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Interstellar Matter
Donald G. York, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
V Properties of the Interstellar Clouds
Given the many possibilities for detection of radiation emitted or modified by interstellar gas, astronomers have pieced together a picture of the interstellar medium. In many cases, the details of the physical processes are vague. In most cases, detailed three-dimensional models cannot be constructed or are very model dependent. On the other hand, in some cases, sufficient knowledge exists to learn about other areas of astrophysics from direct observations. The example of the cosmic-ray distribution has already been given. Others are mentioned subsequently.
Interstellar clouds are complex aggregates of gas at certain velocities, typically moving at ±6 to 20 km/sec with respect to galactic rotation, itself ∼250 km/sec over most of the galaxy. Each cloud is a complex mixture of a volume of gas in a near pressure equilibrium and of isolated regions affected by transient pressure shocks or radiation pulses from star formation or from supernova explosions. Clouds are visible in optical, UV, or X-ray emission (or continuum scattering) when they happen to be close to hot stars and are otherwise detectable in absorption (molecules or atoms) or emission from low-lying excited levels (0.01 eV) or from thermal emission of the grains in the clouds.
V.A Temperatures
Molecular clouds are as cold as 10 K. Diffuse clouds are typically 100 K. HII regions have T ∼ 8000 K, depending on abundances of heavy elements that provide the cooling radiation. Low-column-density regions with 10,000 < T < 400,000 K are seen directly, presumably the result of heating at the cloud edges from shocks, X rays, and thermal conduction. Isolated regions with T > 106 K are seen near sites of supernova explosions.
V.B Densities
Densities, as determined from direct observation of excited states of atoms and molecules, are generally inversely proportional to temperature, implying the existence of a quasi-equilibrium state between the various phases of the medium. The effects of sources of disequilibrium in almost all cases last ≲107 years, or less than 1/10 of a galactic rotation time, itself 1/10 of the age of the sun. The product nT(cm−3K) is ∼3000 to within a factor of three where good measurements exist. Thus, the molecular (dark) clouds have n > 102 cm−3, while in diffuse clouds, n < 102 cm−3. Higher densities (up to 105 cm−3 occur in disequilibrium situations such as star-forming regions inside dense clouds and in HII regions.
V.C Abundances: Gas and Solid Phases
By measuring column densities of various elements with respect to hydrogen, making ionization corrections as necessary, abundances of elements in interstellar diffuse clouds can be determined. Normally, the abundances are compared with those determined in the sun.
Different degrees of depletion are found for different elements. Oxygen, nitrogen, carbon, magnesium, sulfur, argon, and zinc show less than a factor of two depletion. Silicon, aluminum, calcium, iron, nickel, manganese, and titanium show depletions of factors of 5–1000. Correlations of depletion with first ionization potential or with the condensation temperature (the temperature of a gas in thermal equilibrium at which gas-phase atoms condense into solid minerals), have been suggested, but none of these scenarios actually fits the data in detail.
The pattern of depletion suggests no connection with nucleosynthetic processes. Those elements that are depleted are presumed to be locked into solid material, called grains. Such particles are required by many other observations attributed to interstellar gas, as discussed earlier. In principle, the unknown makeup of the grains can be determined in detail by noting exactly what is missing in the gas phase. However, since there must be varying sizes and probably types of grains and since the most heavily depleted elements do not constitute enough mass to explain the total extinction per H atom, most of the grains by mass must be in carbon and/or oxygen. Establishing the exact mass of the grains amounts to measuring the depletions of C and O accurately. Although these measurements can now be made with the Hubble Space Telescope, problems of interpretation of deletions remain.
The grain structure (amorphous or crystalline) is not known. There are unidentified broad absorption features, called diffuse interstellar bands, that have been attributed to impurities in crystalline grains. However, these features may be caused by large molecules. It has been argued that even if grains are formed as crystalline structures, bombardment by cosmic rays would lead to amorphous structures over the life of the galaxy.
Theories of grain formation are uncertain. A general scenario is that they are produced in expanding atmospheres of cool supergiants, perhaps in very small “seed” form. They may then acquire a surface layer, called a mantle, probably in the form of water ice and solid CH4, NH3, etc. This growth must occur in cold, dense clouds. The detailed process and the distribution of atoms between minerals and molecules in solid phase are unknown.
V.D Evolution
Interstellar clouds can be large, up to 106 solar masses, and are often said to be the most massive entities in the galaxy. In this form, they may have a lifetime of more than 108 years. They are presumably dissipated as a result of pressure from stars formed within the clouds. Over the lifetime of the galaxy, interstellar clouds eventually turn into stars, the diffuse clouds being the residue from the star formation process. Growth of new molecular clouds from diffuse material is poorly understood. Various processes to compress the clouds have been suggested, including a spiral density wave and supernova blast waves. No one mechanism seems to dominate and several may be applicable. However, the existence of galaxies with up to 50% of their mass in gas and dust and of others with less than 1% of their mass in interstellar material leads to the inference that diffuse material and molecular clouds are eventually converted into stars.
Millimeter Astronomy
Jeffrey G. Mangum, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
III.A.2 Measurements of Physical Conditions
The study of molecular cloud stability and evolution leads naturally to studies of the physical and chemical evolution of the star formation process. Fundamental to this study of the star formation process is the characterization of the physical conditions in the gas and dust comprising these regions. For the gas, volume density n (cm−3), kinetic temperature TK (K), chemical composition X, turbulent motion Δ υ (km sec−1), and magnetic field strength B (Gauss) are fundamental physical quantities. For the dust, the dust temperature Td (K), dust volume density nd (cm−3), and dust opacity κ describe the physical conditions representative of the dust component of a molecular cloud. Note that all of these quantities are dependent upon time and position.
Most of the material in molecular clouds is in the form of H2, which owing to its lack of a permanent dipole moment has no easily observable rotational transitions. It can be observed through rovibrational and fluorescent transitions, but only within environments which are very specific, such as shocks and regions containing high levels of ultraviolet emission. Therefore, the principal component of molecular clouds is effectively unmeasurable. This fact forces astronomers to use trace constituents, other molecules and dust, to measure the physical conditions in molecular clouds.
III.A.2.a Molecular emission as a tracer of physical conditions
The primary constituent of molecular clouds, H2, is also the main collision partner with other molecular inhabitants of these regions. These collisions lead to the excitation of rotational transitions in a variety of molecules, many of which emit at observable millimeter wavelengths. The most abundant molecule after H2 is carbon monoxide (12C16O, usually simply written CO). It was the first molecule discovered at millimeter wavelengths by Wilson, Penzias, and Jefferts in 1970 using the National Radio Astronomy Observatory 36 ft (now 12 m) millimeter telescope located on Kitt Peak, Arizona. It has been used extensively as a probe of the volume density and kinetic temperature in molecular clouds through measurements of its lowest two rotational transitions at 115.271 and 230.538 GHz. CO has proven to be a very good tracer of the global physical conditions in molecular clouds, but for more compact regions with a larger number of particles along the line of sight [referred to as the column density (N) of a particular molecule], it loses its sensitivity to the bulk of the gas as the opacity in the measured transitions rises. Fortunately, there are other less abundant molecular tracers, including isotopomers (isotopic variants) of CO, such as 13CO, C18O, and 13C18O, which prove to be better probes of these high column density environments.
