3  Elements & Solar System

Part 1: The condensation of the elements and minerals (3.1 - 3.3)
There is evidence that parts of the protoplanetary disk were initially partly evaporated, which cooled again, thereby condensing minerals. Learn the various mechanisms of condensation, as well as the sequence of mineral condensation in such formerly evaporated regions.
Learning Goals
Explain and describe the difference between equilibrium and fractional condensation. Schematic drawing of the condensation sequence, including the most important minerals. Explain the relationship between condensation sequence and cosmochemical character of the elements.

Part 2: The chemical elements and their cosmochemical character (3.4 - 3.6)
The chemical elements behave differently, depending on e.g., ambient pressure or temperature. Various element characteristics have been derived from this and in combination with thermodynamic calculations. Learn about the cosmochemical classification of the elements from tables and the Cosmochemical Periodic Table of the Elements. Further, have a look which elements are primarily concentrates in which of the phases and components in meteorites. Print the Cosmochemical Periodic Table of the Elements for your personal records.
Learning Goals
Explain the 50% condensation temperature. Knowledge of the various cosmochemical characteristics of the elements and their relationship to the condensation temperature. Recall the 10 most abundant elements in meteorites and their cosmochemical character. Knowledge, which elements are abundant in which mineral. Demonstrate the relationship between element/mineral and condensation temperature.

Bonus 1: Formation and Evolution of our Solar System (3.7 - 3.12)
Our Solar System had an eventful start. More and more extra-solar solar- and planetary systems are discovered, in various stages of their early formation. Using this information and together with the much more and much more detailed information from meteorites, it is then possible to reconstruct the formation and evolutionary phases of our Solar System. Describe the events in the first millions of years when planetary systems in general, and our Solar System in particular form and formed, respectively. Learning Goals
Summarise the evolutionary path of a protoplanetary disk from star formation to a planetary system. Relate the isotopic composition of meteorites to what we think happened in the initial stage of our solar system. Describe the growth steps from dust grains to planets.

3.1 Equilibrium Condensation

It is assumed that all material in the protoplanetary disk at some early point in time was more or less completely evaporated, i.e., in gas form. This gas subsequently cooled, and minerals started to condense. The sequence of mineral condensation depends on a number of parameters, such as the composition of the gas, the ambient pressure, or the dust/gas ratio, which is a measure for the O fugacity in the condensing gas. Two principally different condensation possibilities need to be considered: equilibrium and fractional condensation. Equilibrium-condensation plots have the temperature on one, and the fraction of condensed material on the other axis. During equilibrium condensation, all condensed material is always in full equilibrium with the surrounding gas. The first phases to condense during cooling are refractory, e.g., Ca,Al-rich, and the phases found in CAIs. The small amount of Ca,Al-phases and, hence, CAIs, is because of the small amount of refractory elements available. Next major phases condense; these are metal, olivine, pyroxene and feldspar. These form the major components chondrules, matrix and opaque phases. The initial refractory phases are at this stage basically completely consumed during equilibrium condensation.

The material in the inner parts of the protoplanetary disk (few AU) were in the earliest stage of the Solar System most likely completely evaporated. With falling temperatures and low pressures of around 1-10 Pa phases such as spinel, olivine, pyroxene, feldspar, phyllosilicates, etc. started to condense. Equilibrium condensation then means that at each temperature only those phases are present, which energies match the current P- and T-conditions. This means, the phases are in thermodynamical equilibrium with their surrounding. This means, equilibrium condensation is only possible when thermodynamical equilibrium is achieved at each temperature step, i.e., when there is sufficient time at each temperature step for thermodynamical equilibrium.

CAIs contain high-temperature phases such as spinel, hibonite, perovskite oder akermanite. Chondrules primarily contain olivine, pyroxene, feldspar, metal and sulphide. The matrix contains similar minerals, and in addition often phyllosilicates. However, no high-temperature minerals should occur during complete equilibrium condensation. CAIs are therefore evidence for fractional condensation.

✗(Mg,Fe)2Si2O6

✗(Mg,Fe)2Si2O4

✗(Mg,Fe)SiO3

✓(Mg,Fe)2SiO4

✓Anorthite

✗Albite

✗True

✓False

3.2 Fractional Condensation

It is assumed that all material in the protoplanetary disk at some early point in time was more or less completely evaporated, i.e., in gas form. This gas subsequently cooled, and minerals started to condense. The sequence of mineral condensation depends on a number of parameters, such as the composition of the gas, the ambient pressure, or the dust/gas ratio, which is measure for the O fugacity in the condensing gas. Two principally different condensation possibilities need to be considered: equilibrium and fractional condensation. Fractional condensation plots have the temperature on one, and the fraction of material no longer reacting with the surrounding gas (=isolation degree) on the other axis. An isolation degree of 0% represents equilibrium condensation. For a number of minerals, the condensation sequence changes only minor, if at all, with increasing isolation degree. Most changes in the condensation sequence happen at an isolation degree between ~0.4 and 3%. A number of new minerals start to condense, importantly, for example, SiO2. This is interesting, as this appearance can explain the gas-melt exchange that many chondrules, in particular mineralogically zoned chondrules experienced. Another interesting mineral is oldhamite that is found e.g., in enstatite chondrite chondrules. The occurrence of these and other minerals further change the condensation temperatures of other minerals, such as anorthite or albite.

