Solar nebula: initial composition and early development

Initial composition

Initially gas and dust were the only constituents of the cold molecular interstellar cloud, the raw material for the solar nebula. Spectroscopic observations of dense interstellar clouds along with modelling indeed suggest that the difference between the local dust/gas ratio and that in other regions of the Galaxy can be attributed entirely to a difference in metallicity (Vuong et al., 2003). According to the solar metallicity (Table 3.2), the bulk-mass dust/gas ratio in the presolar cloud was ∼ 1/100. The gas included highly volatile elements, first of all H2, He, N2, CH4, CO, the noble gases etc. The dust grains, which were mineral or amorphous condensates, varied in composition from ices (H2O, CO, CO2, NH3, C2H6) to refractory grains (e.g. Al2O3, SiC, graphite and diamond). The mean size of the dust grains was generally ∼ 10−4 cm but could vary from ∼ 10−8 cm up to 1 cm (Elmegreen, 1981).

The dust was chemically and isotopically heterogeneous (Section 3.3), implying a number of presolar sources of the raw material. More than 1000 stars might have contributed to the presolar cloud (Flam, 1991).

Early T Tauri stage: high-temperature processing

There are three approaches allowing PT-t parameters (i.e. pressure and temperature conditions varying in the course of time t), which govern the early evolution of matter of the solar nebula, to be estimated.

Introduction: nuclei and their behaviour

Atoms are the smallest units of matter that characterize a chemical element. An atom consists of a positively charged core or nucleus and negatively charged electrons orbiting around the core. In nuclear physics, a host of different particles is known to make up atomic cores, but for the purpose of cosmochemistry and geochemistry the simplified model suffices, in which we consider just two kinds of nuclear particles (nucleons): positively charged protons, p, and neutral neutrons, n. For a neutral atom the number of protons in the core, Z (the atomic number), is equal to the number of electrons around it. As Z determines the electron configuration and therefore the chemical behaviour, a family of atoms of equal Z constitutes a chemical element. Such a family generally includes nuclei with a varying number of neutrons, N. The atomic mass number A = Z + N, the total number of nucleons, then varies accordingly. Atoms of an element that have different values of N (and therefore A) are called isotopes, a term with Greek roots indicating that these different nuclides occupy the same position in the periodic table. The lightest element, hydrogen, includes three isotopes, 1H, 2H (D) and 3H, having 0, 1 and 2 neutrons in the core, respectively. Most elements consist of a larger number of isotopes; therefore the approximately 100 currently known elements include approximately 1000 isotopes.

Many isotopes exist indefinitely, at least in normal conditions, and these are known as stable isotopes, S.

Introduction: processes governing galactic chemical evolution

After the Big Bang, stellar nucleosynthesis became by far the major element-producing factory. During the mass-loss episodes experienced by most stars, newly produced elements are (partially) transferred to the interstellar medium (ISM), which thereby evolves chemically. The element yields and the transfer mechanisms are both dependent on stellar masses and evolution (Chapters 5 and 6).

The interstellar medium initially consisted almost exclusively of primordial H and He, and this low metallicity favoured the accretion of very massive stars, ∼ 30 M⊙ < M < ∼ 300 M⊙. Even though adequate models of these stars have not yet been developed, they are considered to have a very short life. The total mass of the ejected heavy elements as well as their abundance pattern strongly depend on a stellar model, especially on the details of stellar death. In the case of a successful explosion, heavy r-process and Fe-peak elements are abundant in stellar ejecta whereas the inner core is converted into a neutron star. In the other case, only products of explosive burning in the outer shells are ejected; most material falls back onto the core to generate a massive black hole that could play an important role in the subsequent formation of a galaxy. In both cases the amount of material ejected into the interstellar medium could be enough to overcome the metallicity threshold in a stellar neighbourhood for the formation of low-mass stars, ΣAZ ∼ −3; therefore small low-metallicity long-lived stars could have sampled an array of elements generated by one or very few progenitor(s) (Silk and Bouwens, 2001; Bromm and Larson, 2004; Schneider, 2006; Cowan and Sneden, 2006).

Introduction

The major elements in the Earth's atmosphere and hydrosphere, H, N, C and O, are highly reactive and also play major roles in solid-Earth evolution. Thus the noble gases are the most suitable tracers for early atmo-hydrosphere processes.

The sources of non-radiogenic noble gases in the Earth's interior are reasonably well understood: these include the accreting bodies of chondritic composition that delivered “subsolar” noble gases to the Earth–atmosphere system (in abundances similar to those in E-chondrites), and/or mixed materials such as C1 chondrite-type matter bearing Q gases along with grains containing implanted solar noble gases (Chapter 28).