There are a wide variety of molecules that can be used as tracers of the volume density and kinetic temperature in molecular clouds. The choice of molecular probe depends upon what environment one wishes to study. For example, to measure the physical conditions in the dense cores of molecular clouds, it is best to choose a molecular tracer that is particularly sensitive to the prevalent conditions in this environment. A useful guide used to calculate the sensitivity of a transition to volume density is the critical density ncrit, which is the volume density required to collisionally excite a transition assuming optically thin conditions,
ncrit=AijCij=64π4νij33hc3giCij|μ⇀ji|2=64π4νij33hc3giCijSμ2,
where Aij is the spontaneous emission (Einstein A) coefficient for level i, Cij is the collisional deexcitation rate per molecule in level i, gi is the upper state degeneracy, |μ⇀ji| is the dipole moment matrix element for the transition, S is the line strength for the transition, μ is the dipole moment for the molecule, and the other terms have their usual meanings. Critical densities for common molecules such as CS, HCN, and H2CO are in the range 104−8 cm−3 for a kinetic temperature of 10 K.
Therefore, a simple detection of a transition from one of these molecules implies the existence of dense gas. A second consideration is to choose molecules that allow one to derive accurate measures of the volume density and kinetic temperature in a molecular cloud. Since the collisional excitation of molecular transitions is dependent upon the coupled effects of volume density and kinetic temperature, it is often necessary to use molecules that allow one to decouple these effects. The ability to decouple these physical effects depends upon the properties of the molecular structure. There are three basic types of molecules in this regard; linear, symmetric rotor, and asymmetric rotor. Figures 5–7 show the energy level structure for these three types of molecules. As can be seen from Fig. 5, linear molecules have one ladder of energy levels, the transitions between which are excited by the coupled effects of volume density and kinetic temperature. In general, linear molecules are used to derive the volume density in a molecular cloud by assuming a kinetic temperature or by using a calculation of the kinetic temperature based on measurements of the transitions from another molecule. The energy level structures for the symmetric and asymmetric rotor molecules shown in Figs. 6 and 7 indicate a more complex structure. Like linear molecules, the strengths of transitions within a given ladder (designated by the “K” rotational quantum numbers) are dependent upon the coupled effects of volume density and kinetic temperature. A comparison of the strengths of transitions from the same J levels but from different K ladders, though, is dependent only on the kinetic temperature, thus making it possible to derive a decoupled measurement of the kinetic temperature in a molecular cloud. In general, then, symmetric and asymmetric molecules have molecular level properties that, when the appropriate transitions are compared, allow decoupled measurements of the volume density and kinetic temperature in a molecular cloud. Linear molecules do not possess these decoupling properties, thus requiring an independent measurement of either the volume density or kinetic temperature.
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FIGURE 5. Rotational energy level diagram for HC3N, a typical linear molecule. The rotational quantum number “J” and associated energy above the ground state are indicated for each level.
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FIGURE 6. Rotational energy level diagram for CH3CN, a typical symmetric rotor molecule. The rotational quantum numbers J and K, along with the associated energy above the ground state, are indicated for each level.
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FIGURE 7. Rotational energy level diagram for H2CO, a typical asymmetric rotor molecule. The rotational quantum numbers J, K−1, and K+1, along with the associated energy above the ground state, are indicated for each level.
By comparing the measured intensities of a variety of transitions from a given molecule with molecular line intensity predictions from a molecular cloud model, estimates of the volume density and kinetic temperature within a molecular cloud can be made. In general, these estimates reveal that the volume densities in the regions where stars form within molecular clouds exceed 104 cm−3, while the kinetic temperatures range from 10 to 300 K. These models also indicate that many molecular cores possess density gradients, suggestive of a structure that could evolve into a collapsing protostar.
III.A.2.b Kinematics
The shape of a spectral line is determined by the radial velocity structure along the line of sight through a molecular cloud. The measured widths of spectral lines are generally larger than the thermal width, indicating that the velocity fields within molecular clouds are dominated by Doppler broadening owing to turbulence:
Δυ=υtherm+υturb=22ln2kTKm+υturb
where Δ υ is the full width at half maximum of the spectral line, TK is the kinetic temperature of the gas, and m is the mass of the particles which make up the molecular cloud (principally molecular hydrogen). Unfortunately, a detailed derivation of the spectral line shape has proven elusive, owing to the effects of kinetic temperature and volume density gradients, spatial structure, and radiative transfer effects within the molecular cloud. Measurements of the line center velocity as a function of position over a molecular cloud do indicate that, in general and on large scales, they are neither collapsing nor rotating. This is not the case on small scales. Evidence for rotation and collapse of molecular cloud cores on 0.1 pc scales have yielded interesting clues to the details of the star formation process. Measurements of cloud core rotation indicate that in magnitude it is only 2% of the gravitational potential energy before collapse, making it relatively unimportant to the overall dynamics of a molecular cloud core.
The physical nature of the turbulent component of a spectral line in a molecular cloud is currently a source of considerable debate. Physical processes that have been suggested as sources of the turbulence in molecular clouds are expanding HII regions, supernova remnants, cloud–cloud collisions, galactic differential rotation, and stellar winds. Unfortunately, for all of these processes there are theoretical problems with coupling the energy produced into turbulence.
III.A.2.c Magnetic fields
An understanding of the magnetic field properties of molecular clouds is an important aspect of the overall physical understanding of molecular clouds given their apparent role in providing dynamical support in these environments. There are three methods that have been used to measure the magnetic field strength and direction within molecular clouds: atomic and molecular Zeeman effect splitting, which tell us about the line-of-sight magnetic field component (BZ); polarization of the emission from dust grains, which gives us information about the component of the magnetic fields perpendicular to the line-of-sight (B⊥); and measurements of spectral line emission polarization, which also tells us about B⊥. Measurements of the Zeeman splitting in atoms and molecules have concentrated on studies of HI, OH, and CN, yielding typical values for BZ of 10–20 μG within regions with volume densities of approximately 103 cm−3. At higher volume densities, magnetic fields as large as 700 μG have been measured. Millimeter dust continuum emission polarization levels of at most a few percent have shown that the magnetic field within the high volume density (n ≥ 106 cm−3) cores of molecular clouds is perpendicular to the major axis of the high density structures (such as disks) and parallel to the outflows associated with these objects. Although the possibility of there being measurable levels of polarization of thermal millimeter spectral line emission has been known for years, it has only recently been detected. The percentage of measured polarized emission is equivalent to that detected through millimeter continuum polarimetry.
Turbulence and magnetic fields cannot simply be applied as independent solutions to the problem of cloud support since turbulence should tangle magnetic fields, thus reducing their effectiveness as a source of support. The theory of magnetohydrodynamic turbulence within molecular clouds has shown that magnetic fields should slow the decay of turbulent motions if these motions are less than the propagation speed along the magnetic field lines, referred to as the Alfvén velocity. However, the stability of magnetic support in the presence of turbulence has been called into question, and the interplay between cloud stability and dynamics drives our understanding of the importance of magnetic fields as a means of molecular cloud support. Future measurements of the magnetic field and direction in molecular clouds should clarify their influence on the overall dynamics and evolution of these regions.
Astrochemistry
Steven N. Shore, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
V.B.2 Ionization
Ionization in the densest parts of molecular clouds depends on the penetration of cosmic rays and UV radiation, as well as the presence of shocks generated by such processes as cloud–cloud collisions and internal star formation. In order to probe the electron density in the clouds, it is important to be able to account for the presence of complex polyatomic molecules, whose formation requires ion gas-phase reactions.
The rate for cosmic-ray (CR) ionization, ζCR, is about (3 ± 1) × 10−17 s−1. This is an integral over the collisional ionization cross section for low (MeV)-energy CR protons, but it is approximately a constant for most of the species of interest. An obstacle in our understanding of the detailed structure of molecular clouds is our ignorance of the precise specification of this rate. The low-energy end of the cosmic-ray spectrum is difficult to determine empirically from terrestrial observation, because these particles propagate diffusively through the interplanetary medium, scattering off of turbulence in the solar wind; their spectrum cannot be observed directly, even with in situ measurements from the Voyager and Ulysses spacecraft, and must be inferred from models for their motion through the heliosphere. The more easily observed cosmic ray protons and electrons, in the GeV and higher range, have little or no effect on the ionization of the interstellar medium because of the small interaction cross sections for atoms at such high energies.
