62 S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 3.61 wt% O 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Mg 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 1.91 wt% Table E.8: Predicted bulk elemental planetary abundances for HD108874. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 3-4 4-4 4-5 1-4 2-4 3-4 4-4 4-5 Simulation 346 C 9.09 wt% 12.63 16.64 15.46 20.73 21.07 20.16 20.94 20.72 10.47 N 0.61 0.26 0.78 0.50 0.01 0.01 0.01 0.01 0.01 0.00 wt% Si wt% 12.58 12.94 14.69 12.39 11.67 11.93 11.87 11.43 11.64 12.62 Ti 0.09 0.08 0.10 0.08 0.05 0.05 0.05 0.06 0.05 0.06 wt% P 0.21 0.18 0.25 0.20 0.13 0.14 0.13 0.14 0.14 0.14 wt% Cr 0.45 0.39 0.46 0.40 0.30 0.31 0.30 0.31 0.31 0.32 wt% Ni 2.26 1.89 2.56 2.06 1.40 1.44 1.43 1.48 1.45 1.52 wt% Na 0.00 0.00 0.00 0.00 0.34 0.13 0.32 0.08 0.23 0.46 wt% years years 6 6 10 Ca 10 1.25 1.05 1.13 1.01 0.78 0.80 0.79 0.82 0.80 0.84 × wt% × t=1 t=1.5 Al 2.13 1.59 2.48 1.88 1.17 1.21 1.19 1.24 1.21 1.27 wt% Continued on next page Fe wt% 52.73 36.80 57.32 44.00 24.88 25.56 25.31 26.20 25.67 26.95 Table E.8 – continued from previous page S 2.40 1.43 2.70 1.95 0.32 0.64 0.52 0.66 0.47 3.52 wt% O 6.83 4.28 wt% 16.25 11.85 23.99 23.45 24.01 22.89 23.48 26.40 Mg 5.82 4.17 8.21 wt% 10.50 14.22 13.26 13.90 13.74 13.83 15.42 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 3-4 4-4 4-5 1-4 2-4 3-4 4-4 4-5 Simulation 347 C 0.00 wt% 19.03 13.42 20.01 17.13 14.16 11.72 11.27 15.13 10.76 N 0.00 0.00 0.01 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 11.81 12.38 11.75 11.97 11.83 12.58 12.43 12.19 12.62 13.70 Ti 0.05 0.06 0.05 0.05 0.06 0.06 0.06 0.06 0.06 0.07 wt% P 0.13 0.14 0.13 0.14 0.13 0.14 0.14 0.14 0.14 0.15 wt% Cr 0.30 0.32 0.30 0.31 0.30 0.32 0.32 0.31 0.32 0.35 wt% Ni 1.42 1.49 1.41 1.44 1.42 1.51 1.49 1.47 1.52 1.64 wt% Na 0.42 0.45 0.37 0.41 0.43 0.46 0.46 0.45 0.46 0.50 wt% years years 6 6 10 Ca 10 0.79 0.83 0.78 0.80 0.79 0.84 0.83 0.81 0.84 0.91 × wt% × t=2 t=2.5 Al 1.19 1.25 1.18 1.21 1.19 1.27 1.25 1.23 1.27 1.38 wt% Continued on next page Fe wt% 25.17 26.37 25.05 25.51 25.20 26.80 26.47 26.00 26.89 29.13 Table E.8 – continued from previous page S 0.84 2.41 0.37 1.61 5.21 2.62 4.09 1.87 3.29 6.55 wt% O wt% 24.42 25.77 24.23 24.82 24.83 26.30 26.01 25.46 26.40 28.86 Mg wt% 14.42 15.12 14.33 14.61 14.45 15.37 15.18 14.90 15.42 16.75 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 3-4 4-4 4-5 1-4 2-4 3-4 4-4 4-5 Simulation 348 C 8.81 5.97 6.36 9.48 0.00 wt% N 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 12.70 12.97 13.10 12.56 13.69 Ti 0.06 0.06 0.06 0.06 0.07 wt% P 0.16 0.16 0.15 0.16 0.15 wt% Cr 0.32 0.33 0.33 0.32 0.35 wt% Ni 1.53 1.56 1.58 1.51 1.64 wt% Na 0.47 0.47 0.48 0.46 0.50 wt% years 6 Ca 10 0.89 0.90 0.87 0.87 0.91 wt% × t=3 Al 1.28 1.30 1.32 1.26 1.38 wt% Fe wt% 27.08 27.63 27.92 26.78 29.11 Table E.8 – continued from previous page S 4.59 5.53 4.35 4.83 6.63 wt% O wt% 26.60 27.26 27.46 26.36 28.83 Mg wt% 15.51 15.85 16.01 15.35 16.73 H 0.00 0.00 0.00 0.00 0.00 wt% 1-4 2-4 3-4 4-4 4-5 Simulation 349 C wt% 55.27 55.24 55.33 55.25 52.93 55.31 54.39 50.28 N 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 