The present-day abundance of noble-gas isotopes in the atmosphere could have developed via (1) fractionation of the initial composition during early loss event(s) and (2) the addition of radiogenic or fissiogenic isotopes, transferred to the atmosphere from the planetary interior. Comparison of the subsolar noble-gas isotope pattern, seen in the Earth's mantle and postulated as the initial one for the atmosphere, with the present-day atmospheric abundances of unradiogenic isotopes reveals an elemental and mass-dependent isotopic fractionation of the atmospheric gases. Therefore adequate models of atmosphere-loss mechanisms must account for these features (Section 20.3).

The inventory of radiogenic noble-gas isotopes, first of all the 129I–129Xe(I) and 136Xe(Pu)–244Pu systematics, indicates a major loss of atmophile elements, even as heavy as xenon, from the Earth's and the martian atmospheres and also presents rough estimates of the relevant time scales (Section 20.2).

In this part of the book the processes of nucleosynthesis and the environments in which they are occurring, and have occurred are sketched out.

To understand the principles of nucleosynthesis, it is important to appreciate the factors that determine the relative stability of different nuclides, and this subject is treated in Chapter 1. The grand scene is introduced in Chapter 2, without too much detail. Chapter 3 deals with data and observations concerning the chemical and isotopic composition of stars, galaxies and the solar system. This follows a broad chronological order, starting with the D/H and He/H ratios that lend support to the hypothesis of Big Bang nucleosynthesis, following through with the most primitive stellar matter and heterogeneities in presolar grains and then focussing on the composition of the solar system. Models and explanations of these data are contained in Chapters 4 to 8, which relate the data to results derived from astrophysical modelling. This helps us to understand first how the chemical elements were and are produced and second how they were scattered in space, to be incorporated in stars and solar systems that formed later.

Introduction: non-chondritic meteorites and their relationships

The two main groups of meteorites other than chondrites are achondrites and iron meteorites (Wasson, 1985). Meteorites from the first group consist essentially of silicate minerals and are distinct from the chondrites both in their compositions and internal rock textures (the intergrowths of minerals), which indicate that they crystallized from a silicate melt or, at least, equilibrated extensively with such a melt. The second group comprises iron meteorites consisting chiefly of metal, frequently with sulphide and occasionally silicate inclusions. They tend to be coarse-grained (up to cm-size grains) and, by etching of a polished surface, the crystalline structure of the metal is made visible in the so-called Widmanstatten structures so often admired in museum collections.

There are ample arguments that both these meteorite types are products of the thermal processing of a chondritic precursor. These are, for example, the identical oxygen-isotopic compositions of some chondrites, achondrites and irons; the complementary chemical compositions of achondrites and iron meteorites: the mineralogy of silicate inclusions in stone-irons; the similar densities of processed asteroids and chondritic meteorites and of the unprocessed clasts of chondritic material often seen in non-chondritic meteorites (Zolensky et al., 1996; Meibom and Clark, 1999).

Oxygen isotopes appear to be a powerful tool for tracing the relations between primitive and processed meteorites. Enstatite chondrites, enstatite achondrites and aubrites have O-isotopic compositions close to the terrestrial fractionation line.

Hydrogen and helium and their special significance

The lightest isotope, hydrogen, with A = 1, is the prime building block for the elements, and spectroscopic measurements show that H is the most abundant element in stars and interstellar clouds in our Galaxy and in the Universe as a whole: nine out of ten atoms are hydrogen. The second stable hydrogen isotope, deuterium (D), with A = 2, is much less abundant: it has a low binding energy per nucleon (Fig. 1.2), and upon collision with baryons and heavier particles it is readily fused to form 3He or He. As nuclear processes in stars thus tend to destroy D, this nuclide must have been generated in another process, i.e. in the earliest prestellar nucleosynthesis.

This hypothesis can be tested by measuring the prestellar D/H ratio of galaxies. This can be achieved by the spectroscopy of interstellar clouds lying along the line of sight to a remote very bright object. Some of these clouds represent nearly virgin samples of prestellar cosmic material with ∼ 1000 times lower metallicity than that in the solar system (the metallicity of a system is a measure of the abundance of elements heavier than helium). High-resolution spectroscopic measurements of the H and D abundances in these clouds gave D/H = (30 ± 2) × 10−6 (Schramm and Turner, 1998).

Major geotectonic units: the plates

Let us imagine that we observe the Earth from a remote point, for example from a satellite, and that the atmosphere and oceans are transparent and therefore we can see the whole solid surface of the planet (Fig. 23.1(a), (b)). The most intriguing features of this surface are long narrow zones or cuts, characterized by enhanced seismicity and heat flow and also by intense volcanic activity: volcanic exhalations are easily seen from our satellite. These narrow zones outline much “quieter” (aseismic and non-volcanic) regions, which cover most of the surface. Within the plate-tectonic concept the stable regions of the Earth surface are termed lithospheric plates and the active zones are plate boundaries (e.g. van Andel, 1992). Using the satellite-positioning system, we can observe that the plates move with relative velocities between 1 and 10 cm yr−1. The active zones separate plates that are moving in different directions. As discussed in Section 17.5, the largest terrestrial reservoir, the mantle, consists of solid silicate material convecting at the rate of a few cm yr−1. This is similar to the observed plate velocities, and the similarity implies that there is some relationship between the motion of the plates and the convection of the underlying mantle: the plates form the lid of the mantle.