In molecular clouds, atomic species with ionization energies greater than 13.6 eV must be predominantly neutral because of the shielding effects of neutral hydrogen. It is mainly the heavier elements, such as C, N, and O, which are observed in the peripheral portions of the clouds to be in the partially ionized state. For circumstellar envelopes, cosmic rays lose out to photo processes and the chemistry is mediated by the input of stellar photospheric radiation (in the hotter stars and in novae and supernovae) and from the diffuse interstellar radiation field.
The basic equations for two body interactions can be written in the form
(24)dNidt=∑j,k≠iKijkNjNk−∑jKij′NiNj,
where Kijk is the formation rate for the ith molecular species, while K′ij is the destruction rate for the molecule. The inclusion of UV photo processes is accomplished by the photodissociation rate:
(25)Rpd=∫ν0∞κνFνe−τνdνhν,
where Fν is the incident photon flux, τν is the opacity of the ambient medium (presumed to be from dust), κν is the continuous absorption coefficient for the dissociative continuum, and the dissociation energy is hν0.
An aspect in which circumstellar environments differ from interstellar is the net mass advection through the medium. Abundances become time dependent—and hence space dependent—in the envelope, due both to the implicit time dependence of the reactions and to the transport of matter through different radii via stellar wind flow. The atomic abundances are fixed at stellar photosphere, rather than having to be assumed for some mixture of physical parameters of temperature and pressure as they must for molecular clouds. It is then essentially an initial-value problem to compute the abundances which will be a function of radius in the envelope. For a steady-state wind, the abundances become strictly a function of radius. Also, unlike a molecular cloud, the density profile of the envelope is specified from the assumption of steady mass loss at the terminal velocity for the wind, so thatρ(r)=M./(4πr2υ∞), where M. is the mass loss rate and υ∞ is the terminal velocity of the wind.
An interesting aspect of stellar envelopes is that they may have two different sources of UV radiation, internal and external. Work on the envelope of two extreme, low-temperature, evolved supergiants, IRC + 10216 and α Ori, showed that the outer limit of the molecular envelope is determined by the DIRF, which destroys the outermost molecular species by photodissociation, while the inner boundary is set by both the temperature and UV emission from the stellar chromospheres. In this respect, since the dynamics can be probed in exquisite detail for several of the nearer supergiants through molecular observations, and since the input abundances are known and atomic in nature, it is possible to use these stars as very well-conditioned laboratories for the study of the same processes which must be involved in at least some aspects of molecular cloud chemistry. For the densest envelopes, which are completely optically thick and hence very similar to molecular clouds, cosmic rays are significant in governing the ion fractions but can be neglected in thin envelopes (low mass loss rates).
The Hadean and Archaean Atmosphere and the Oldest Records of Life as Micro- or Chemofossils
Akio Makishima, in Origins of the Earth, Moon, and Life, 2017
8.2 The Perspective of Atmospheric Evolution From the Hadean to the Archean Earth
As the sun formed from its molecular cloud, it was accompanied by disk material that consisted of gas and small dust particles. Over several tens of millions of years, these dust particles formed the planets. This process occurred in several stages in the terrestrial planet zone, including moon-forming impacts on the proto-Earth (Canup and Asphaug, 2001).
A Hadean atmosphere containing N2 and CO2 and a Hadean ocean containing H2O seems to have formed as a natural consequence of planetary accretion in the terrestrial planet region. The atmosphere which had weak reducing potential with relatively high partial pressure of CO2 should have formed (Holland, 1984; Walker, 1985; Kasting, 1993; Ferus et al., 2015; Furukawa et al., 2015). In this condition the important biological precursor compounds for life were synthesized. It should be noted that high partial pressure of CO2 in the early Hadean atmosphere can be presumed. Such gas is like intestinal gas, so it is reasonable to suppose that life may have fermented.
Atmospheric O2 levels rose naturally and gradually, but not immediately, occurring by photosynthesis and organic carbon burial. At the same time, the concentrations of CO2 and other greenhouse gases did not compensate for the brightening sun. The Earth’s relatively stable climate was a result of the negative feedback between atmospheric CO2, surface temperature, and the weathering rate of silicate rocks (Kasting, 1993).
This atmospheric evolution implies that Earth is not a unique planet. If planets exist around other stars, some of them could reside in orbits where the illumination is similar to that received by the Earth. Planetary climates are buffered by the carbonate–silicate cycle. Therefore, the habitable zone around late F to mid K stars (see the top scale of Fig. Box 1.4; the Hertzsprung–Russel diagram) may be wider, and other habitable planets may exist. If the origin of life was not a fortuitous event, many of these planets could be inhabited and on some, intelligent life may even have evolved. Both of these speculations can be tested: the first by spectroscopic investigations from large, space-based telescopes; and the second by monitoring microwave and radio emissions from space (Kasting, 1993).
Morbidelli et al. (2000) suggested that the most plausible sources of the water accreted by the Earth were in the outer asteroid belt, in the giant planet regions, and in the Kuiper Belt. It is plausible that the Earth accreted water from the early phases when the solar nebula was still present to the late stages of gas-free scattered planetesimals. Asteroids and comets from the Jupiter-Saturn region were the first water deliverers, when the Earth was less than half its present mass. The bulk water presently on the Earth was carried by a few planetary embryos, originally formed in the outer asteroid belt and accreted by the Earth at the final stage of its formation. (see Fig. 2.10: The blue planetesimals are water rich, which could be a source of water for Earth).
Finally, a late veneer (this could be the same as or different from the late veneer for the highly siderophile elements discussed in Chapter 4), accounting for at most 10% of the present water mass, occurred due to comets from the Uranus-Neptune region and from the Kuiper Belt. The net result of accretion from these several reservoirs is that the D:H ratio of water on Earth is essentially the typical water condensed in the outer asteroid belt. This is in agreement with the observation that the D:H ratio in the oceans is very close to the mean value of the D:H ratio of water inclusions in carbonaceous chondrites.
Stars, Massive
Steven N. Shore, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
IV.B Induced Star Formation
Should the star formation begin within a molecular cloud, the winds and H II regions can either destroy the cloud by heating it up through radiative and mechanical processes or they can break free of the cloud. This latter type of star formation event is called a champagne flow, analogous to that commonly observed at weddings. The sudden release of pressure by the breakout of the shock from its environment causes a rapid outflow of material from the hot H II region surrounding the OB stars. This flow further drives a shock into the molecular cloud via momentum conservation and compresses the already dense material of the cloud core. Such an event may initiate collapse from gravitational instability, at least locally to the OB stars. Thus, the massive stars are not only capable of destoying the cloud environments in which they have formed, they can also serve as agents for propagating star formation through the cloud.
The formation of the massive stars is well traced by the radio and infrared flux emitted by the H II regions that surround them. Even the densest parts of molecular clouds are not optically thick longward of about 60 μm, so that far infrared (FIR) and radio photons freely escape the cloud. The argument that permits these to determine the rate of massive star formation is as follows. Every massive star that forms an H II region will be surrounded by a thermal, radio-emitting plasma with a temperature of about 104 K. The rate of radio emission depends on the rate of ionization, which in turn depends on the luminosity of the central stars. This in turn depends on the mass of the stars. Each radio photon is associated with the ionized medium, while each recombination eventually leads to a Lyα photon through radiative cascades. These photons, which are trapped in the optically thick H II regions, eventually collide with dust grains in the cloud and the H II region and are absorbed. The dust reradiates this energy at equilibrium in the FIR. Thus, there is an expected correlation between LFIR, the far infrared luminosity, Lradio, and ψ, the star formation rate per unit mass.