44.46 44.49 44.40 44.49 42.52 44.42 43.71 40.32 Ti 0.27 0.27 0.27 0.27 0.25 0.26 0.26 0.24 wt% P 0.00 0.00 0.00 0.00 0.26 0.00 0.19 0.52 wt% Cr 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Ni 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years years 5 5 10 Ca 10 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 × wt% × t=5 t=2.5 Al 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Continued on next page Fe 0.00 0.00 0.00 0.00 4.04 0.00 1.45 8.64 wt% S 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% O 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Mg 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% Table E.9: Predicted bulk elemental planetary abundances for HD4203. All values are in wt%. H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 1-3 2-3 3-3 4-3 Simulation 350 C wt% 23.22 24.17 22.85 24.74 27.82 27.17 28.08 27.83 N 1.15 1.11 1.16 0.00 0.01 0.00 0.02 0.01 wt% Si wt% 19.53 19.97 19.36 19.87 11.95 11.36 12.28 11.96 Ti 0.11 0.11 0.11 0.11 0.05 0.06 0.05 0.05 wt% P 0.26 0.27 0.25 0.25 0.13 0.13 0.12 0.12 wt% Cr 0.47 0.37 0.51 0.53 0.27 0.28 0.26 0.25 wt% Ni 2.88 2.92 2.87 2.85 1.44 1.53 1.42 1.37 wt% Na 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% years years 6 6 10 Ca 10 1.41 0.95 1.56 1.55 0.78 0.83 0.77 0.75 × wt% × t=1 t=1.5 Al 2.21 2.13 2.24 2.23 1.13 1.20 1.11 1.07 wt% Continued on next page Fe wt% 43.02 42.78 43.15 42.99 21.71 23.10 21.37 20.71 Table E.9 – continued from previous page S 3.75 3.30 3.91 1.24 0.62 0.66 0.61 0.60 wt% O 0.00 0.00 0.00 2.64 wt% 21.91 19.99 22.49 22.80 Mg 1.99 1.92 2.02 1.00 wt% 12.19 13.67 11.41 12.49 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 1-3 2-3 3-3 4-3 Simulation 351 C wt% 27.45 27.63 27.40 27.04 26.03 27.04 26.04 23.69 N 0.01 0.01 0.01 0.00 0.00 0.00 0.00 0.00 wt% Si wt% 11.82 11.82 11.80 11.90 11.93 11.90 11.91 12.08 Ti 0.05 0.05 0.05 0.03 0.05 0.03 0.05 0.05 wt% P 0.12 0.12 0.12 0.12 0.12 0.12 0.12 0.12 wt% Cr 0.25 0.25 0.25 0.25 0.25 0.25 0.25 0.26 wt% Ni 1.36 1.36 1.36 1.37 1.37 1.37 1.37 1.39 wt% Na 0.42 0.32 0.45 0.47 0.47 0.47 0.47 0.47 wt% years years 6 6 10 Ca 10 0.74 0.74 0.74 0.74 0.74 0.74 0.74 0.75 × wt% × t=2 t=2.5 Al 1.06 1.06 1.06 1.07 1.07 1.07 1.07 1.08 wt% Continued on next page Fe wt% 20.48 20.49 20.44 20.58 20.65 20.58 20.59 20.90 Table E.9 – continued from previous page S 0.43 0.59 0.59 0.00 0.69 0.00 0.81 2.07 wt% O wt% 23.04 22.87 22.95 23.50 23.67 23.50 23.66 24.02 Mg wt% 12.77 12.69 12.78 12.93 12.96 12.93 12.93 13.12 H 0.00 0.00 0.00 0.00 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 1-3 2-3 3-3 4-3 Simulation 352 C wt% 21.98 24.52 22.02 15.79 N 0.00 0.00 0.00 0.00 wt% Si wt% 12.31 11.99 12.29 13.17 Ti 0.05 0.05 0.05 0.06 wt% P 0.12 0.12 0.12 0.13 wt% Cr 0.26 0.25 0.26 0.28 wt% Ni 1.41 1.38 1.41 1.51 wt% Na 0.48 0.47 0.48 0.52 wt% years 6 Ca 10 0.77 0.75 0.77 0.82 wt% × t=3 Al 1.10 1.07 1.10 1.18 wt% Fe wt% 21.29 20.76 21.24 22.79 Table E.9 – continued from previous page S 2.38 1.77 2.48 3.25 wt% O wt% 24.47 23.84 24.43 26.19 Mg wt% 13.36 13.02 13.34 14.30 H 0.00 0.00 0.00 0.00 wt% 1-3 2-3 3-3 4-3 Simulation 353

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