For fractional condensation, the extensive parameters T and P change quicker than it takes for thermodynamical equilibrium between condensed phases and their surrounding. Hence, the condensed phases – i.e., part of these – are not in thermodynamical equilibrium with each other. It is then possible for refractory phases such as e.g., spinel to survive and still be present at lower temperatures, at which these might be incorporated into chondrules or the matrix. This process is analogous to fractional crystallisation.

SiO2, i.e., quartz/tridymite/cristobalite – depending on temperature and cooling rate. The protoplanetary disk has a Mg/Si ratio of about 1. The only minerals condensing from a gas should therefore be olivine and pyroxene, if some Fe is added. As forsteritic olivine does not completely react to pyroxene during fractional condensation, the gas becomes increasingly enriched in SiO, which can then condense SiO2. This might happen at an isolation degree as low as about 0.5.

✗(Mg,Fe)2SiO6

✗(Mg,Fe)SiO4

✓(Mg,Fe)2Si2O6

✓(Mg,Fe)SiO3

✗1220 K

✓1320 K

✗1420 K

✗1520 K

✗Enstatite

✓Diopside

3.3 Condensation Sequence of the Minerals

The mineral condensation sequence is defined by the 50% condensation temperatures of the elements: first Ca,Al-oxides, then silicates, and finally OH-containing minerals.

Refractory: >1400 K; Main: 1400 K; Volatile: 1230 K; Highly Volatile: 640 K

CAI & AOA: corundum, hibonite, perovskite, spinel, melilite, diopside; chondrules: forsterite, enstatite, feldspar, metal; matrix: olivine, pyroxene, metal, sulphide.

✓CaMgSi2O6

✗CaMgSiO4

✓(Mg,Ca)SiO3

✗CaMgSi3O6

✓Fe(2+)Fe(3+)2O4

✗Fe(2+)Fe(3+)O3

✗Fe(3+)Fe(2+)O4

✗Mg(2+)Fe(3+)2O4

✓spinel – forsterite – enstatite – troilite

✓forsterite – metal – albite – serpentine

✗spinel – Korund – enstatite – phyllosilicate

✓perovskite – enstatite – troilite – metal – phyllosilicate

✗enstatite – metal – forsterite – troilite

3.4 Cosmochemical Periodic Table of the Elements

A specific periodic table for Cosmochemistry is highly beneficial to quickly reference fundamental information such as abundances, condensation temperatures, isotope data, etc.

For example, CI abundance, 50% condensation temperature, oxid-element conversion factor, condensing phase, number of stable isotopes, most important radioactive isotope.

Colour: cosmochemical property – refractory, main, volatile, highly-volatile. Hatching: litophile, sidero-/chalcophile.

✗… lithophile volatile

✗… siderophile volatile

✓… lithophile refractory

✗… siderophile refractory

✗<1400 K

✗<1350 K

✗<1250 K

✓>1400 K

✗>1350 K

✗>1250 K

3.5 Cosmochemical Characteristics of the Elements

The chemical elements are broadly classified into what is called their ›cosmochemical character‹. This consists of two properties: (i) the 50% condensation temperature (T50%), and (ii) whether an element is lithophile or siderophile/chalcophile. Elements with T50% above 1400 K are called ›refractory‹ elements. The most important lithophile refractory elements are Al, Ca, Ti and the REE. The more important siderophile refractory elements are W, Ir and Pt. Elements with T50% between 1250 and 1400 K are called ›main‹ elements, as the 3 main elements Mg, Si and Fe condense in this interval. An additional important lithophile main element is Cr. Additionally important siderophile main elements are Ni and Co. Elements with T50% between 640 and 1250 K are called ›volatile‹ elements. The most important lithophile volatile elements are Mn, Na and P. Important siderophile volatile elements are S and Cu. Elements with T50% below 640 K are called ›highly volatile‹ elements. The most important lithophile highly volatile elements are H, O, C, and the noble gases. The most important siderophile highly volatile element is Pb. Interestingly, although O is a highly volatile element, it condenses as the first element together with Al in the most refractory phase corundum. This illustrates the important fact that not only the T50% is important, but also the phase into which an element condenses.