Such a correlation is observed in regions of active star formation in the Galaxy and nearby galaxies, although it is still not completely certain what the implied rates of star formation mean. For some galaxies, called starburst galaxies, the rate of massive star formation can reach more than 103 M⊙ yr−1. This enormous rate cannot continue for very long before molecular cloud material is consumed, but the triggering mechanism is not currently understood. Such systems may be the result of collisions between galaxies and are often seen in galactic systems undergoing merger.
UV astronomy and the investigation of the origin of life
Ana I. Gómez de Castro, Ada Canet, in Ultraviolet Astronomy and the Quest for the Origin of Life, 2021
2.3 UV spectral energy distribution and the photoelectric yield
In diffuse environments such as the envelopes of the molecular clouds, the atmosphere of young planetary disks or the interplanetary medium, UV photons are absorbed by materials in the dust grains (graphites, silicates and ferrites, mainly) yielding to the emission of photoelectrons that heat and ionize the environment and accelerate some chemical processes. Only UV radiation, with energy above ∼6 eV, is able to produce a significant photoelectric yield.
The heating rate depends on the density and energy of the ejected photoelectrons which, in turn, depends on the dust grains size (Watson, 1972), composition and charge state as well as on the spectral energy distribution (SED) of the radiation feeding the photoelectric flow (i.e. Pedersen and Gómez de Castro, 2011). The photoelectric yields are displayed in Fig. 2.4 (after Weingartner and Draine, 2001) for silicates and carbonaceous particles. Photons with energies above some 6 eV produce a significant yield from carbonaceous particles while the threshold goes up to 8 eV for silicates.
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Figure 2.4. Photoelectric yield of silicate and carbonaceous particles in space. The various curves correspond to different particle size. In general, the smallest the size the larger the yield because of the finite electron escape length.
After Weintgarner & Draine (2001).
The effectiveness of the process depends on the kinetic energy acquired by the photoelectrons to overcome the Coloumb field of the charged dust grain. Hence, there is a strong dependency between the final charge of the grains and the SED of the UV spectrum irradiating them. This effect shows in the temporal evolution of the charge of the grains (or charging profile) when submitted to a specific UV spectrum, as shown in Fig. 2.5. Two very different spectra are selected to illustrate this process: the very hot, but soft SED of an O-type star and the emission lines spectrum of a pre-main sequence (PMS) solar-like star. Above 6 eV, PMS stars only radiate significantly in very narrow bands associated to the main spectral lines at: ∼8 eV (the He II at 164 nm and CIV at 1550 nm lines), 8.9 eV (Si IV), 9.3 eV (CII), 9.5 eV (OI) and, of course, the Lyα hydrogen line at 10.2 eV. This is the radiation affecting the dust grains in young planetary disks. Dust in the HII regions and the bright reflection nebula are however, charged by the smooth UV spectrum of hot stars. Some examples are the B33 (Horsehead) nebula irradiated by a cluster of early B-type late O-type stars (Caballero, 2007) and the Rosette nebula by the cluster NGC 2244 of B8-9 type stars and 7 O4–O9 stars (Martins et al., 2012). As shown in the bottom panel of Fig. 2.6, dust charging profiles are step-wise in young planetary disks instead of the smooth rising curves from HII regions which are, otherwise consistent with the experimental observation (Sickafoose et al., 2000).
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Figure 2.5. Top: STIS spectra of the Sun (in red, from Meftah et al., 2020), the PMS binary AK Sco (blue, from Gómez de Castro et al., 2020b) and the O-type star HD46223 (yellow, data from the Hubble archive). AK Sco and HD46223 fluxes are scaled so the integrated flux is the same as the solar flux received from the Earth at 1 AU in the STIS wavelength range. Bottom: Charging profile of a silicate dust grain (a = 30 nm) submitted to the radiation fields of an O-type and a PMS star (Pedersen & Gómez de Castro, 2011).
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Figure 2.6. Annular bright field (ABF) images of a Highly Ordered Pyrolytic Graphite (HOPG) sample after irradiation with a Deuterium UV lamp. The high-resolution scanning transmission electron microscopy (STEM) images resolve the separation between the individual graphene layers in the HOPG. The image evidences the dislocation of the superficial graphitic layers after irradiation and the formation of bumps. Electron Energy Loss spectroscopy (EELs) of the bumps show that the sp2 to sp3 hybridation transition is a surface effect that is maximum at the irradiated surface. High sensitivity X-ray diffraction experiments confirm the formation of diamond within the bumps. On the right, the separation between the graphene layers is plotted. There is a transitional area where separation between the layers increases from 3.53 Å to 3.87 Å.
After Gómez de Castro et al., (2020a).
As interstellar and interplanetary dust grains are not in isolation, the final grain charge will also depend on the temperature of the surrounding gas that acts as a regulator since dust grain’s charge is modified in the gas-dust collisions (Spitzer, 1978).
Formation of the Proto-Earth in the Solar Nebula
Akio Makishima, in Origins of the Earth, Moon, and Life, 2017
2.1 Evolution of Molecular Clouds to the Solar Nebula
The evolution of an early star and its surrounding molecular cloud (nebula) are controlled by four components: (1) the centrifugal force; (2) mass ejection (mass and angular momentum loss as ejection jet); (3) mass accretion (mass gain from the disk surrounding the star); (4) a magnetic field (the mass ejection and the accretion follow the magnetic field). These four components are schematically depicted in Fig. 2.1 (Montmerle et al., 2006).
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Figure 2.1. The four components for evolution of the early star and surrounding disk (each component is not in scale).
The star has two strong magnetic fields (B). One magnetic field extends vertically along the rotation axis of the star. Another magnetic field stretches horizontally. In the accretion-ejection model, these two magnetic fields meet along the X-ring (or at the X-point if plotted two dimensionally). The accreting materials from the horizontal direction are ejected at the X-ring.
Our solar system is considered to have been formed from cosmic gas clouds as shown in Fig. 2.2. The Fig. 2.2 was obtained by the Hubble Space Telescope (HST) (see Box 2.1). These gas clouds are made of hydrogen, helium, and heavy elements that are the remnants of a supernova (see Chapter 1). The shockwave from a supernova produced inhomogeneity of the gas cloud near the location where our Sun would appear. The gas cloud about 3 light years began to gather and form a dense gas disk (103 ∼ 104 AU; AU is the present distance between the Sun and Earth). At the center of the gas disk was a huge gas ball (protostar) mainly composed of hydrogen.
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Figure 2.2. The close-up of cosmic clouds and stellar winds at the Orion Nebula.
Image credit: NASA, European Space Agency, and the Hubble Heritage Team www.nasa.gov/sites/default/files/images/724329main_hubble_feature_full_full.jpg.
Box 2.1
The Hubble Space Telescope (HST)
The HST was launched into Earth orbit in 1990 (see Fig. Box 2.1A and B). Its 2.4-m mirror and main instruments can observe the near ultraviolet, visible, and near infrared spectra. The telescope was named after astronomer Edwin Hubble. The HST can take high-resolution images without atmospheric distortion and the background light of the atmosphere. The HST was built by the National Aeronautics and Space Administration (NASA, the U.S. space agency), with contributions from the European Space Agency.
The various stellar phases with timescales and sizes are shown in Fig. 2.3 after Feigelson and Montmerle (1999). These models were established due to the advancement of observation techniques using infrared and X-ray telescopes. The initial stage, which is previously explained, is called the infalling protostar stage (see the top figure of Fig. 2.3). The accretion and ejection of the molecular gases occur. Although the gas ball is very large, it cannot be seen from outside because of the surrounding clouds. This stage lasts fewer than 104 years.