The plates vary greatly in size (and therefore in the lengths of their boundaries), from ∼ 1000 km or less (Cocos, Corda, Fig. 23.1(b)) to more than 10 000 kilometres (Pacific, Antarctic, Eurasia).

Introduction

In this chapter we discuss the present-day isotopic characteristics of the principal terrestrial reservoirs and the evolutionary trends that could have led to their diversity, starting from a common chondrite-type origin. Ultimately, these data allow the modelling of Earth's evolution. The main reservoirs considered in this book are the core (COR, Chapter 18), the core–mantle transition zone (D′′ or DDP, Chapter 19), the depleted MORB-source mantle (DMM), the continental crust (CCR) and the atmo-hydrosphere (AHS). Of these the latter three can be considered to be the Earth's accessible reservoirs (EARs), as they can be directly or indirectly sampled, and these are the main reservoirs discussed below. The two silicate reservoirs, DMM and CCR, constitute the major part of the bulk silicate Earth (BSE). This is generally considered to have chondrite-like relative abundances of the refractory lithophile elements, and hence the name CHUR (chondritic uniform reservoir) is also used.

The complexity of the terrestrial transport processes discussed in Chapters 24–26 in many cases prevents the reconstruction of remote source domains using the chemical compositions of the sedimentary or magmatic rocks that resulted. However, isotope ratios preserve a unique record of the chemical composition and evolution of the sources, and reading this record is the major aim of isotope geochemistry and geochronology. These are multifarious branches of Earth sciences, and in this chapter we mainly (but not exclusively) concentrate on how radiogenic isotope systematics shed light on both the composition and evolution of major terrestrial reservoirs.

Introduction

From the discussion of the Big Bang nucleosynthesis in the previous chapter it follows that H and He, but not the heavier elements, could have been produced in this very early event. Since the classical publication by Burbidge et al. (1957) there has been no doubt that these heavier elements are produced in stellar nucleosynthesis, and the real problem is to discover how similar model stars and the nuclear processes in them are to reality.

After the collapse of a cold molecular cloud and the formation of a star, nuclear burning becomes the major source of energy in a star and also the process producing heavier elements from lighter ones. The production of heavy elements with A > 60 is controlled by neutron fluxes, which can be slow (Section 5.4) or rapid (Section 6.1), and also by several associated nuclear processes. The elements used by a star at its formation as well as newly produced elements are (partially) returned to the interstellar medium by stellar winds or explosions and mix with the matter already there. Another part of the stellar matter could be held in stellar remnants for a long time.

In this chapter we describe nucleosynthesis processes in stars up to ∼ 4 solar masses. These include H and He burning, which produces C, N and O in abundance and small amounts of the elements up to Si (Section 5.3), as well as the generation of elements heavier than Fe by the slow neutron capture s-process.

Introduction: siderophile elements in the silicate mantle and light elements in the core

Siderophile-element abundance pattern in the Earth's mantle

As discussed above, it is an important feature of the silicate Earth that its refractory-lithophile-element abundances are like those of chondrites, thus documenting a relation between the Earth and chondritic matter. In contrast, the siderophile elements are depleted in the silicate Earth, almost certainly because of their segregation into the metallic core. Their relative depletion pattern provides important clues on the core-segregation process (Wänke et al., 1984; Ringwood, 1984; Jones and Drake, 1986; Newsom, 1990; O'Neill, 1991a, b).

Transitional, moderately and highly siderophile elements constitute three groups that are progressively depleted, thus pointing to metal–silicate fractionation having played an important role in core formation (Fig. 18.1). However, (1) within each group the observed chondrite-normalized abundance pattern is almost flat and contrasts with the variable abundances predicted from metal–silicate equilibrium fractionation. Also (2) for strongly siderophile and for most moderately siderophile elements, the observed abundances greatly exceed those predicted from equilibrium fractionation, while for W and the transition elements they are lower.

Becker et al. (2006) recently reported new data for highly siderophile element (HSE) concentrations in the mantle. Their best estimate, [Ir] = 3.5 ppb, is indistinguishable from the previous average value suggested by Morgan et al. (2001), and the C1-normalized HSE pattern is flat except for a small enhancement of [Ru] (by a factor 1.4) that is not revealed by previous data sets.