Cosmochemical Charakcer: refractory – main – volatile – highly-volatile; Geochemichal Character: lithophile – sider- & chalcophile.

For example: refractory lithophile: Ca, Al, Ti, SEE; refractory sidero-/chalcophile: Ir, W; main lithophile: Mg, Si, Cr; main sidero-/chalcophile: Fe, Ni, Co; volatile lithophile: Na, Mn; volatile sidero-/chalcophile: S, Cu, Au; highly-volatile lithophile: O, C, Cl, Edel-Gase; highly-volatile sidero-/chalcophile: Pb, Hg.

✗U – Pb

✓Hf – W

✗Sm – Nd

✗Rb – Sr

✗0.001 Pa

✗0.1 Pa

✓10 Pa

✗1000 Pa

✓0.001 Bar

✗0.1 Bar

✗10 Bar

✗1000 Bar

✓1250-1400 K

✗1250-1400 ºC

✗1150-1400 K

✗1250-1400 ºC

✗1350-1500 K

✗1350-1500 ºC

3.6 Fundamental Properties of Chondrites

Among the fundamental chondrite properties are the modal abundances of their components: chondrules (plus their average sizes), matrix, Ca,Al-rich inclusions (CAIs) and metal, as well as sulphide. It is often not discriminated between the latter two. Further, it is interesting to know either the type meteorite for a specific chondrite group, or which criteria are else used to name a chondrite group.

Carbonaceous chondrites, except for CI. Also, CK chondrites have low chondrule abundances, likely because of their often higher petrologic types, which recrystallises their initial texture. Further R and K chondrites.

Carbonaceous chondrites (8 groups) contain ›C‹ (for carbonaceous) as first letter, followed by a second letter ,which is the first letter of the type meteorite for the group. An exception are CH chondrites, here ›H‹ and for ›high iron‹. Ordinary chondrite (3 groups) are discriminated according to their iron and metal content in ›H‹ for high iron, ›L‹ vor low iron, and ›LL‹ for low iron, low metal. The two enstatite groups have an ›E‹ as first letter, followed by either ›H‹ or ›L‹, with again standing for high and low iron, respectively. The individual groups R and K are named after their respective type meteorites.

✗<0.1 mm

✓<1 mm_x000B_<10 mm

✗ca. 3 vol% in CR chondrites

✗ca. 4 vol% in CR chondrites

✗ca. 4 vol% in CK chondrites

✓ca. 3 vol% in CV chondrites

✗ca. 4 vol% in CV chondrites

✓CC

✗OC

✗EC

✗K

✗R

3.7 Modelling Star Formation from a Molecular Cloud

Stars – including our Sun – formed in clusters, not as individual stars. It is possible to model simulations of such star cluster formations within a collapsing molecular cloud. It can then for example be studied how many stars develop a protoplanetary disk and how stable and long-lived such disks are. However, this needs to be taken cautiously, as the models currently might still be too limited for quantitative studies of in particular protoplanetary disks.

So for more for star formation. The formation and abundances of protoplanetary disks can more qualitatively be understood from the simulations.

Smaller star forming regions are more likely to produce stars with protoplanetary disks, while in massive star clusters these are quickly destroyed by nearby stars – if protoplanetary disks form in these regains at all.

✗ca. 0.1 light year

✗ca. 0.5 light years

✓ca. 1 light year

✗ca. 10 light years

✗ca. 15 light years

✗True

✓False

✗… as a single star

✗… in massive star-clusters

✓… in the vicinity with other stars

3.8 ALMA Image of a Protoplanetary Disk

This image has been taken with the Atacama Large Millimeter/Submillimeter Array (ALMA) in Chile, and shows a planet-forming, i.e., a protoplanetary disk around the young, Sun-like star TW Hydrae, about 176 light-years away from Earth. The dark, narrow gaps between the broad red rings have low dust densities. Planets that formed from agglomerating dust likely have formed the gaps. Although, alternative interpretations are suggested.

There are gaps in these protoplanetary disks. These gaps are only a few tens of AU away from the central star, which in our solar system are about the distances of e.g., Saturn or Uranus.

The gaps might have formed by planets forming and sweeping up all the material from what are now the gaps. However, alternative processes have been suggested how these gaps might have formed. So it is not yet clear, whether these are really from planets

✗yes

✓no

✗only a few

✗in most gaps, yes

✗Tens of Lightyears

✓Hundreds of Lightyears

✗Thousands of Lightyears

✗Hundred Thousands of Lightyears

✓dust

✗gas

✗planets

✗chondrules

3.9 Formation of the Protoplanetary Disk

The solar system formed when a large molecular cloud of maybe one lightyear in diameter collapsed. It is yet unclear what triggered the collapse. It might have been some statistical fluctuations within the molecular cloud, or the shock wave from a nearby supernova. The latter is support by isotopic evidence in meteorites that indicates the late injection of presolar grains, which might have formed in a supernova. The collapsing molecular cloud started to rotate, flatten, and finally a protoplanetary disk formed. The density in the center increased to the point when H atoms started to fuse to He, which ignited the Sun. Material growth in the protoplanetary disk and accretion over 10-15 orders of magnitudes formed the planets, and opened up gaps in the protoplanetary disk that dissipated within about less then 10 Ma.