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Figure 2.3. Various stellar phases with timescales and sizes (AU = Sun-Earth distance = 1.5 × 108 km; each stage cartoon is not in scale).
The next stage is the evolved protostar stage. The accretion and ejection of the gases are still vigorously occurring. However, the size of the gas ball becomes 10 times smaller. This stage lasts for 0.1 Myr. The gas disk also glows with accretion.
Then the classic T Tauri stage begins. This stage is also the ejection stage. The gas clouds are vigorously ejected, and the ejection jet is clearly observed. This stage lasts for 1–10 Myr, and the thick disk shrinks to about 100 AU. In this stage, the planetary system would be formed.
The next stage is the weak-lined T Tauri stage. “Weak-lined” means that a “weak-Hα” line of ionized hydrogen is observed in the spectrum of the (proto)star.
Finally, the clouds dissipate and the Sun can be seen. The Sun becomes a main sequence star in the H-R diagram (see Fig. Box 1.4).
The evolution model described here is for the evolution of the Sun (star) based on actual observation of stars and theories. However, when the masses of the presolar nebula or the planetary system are too small compared to the Sun, the evolution of the presolar nebula requires different views and theories.
Planets, Asteriods, Comets and The Solar System
T.H. Burbine, in Treatise on Geochemistry (Second Edition), 2014
2.14.7.3 Formation of Material in the Solar Nebula
Mineralogical and isotopic studies of meteorites, observations of molecular clouds and young stellar objects, and theoretical models give insight on the history of the early Solar System (e.g., Alexander et al., 2001). Grossman (1972) calculated the order (Table 1) that minerals would condense out of the solar nebula with decreasing temperature. Later studies (e.g., Ebel, 2006) confirmed this basic sequence. The first minerals that are predicted to condense are oxides and members of the melilite groups, which are common constituents of CAIs. CAIs are the oldest components of meteorites and are commonly found in chondritic meteorites. The oldest known CAIs have formation ages of 4.567–4.568 Ga (Amelin et al., 2010; Bouvier and Wadhwa, 2010) with uncertainties of hundreds of thousands of years. After CAIs, metallic iron and a number of FeO-free minerals (diopside, forsterite, anorthite, enstatite) would condense. Enstatite chondrites, which are predominately enstatite and metallic iron, are examples of such material. As the temperature decreases, more alkali feldspars and FeO-bearing olivines and pyroxenes would condense. Then, troilite will condense at ~ 700 K. At temperatures below ~ 490 K (Lewis, 1997), phyllosilicates will form from the reaction of water vapor with Fe-bearing olivines and pyroxenes. Magnetite will condense at ~ 405 K.
Table 1. Condensation sequence for different minerals from Grossman (1972) for a gas of solar composition at 10−3 atm
Temperature (K) Mineral
1758 (1513) Corundum (oxide), Al2O3
1647 (1393) Perovskite (oxide), CaTiO3
1625 (1450) Melilite group, Ca2Al2SiO7–Ca2MgSi2O7
1513 (1362) Spinel (oxide), MgAl2O4
1471 FeNi metal
1450 Diopside (pyroxene), CaMgSi2O6
1444 Forsterite (olivine), Mg2SiO4
1362 Anorthite (feldspar), CaAl2Si2O8
1349 Enstatite (pyroxene), Mg2Si2O6
< 1000 Alkali-bearing feldspar, (Na,K)AlSi3O8–CaAl2Si2O8
< 1000 Ferrous olivines, (Mg,Fe)2SiO4; ferrous pyroxenes, (Mg,Fe)2Si2O6
700 Troilite (sulfide), FeS
405 Magnetite (oxide), Fe3O4
The temperatures in parentheses are the temperatures where the condensate disappears.
Water ice will condense at temperatures of ~ 160–170 K in the solar nebula (e.g., Lunine, 2006). Organics can form in a variety of ways. One way to form organics is through Fischer–Tropsch reactions (e.g., Kress and Tielens, 2001) where CO and H2 are converted to hydrocarbons through heating with a transition metal catalyst. Simple organics can also form (e.g., Ehrenfreund and Charnley, 2000; Nuevo et al., 2011) in extremely cold environments through irradiation by ultraviolet irradiation and cosmic-ray bombardment of H2O, CO, CO2, CH3OH, and NH3 ices on the surfaces of silicate and carbonaceous dust grains.
Even though hydrated silicates should have formed in the solar nebula, the hydrated silicates in meteorites appear to have formed through aqueous alteration (e.g., Brearley, 2006; Tomeoka, 1990). For aqueous alteration to take place, the body must have liquid water and be exposed to a heat source (e.g., Grimm and McSween, 1989; McSween et al., 2002). Temperatures of at least ~ 300 K (e.g., Jones and Brearley, 2006) are thought to be needed to aqueously alter asteroids.
The temperature of the disk decreases with increasing distance from the Sun (e.g., Boss, 1998; Chapter 2.3) with the slope of the temperature decrease being very model-dependent. CAIs would condense closest to the Sun but have been shown to be able to dynamically diffuse outward (e.g., Boss, 2012; Ciesla, 2010; Cuzzi et al., 2003) so they can be incorporated into various chondritic meteorite classes (e.g., MacPherson and Boss, 2011; Scott and Krot, 2005) and comet-forming regions (e.g., Simon et al., 2008). X-winds (high speed bipolar collimated jets) around young stellar objects have also been proposed (e.g., Hu, 2010; Shu et al., 1996, 2001) to be the mechanisms for producing CAIs (and chondrules) and transporting them outward from the Sun; however, some researchers (e.g., Desch et al., 2010) do not believe the X-wind model is viable for producing and transporting CAIs.
Partially melted and fully differentiated parent bodies would form from the melting of chondritic planetesimals with the heat source usually assumed to be due to the decay of 26Al. Dating of differentiated and chondritic material find that the accretion of differentiated planetesimals occurred at an earlier time than the accretion of undifferentiated planetesimals (e.g., Baker et al., 2005; Qin et al., 2008). Aubrites are commonly assumed to be the result of melting of enstatite chondrite-like material (e.g., Keil, 1989, 2010). Experiments have shown (Jurewicz et al., 1991, 1993) partial melting of carbonaceous chondrites produce angrite- or eucrite-like depending on the oxygen fugacity. Ford et al. (2008) found that partial melting of ordinary chondrites does not appear to form primitive achondrite material (e.g., acapulcoites/lodranites) and argue that the precursor material for these meteorites was not as FeO-rich and oxidized as ordinary chondrites.
Ebel and Alexander (2011) predict that enstatite chondrites would have formed under similar conditions that formed Mercury at 0.4 AU, which also argues that enstatite chondrites formed in the inner solar system. From MESSENGER's x-ray measurements (Nittler et al., 2011; Weider et al., 2012) of its surface, Mercury's elemental abundance ratios are best matched by a partially melted enstatite chondrite mineralogy (Burbine et al., 2002b). Wasson and Wetherill (1979) previously proposed that enstatite-rich meteorites formed near or interior to 1 AU. Depending on the model, water ice would start condensing between ~2 and ~5 AU (e.g., Lunine, 2006). FeO-bearing silicates would condense at intermediate heliocentric distances between the regions where enstatite and water ice condense.
Meteorites, Comets, and Planets
J.E. Chambers, in Treatise on Geochemistry, 2007
1.17.2.1 Protoplanetary Disks
The solar system probably formed from the collapse of a “molecular cloud core,” a cold, dense portion of the interstellar medium containing gas and dust with a temperature of 10–20 K (see Chapter 1.04). Collapse may have occurred spontaneously or may have been triggered externally, for example by a supernova (Cameron, 1996). As the cloud core collapsed, most of its mass fell to the center to become a T Tauri star, while the remaining material formed a rotationally supported disk or protoplanetary nebula. Such disk-shaped regions are clearly seen around young stars silhouetted against brighter background material in the Orion nebula (O'Dell and Wen, 1994).
Many T Tauri stars give off unexpectedly large amounts of infrared radiation, which comes from fine grains of dust orbiting the star in an optically thick disk (Lada et al., 2006). This dust is typically larger than dust in the interstellar medium and shows signs of thermal processing (Kessler-Silacci et al., 2006). The fraction of stars with “infrared excesses” declines with age. Few stars older than 6 Myr have detectable amounts of hot dust orbiting within 0.1 AU (astronomical units) of the star (Haisch et al., 2001). Cooler dust, orbiting at larger distances, disappears on the same timescale as the hot dust (Andrews and Williams, 2005). The spectra of many T Tauri stars have ultraviolet and visible emission lines caused by hot gas from the disk accreting onto the star’s surface. The inferred accretion rates are 10−6–10−9 solar masses per year (Muzerolle et al., 2001).
The Sun’s protoplanetary disk must have contained at least 0.01 solar masses of material. This minimum-mass nebula is obtained by adding up the amount of rocky and icy material in the planets, and adding enough hydrogen and helium to give a nebula with the same composition as the Sun (Weidenschilling, 1977a). Theoretical models suggest that planet formation is actually an inefficient process, so the protoplanetary nebula was almost certainly more massive than this.
Most stars form in dense clusters containing hundreds or thousands of stars, similar to the Trapezium cluster in Orion. Large clusters generally contain at least one massive, OB star. These stars produce large amounts of ultraviolet radiation that can erode the outer parts of nearby protoplanetary disks by photoevaporation. Massive stars also release short-lived radioactive isotopes into the interstellar medium around them. There is evidence that many of these isotopes existed in the early solar system (see Section 1.17.3.2; Chapter 1.16), which suggests the Sun formed in a large stellar cluster.
The Precambrian Earth
P.G. Eriksson, ... O. Catuneanu, in Developments in Precambrian Geology, 2004
Formation of Our Solar System
Due at least in part to the lower mass of its precursor molecular cloud, the development of our solar system was fortunately very different from that outlined above for supernovae. Smaller stars of up to a few solar masses have longer pre-main sequence histories, and substantially longer lifetimes, than more massive stars. The interstellar cloud from which our solar system formed was derived from the ejecta of a range of stellar sources, including red giants and supergiants, AGB, nova, supernova and possibly also Wolf–Rayet stars (massive, high-temperature stars with extremely high mass-loss rates). A nearby supernova event injected newly synthesised, short-lived nuclides into the interstellar cloud and triggered its collapse to form a proto-Sun with radius about 5 times that of the present Sun over a period of < 105 years (see Cameron, 1995). Collapse will have occurred progressively from the inner to the outer part of the cloud, with conservation of angular momentum causing the collapsing cloud to spin faster. Collisions of dust and gas particles orbiting the proto-Sun in the same direction caused the loss of their energy, resulting in the flattening of the cloud, particularly near the centre where the densities are highest. Rotation of both the disk and the proto-Sun around a common centre of mass generated spiral density waves in the surrounding nebula. Within the evolving nebula, gravitational energy can be converted to heat during collapse and can initially be radiated away, so temperatures initially decrease with increasing distance from the cloud core. However, as the density of the cloud increased, heat could not be lost efficiently. At some time within 105 years from the onset of collapse of the cloud core, the proto-Sun commenced the violent early H-burning (or T-Tauri) phase of its evolution (Cameron, 1995). The infall rate of material from the rotating accretion disk increased episodically during this phase, although a significant proportion of the mass of the Sun was lost or recycled back into the disk via energetic bipolar outflows emitted along the axis of rotation. Magnetic field instabilities emanating from the Sun and vigorous flares, violent eruptions and strong stellar winds will have caused turbulence and mixing of ionised gas and dust within the accretion disk.
Temperatures within the accretion disk will have changed dynamically as the disk evolved, with parts shielded from the increasing temperatures of the inner nebula by the increasing density of the accumulating dust and gas closer to the cloud core and near the mid-plane of the disk (see Meibom et al., 2000). Where temperatures dropped below 1500 K, high-temperature refractory elements, such as Ca, Ti (titanium) and Al, condensed to form fragile “fluffy” sub-micron sized grains. These dust grains collided and accreted to form precursors of the Ca-and Al-rich inclusions (CAIs) of primitive meteorites. These inclusions are composed of a variety of high condensation temperature minerals, such as corundum (Al2O3), perovskite ([Ca, Na, Fe, Ce]TiO3), melilite ([Ca, Na]2 [Mg, Fe, Al, Si]3O7), hibonite ([Ca, Ce][Al, Ti, Mg]12O19) and spinel ([Mg, Fe, Ni, Cr] Al2O4). The more abundant Fe, Ni (nickel), and silicate-rich components condensed within parts of the nebula at lower temperatures, whereas volatile components such as water, ammonia, and methane ices, condensed only in the cold outer regions of the accretion disk. Volatile components in the inner solar system may have been carried to the outer regions of the solar system as ionised gas and dust by the solar wind (Shu et al., 1994). The Sun would have settled into the hydrogen-burning “main sequence” phase within 3–30 My of the onset of collapse of the cloud core (Strom et al., 1993; Cameron, 1995).
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Interstellar Matter
Donald G. York, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
V Properties of the Interstellar Clouds
Given the many possibilities for detection of radiation emitted or modified by interstellar gas, astronomers have pieced together a picture of the interstellar medium. In many cases, the details of the physical processes are vague. In most cases, detailed three-dimensional models cannot be constructed or are very model dependent. On the other hand, in some cases, sufficient knowledge exists to learn about other areas of astrophysics from direct observations. The example of the cosmic-ray distribution has already been given. Others are mentioned subsequently.
Interstellar clouds are complex aggregates of gas at certain velocities, typically moving at ±6 to 20 km/sec with respect to galactic rotation, itself ∼250 km/sec over most of the galaxy. Each cloud is a complex mixture of a volume of gas in a near pressure equilibrium and of isolated regions affected by transient pressure shocks or radiation pulses from star formation or from supernova explosions. Clouds are visible in optical, UV, or X-ray emission (or continuum scattering) when they happen to be close to hot stars and are otherwise detectable in absorption (molecules or atoms) or emission from low-lying excited levels (0.01 eV) or from thermal emission of the grains in the clouds.
V.A Temperatures
Molecular clouds are as cold as 10 K. Diffuse clouds are typically 100 K. HII regions have T ∼ 8000 K, depending on abundances of heavy elements that provide the cooling radiation. Low-column-density regions with 10,000 < T < 400,000 K are seen directly, presumably the result of heating at the cloud edges from shocks, X rays, and thermal conduction. Isolated regions with T > 106 K are seen near sites of supernova explosions.
V.B Densities
Densities, as determined from direct observation of excited states of atoms and molecules, are generally inversely proportional to temperature, implying the existence of a quasi-equilibrium state between the various phases of the medium. The effects of sources of disequilibrium in almost all cases last ≲107 years, or less than 1/10 of a galactic rotation time, itself 1/10 of the age of the sun. The product nT(cm−3K) is ∼3000 to within a factor of three where good measurements exist. Thus, the molecular (dark) clouds have n > 102 cm−3, while in diffuse clouds, n < 102 cm−3. Higher densities (up to 105 cm−3 occur in disequilibrium situations such as star-forming regions inside dense clouds and in HII regions.