The shock front of a nearby supernova might have collided and trigged the collapse. Alternatively, statistical fluctuations in the molecular cloud itself might have triggered its collapse.

The molecular cloud collapsed, started to rotate, and flatten.

✓The isotope composition of meteorites

✗The element composition of meteorites The Ca,Al-rich inclusions (CAIs) in meteorites

✗The CI chondrite with presolar compositions

✗True

✓False

✗True

✓False

3.10 Initial Composition of the Protoplanetary Disk

The initial composition of the protoplanetary disk resembles the composition of the Sun, i.e., the Solar System: most was H and He, only 2% were something else, and are often collectively called ›metals‹. A litte more than a quarter of these metals are rocks, from which the meteorites, asteroids and terrestrial planets formed. The rest are various ices, from which the gas giants formed.

All elements except for H and He.

Rocky material, H2O ice (0.9 wt%), CH4 ice (0.5 wt%), and NH3 ice (0.1 wt%).

✗0.002 wt%

✗0.02 wt%

✗0.2 wt%

✓2 wt%

✗0.06 wt%

✓0.6 wt%

✗6 wt%

✓25 wt%

✗75 wt%

3.11 Material Growth in the Early Solar System

Our Solar System formed from a molecular cloud of about one light year in diameter. Such molecular clouds consist primarily of gas, but also interstellar material. These are sub-µm small grains. Subsequent to the collapse of the molecular cloud, a protoplanetary disk forms around forming Sun, in which the first of 3 discrete growth steps proceeds. The sub-µm small interstellar grains grow to hundreds of µm to mm sized grains, the chondrules. Further µm to tens of µm sized material form and constitute the matrix of the chondrites. This material, however, was certainly also produced on asteroids during alteration processes. In the second growth step, chondrules and matrix agglomerated into m to km sized planetesimals, including minor components such as CAIs and opaque phases. During this time the protoplanetary disk dissolves, and the asteroids orbit the sun on Keplerian paths. In the final third growth step these asteroids collide and aggregate into hundreds of thousands km sized planets. Only these there discrete growth steps were necessary to build planets across 10-15 orders of magnitude from sub-µm small interstellar grains.

1: sub-µm; 2: µm to mm – growth by 3-6 orders of magnitude; 3: m to km – growth by 6-9 orders of magnitude; 4: hundred thousands of km – growth by 4-9 orders of magnitude. In Total: from sub-µm to hundred thousandds of km ca. 10-15 orders of magnitude.

1: interstellare material: 2: chondrules, matrix, CAIs, opaque phases; 3: planetesimals (now called asteroids); 4: planets and moons.

✗ca. 0.1 light year

✗ca. 0.5 light years

✓ca. 1 light year

✗ca. 10 light years

✗ca. 15 light years

✗2

✓3

✗4

✗5

✗6

✗1

✓2

✗3

✗4

3.12 Lifetimes of Protoplanetary Disks

Stars form contemporaneously in stellar clusters. Circumstellar disks are observed around these young stars. The fraction of the near-infrared (Near-IR) excess of a circumstellar disk is a measure of how much of a circumstellar disk is still present around a star. A near-IR excess fraction of 0% means that the circumstellar disk entirely dissipated. It is further possible to determine mean ages for stellar clusters. The negative, linear correlation between near-IR excess fraction and age of stellar clusters can then be interpreted as an almost linear dissipation of the circumstellar disk over time. Even more important, an age can be derived, after which circumstellar disks are fully dissipated, which is around after 6 Ma. This is in accordance with what is inferred from the meteorite record.

x-axis: Mean age of a star-cluster (in Ma), in which stars with circumstellar disks are observed. y-axis: Fraction of the near-infrared (Near-IR) excess of a circumstellar disk (in %).

Circumstellar disks dissipate after around 10 Ma. Chondrules and planetesimals – the basic building blocks in our early Solar System – together with at least Jupiter formed in the first 4-10 Ma of our Solar System. After this time, the initial protoplanetary disk of our Solar System most likely dissipated. Hence, the timescales for protoplanetary disk dissipation are comparable between our Solar System and circumstaller disks observed around other stars. Our Solar System is therefore likely a typical member of planetary system, and a great place to study planetary system formation.

✓about linear

✗about logarithmic

✗about quadratic

✓True

✗False

✓10 Ma

✗20 Ma

✗30 Ma

✗100 Ma