V.C Abundances: Gas and Solid Phases
By measuring column densities of various elements with respect to hydrogen, making ionization corrections as necessary, abundances of elements in interstellar diffuse clouds can be determined. Normally, the abundances are compared with those determined in the sun.
Different degrees of depletion are found for different elements. Oxygen, nitrogen, carbon, magnesium, sulfur, argon, and zinc show less than a factor of two depletion. Silicon, aluminum, calcium, iron, nickel, manganese, and titanium show depletions of factors of 5–1000. Correlations of depletion with first ionization potential or with the condensation temperature (the temperature of a gas in thermal equilibrium at which gas-phase atoms condense into solid minerals), have been suggested, but none of these scenarios actually fits the data in detail.
The pattern of depletion suggests no connection with nucleosynthetic processes. Those elements that are depleted are presumed to be locked into solid material, called grains. Such particles are required by many other observations attributed to interstellar gas, as discussed earlier. In principle, the unknown makeup of the grains can be determined in detail by noting exactly what is missing in the gas phase. However, since there must be varying sizes and probably types of grains and since the most heavily depleted elements do not constitute enough mass to explain the total extinction per H atom, most of the grains by mass must be in carbon and/or oxygen. Establishing the exact mass of the grains amounts to measuring the depletions of C and O accurately. Although these measurements can now be made with the Hubble Space Telescope, problems of interpretation of deletions remain.
The grain structure (amorphous or crystalline) is not known. There are unidentified broad absorption features, called diffuse interstellar bands, that have been attributed to impurities in crystalline grains. However, these features may be caused by large molecules. It has been argued that even if grains are formed as crystalline structures, bombardment by cosmic rays would lead to amorphous structures over the life of the galaxy.
Theories of grain formation are uncertain. A general scenario is that they are produced in expanding atmospheres of cool supergiants, perhaps in very small “seed” form. They may then acquire a surface layer, called a mantle, probably in the form of water ice and solid CH4, NH3, etc. This growth must occur in cold, dense clouds. The detailed process and the distribution of atoms between minerals and molecules in solid phase are unknown.
V.D Evolution
Interstellar clouds can be large, up to 106 solar masses, and are often said to be the most massive entities in the galaxy. In this form, they may have a lifetime of more than 108 years. They are presumably dissipated as a result of pressure from stars formed within the clouds. Over the lifetime of the galaxy, interstellar clouds eventually turn into stars, the diffuse clouds being the residue from the star formation process. Growth of new molecular clouds from diffuse material is poorly understood. Various processes to compress the clouds have been suggested, including a spiral density wave and supernova blast waves. No one mechanism seems to dominate and several may be applicable. However, the existence of galaxies with up to 50% of their mass in gas and dust and of others with less than 1% of their mass in interstellar material leads to the inference that diffuse material and molecular clouds are eventually converted into stars.
Millimeter Astronomy
Jeffrey G. Mangum, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
III.A.2 Measurements of Physical Conditions
The study of molecular cloud stability and evolution leads naturally to studies of the physical and chemical evolution of the star formation process. Fundamental to this study of the star formation process is the characterization of the physical conditions in the gas and dust comprising these regions. For the gas, volume density n (cm−3), kinetic temperature TK (K), chemical composition X, turbulent motion Δ υ (km sec−1), and magnetic field strength B (Gauss) are fundamental physical quantities. For the dust, the dust temperature Td (K), dust volume density nd (cm−3), and dust opacity κ describe the physical conditions representative of the dust component of a molecular cloud. Note that all of these quantities are dependent upon time and position.
Most of the material in molecular clouds is in the form of H2, which owing to its lack of a permanent dipole moment has no easily observable rotational transitions. It can be observed through rovibrational and fluorescent transitions, but only within environments which are very specific, such as shocks and regions containing high levels of ultraviolet emission. Therefore, the principal component of molecular clouds is effectively unmeasurable. This fact forces astronomers to use trace constituents, other molecules and dust, to measure the physical conditions in molecular clouds.
III.A.2.a Molecular emission as a tracer of physical conditions
The primary constituent of molecular clouds, H2, is also the main collision partner with other molecular inhabitants of these regions. These collisions lead to the excitation of rotational transitions in a variety of molecules, many of which emit at observable millimeter wavelengths. The most abundant molecule after H2 is carbon monoxide (12C16O, usually simply written CO). It was the first molecule discovered at millimeter wavelengths by Wilson, Penzias, and Jefferts in 1970 using the National Radio Astronomy Observatory 36 ft (now 12 m) millimeter telescope located on Kitt Peak, Arizona. It has been used extensively as a probe of the volume density and kinetic temperature in molecular clouds through measurements of its lowest two rotational transitions at 115.271 and 230.538 GHz. CO has proven to be a very good tracer of the global physical conditions in molecular clouds, but for more compact regions with a larger number of particles along the line of sight [referred to as the column density (N) of a particular molecule], it loses its sensitivity to the bulk of the gas as the opacity in the measured transitions rises. Fortunately, there are other less abundant molecular tracers, including isotopomers (isotopic variants) of CO, such as 13CO, C18O, and 13C18O, which prove to be better probes of these high column density environments.
There are a wide variety of molecules that can be used as tracers of the volume density and kinetic temperature in molecular clouds. The choice of molecular probe depends upon what environment one wishes to study. For example, to measure the physical conditions in the dense cores of molecular clouds, it is best to choose a molecular tracer that is particularly sensitive to the prevalent conditions in this environment. A useful guide used to calculate the sensitivity of a transition to volume density is the critical density ncrit, which is the volume density required to collisionally excite a transition assuming optically thin conditions,
ncrit=AijCij=64π4νij33hc3giCij|μ⇀ji|2=64π4νij33hc3giCijSμ2,
where Aij is the spontaneous emission (Einstein A) coefficient for level i, Cij is the collisional deexcitation rate per molecule in level i, gi is the upper state degeneracy, |μ⇀ji| is the dipole moment matrix element for the transition, S is the line strength for the transition, μ is the dipole moment for the molecule, and the other terms have their usual meanings. Critical densities for common molecules such as CS, HCN, and H2CO are in the range 104−8 cm−3 for a kinetic temperature of 10 K.
Therefore, a simple detection of a transition from one of these molecules implies the existence of dense gas. A second consideration is to choose molecules that allow one to derive accurate measures of the volume density and kinetic temperature in a molecular cloud. Since the collisional excitation of molecular transitions is dependent upon the coupled effects of volume density and kinetic temperature, it is often necessary to use molecules that allow one to decouple these effects. The ability to decouple these physical effects depends upon the properties of the molecular structure. There are three basic types of molecules in this regard; linear, symmetric rotor, and asymmetric rotor. Figures 5–7 show the energy level structure for these three types of molecules. As can be seen from Fig. 5, linear molecules have one ladder of energy levels, the transitions between which are excited by the coupled effects of volume density and kinetic temperature. In general, linear molecules are used to derive the volume density in a molecular cloud by assuming a kinetic temperature or by using a calculation of the kinetic temperature based on measurements of the transitions from another molecule. The energy level structures for the symmetric and asymmetric rotor molecules shown in Figs. 6 and 7 indicate a more complex structure. Like linear molecules, the strengths of transitions within a given ladder (designated by the “K” rotational quantum numbers) are dependent upon the coupled effects of volume density and kinetic temperature. A comparison of the strengths of transitions from the same J levels but from different K ladders, though, is dependent only on the kinetic temperature, thus making it possible to derive a decoupled measurement of the kinetic temperature in a molecular cloud. In general, then, symmetric and asymmetric molecules have molecular level properties that, when the appropriate transitions are compared, allow decoupled measurements of the volume density and kinetic temperature in a molecular cloud. Linear molecules do not possess these decoupling properties, thus requiring an independent measurement of either the volume density or kinetic temperature.
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FIGURE 5. Rotational energy level diagram for HC3N, a typical linear molecule. The rotational quantum number “J” and associated energy above the ground state are indicated for each level.
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FIGURE 6. Rotational energy level diagram for CH3CN, a typical symmetric rotor molecule. The rotational quantum numbers J and K, along with the associated energy above the ground state, are indicated for each level.
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FIGURE 7. Rotational energy level diagram for H2CO, a typical asymmetric rotor molecule. The rotational quantum numbers J, K−1, and K+1, along with the associated energy above the ground state, are indicated for each level.
By comparing the measured intensities of a variety of transitions from a given molecule with molecular line intensity predictions from a molecular cloud model, estimates of the volume density and kinetic temperature within a molecular cloud can be made. In general, these estimates reveal that the volume densities in the regions where stars form within molecular clouds exceed 104 cm−3, while the kinetic temperatures range from 10 to 300 K. These models also indicate that many molecular cores possess density gradients, suggestive of a structure that could evolve into a collapsing protostar.
III.A.2.b Kinematics
The shape of a spectral line is determined by the radial velocity structure along the line of sight through a molecular cloud. The measured widths of spectral lines are generally larger than the thermal width, indicating that the velocity fields within molecular clouds are dominated by Doppler broadening owing to turbulence:
Δυ=υtherm+υturb=22ln2kTKm+υturb
where Δ υ is the full width at half maximum of the spectral line, TK is the kinetic temperature of the gas, and m is the mass of the particles which make up the molecular cloud (principally molecular hydrogen). Unfortunately, a detailed derivation of the spectral line shape has proven elusive, owing to the effects of kinetic temperature and volume density gradients, spatial structure, and radiative transfer effects within the molecular cloud. Measurements of the line center velocity as a function of position over a molecular cloud do indicate that, in general and on large scales, they are neither collapsing nor rotating. This is not the case on small scales. Evidence for rotation and collapse of molecular cloud cores on 0.1 pc scales have yielded interesting clues to the details of the star formation process. Measurements of cloud core rotation indicate that in magnitude it is only 2% of the gravitational potential energy before collapse, making it relatively unimportant to the overall dynamics of a molecular cloud core.
The physical nature of the turbulent component of a spectral line in a molecular cloud is currently a source of considerable debate. Physical processes that have been suggested as sources of the turbulence in molecular clouds are expanding HII regions, supernova remnants, cloud–cloud collisions, galactic differential rotation, and stellar winds. Unfortunately, for all of these processes there are theoretical problems with coupling the energy produced into turbulence.
III.A.2.c Magnetic fields
An understanding of the magnetic field properties of molecular clouds is an important aspect of the overall physical understanding of molecular clouds given their apparent role in providing dynamical support in these environments. There are three methods that have been used to measure the magnetic field strength and direction within molecular clouds: atomic and molecular Zeeman effect splitting, which tell us about the line-of-sight magnetic field component (BZ); polarization of the emission from dust grains, which gives us information about the component of the magnetic fields perpendicular to the line-of-sight (B⊥); and measurements of spectral line emission polarization, which also tells us about B⊥. Measurements of the Zeeman splitting in atoms and molecules have concentrated on studies of HI, OH, and CN, yielding typical values for BZ of 10–20 μG within regions with volume densities of approximately 103 cm−3. At higher volume densities, magnetic fields as large as 700 μG have been measured. Millimeter dust continuum emission polarization levels of at most a few percent have shown that the magnetic field within the high volume density (n ≥ 106 cm−3) cores of molecular clouds is perpendicular to the major axis of the high density structures (such as disks) and parallel to the outflows associated with these objects. Although the possibility of there being measurable levels of polarization of thermal millimeter spectral line emission has been known for years, it has only recently been detected. The percentage of measured polarized emission is equivalent to that detected through millimeter continuum polarimetry.
Turbulence and magnetic fields cannot simply be applied as independent solutions to the problem of cloud support since turbulence should tangle magnetic fields, thus reducing their effectiveness as a source of support. The theory of magnetohydrodynamic turbulence within molecular clouds has shown that magnetic fields should slow the decay of turbulent motions if these motions are less than the propagation speed along the magnetic field lines, referred to as the Alfvén velocity. However, the stability of magnetic support in the presence of turbulence has been called into question, and the interplay between cloud stability and dynamics drives our understanding of the importance of magnetic fields as a means of molecular cloud support. Future measurements of the magnetic field and direction in molecular clouds should clarify their influence on the overall dynamics and evolution of these regions.
Astrochemistry
Steven N. Shore, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
V.B.2 Ionization
Ionization in the densest parts of molecular clouds depends on the penetration of cosmic rays and UV radiation, as well as the presence of shocks generated by such processes as cloud–cloud collisions and internal star formation. In order to probe the electron density in the clouds, it is important to be able to account for the presence of complex polyatomic molecules, whose formation requires ion gas-phase reactions.
The rate for cosmic-ray (CR) ionization, ζCR, is about (3 ± 1) × 10−17 s−1. This is an integral over the collisional ionization cross section for low (MeV)-energy CR protons, but it is approximately a constant for most of the species of interest. An obstacle in our understanding of the detailed structure of molecular clouds is our ignorance of the precise specification of this rate. The low-energy end of the cosmic-ray spectrum is difficult to determine empirically from terrestrial observation, because these particles propagate diffusively through the interplanetary medium, scattering off of turbulence in the solar wind; their spectrum cannot be observed directly, even with in situ measurements from the Voyager and Ulysses spacecraft, and must be inferred from models for their motion through the heliosphere. The more easily observed cosmic ray protons and electrons, in the GeV and higher range, have little or no effect on the ionization of the interstellar medium because of the small interaction cross sections for atoms at such high energies.
In molecular clouds, atomic species with ionization energies greater than 13.6 eV must be predominantly neutral because of the shielding effects of neutral hydrogen. It is mainly the heavier elements, such as C, N, and O, which are observed in the peripheral portions of the clouds to be in the partially ionized state. For circumstellar envelopes, cosmic rays lose out to photo processes and the chemistry is mediated by the input of stellar photospheric radiation (in the hotter stars and in novae and supernovae) and from the diffuse interstellar radiation field.
The basic equations for two body interactions can be written in the form
(24)dNidt=∑j,k≠iKijkNjNk−∑jKij′NiNj,
where Kijk is the formation rate for the ith molecular species, while K′ij is the destruction rate for the molecule. The inclusion of UV photo processes is accomplished by the photodissociation rate:
(25)Rpd=∫ν0∞κνFνe−τνdνhν,
where Fν is the incident photon flux, τν is the opacity of the ambient medium (presumed to be from dust), κν is the continuous absorption coefficient for the dissociative continuum, and the dissociation energy is hν0.
An aspect in which circumstellar environments differ from interstellar is the net mass advection through the medium. Abundances become time dependent—and hence space dependent—in the envelope, due both to the implicit time dependence of the reactions and to the transport of matter through different radii via stellar wind flow. The atomic abundances are fixed at stellar photosphere, rather than having to be assumed for some mixture of physical parameters of temperature and pressure as they must for molecular clouds. It is then essentially an initial-value problem to compute the abundances which will be a function of radius in the envelope. For a steady-state wind, the abundances become strictly a function of radius. Also, unlike a molecular cloud, the density profile of the envelope is specified from the assumption of steady mass loss at the terminal velocity for the wind, so thatρ(r)=M./(4πr2υ∞), where M. is the mass loss rate and υ∞ is the terminal velocity of the wind.
An interesting aspect of stellar envelopes is that they may have two different sources of UV radiation, internal and external. Work on the envelope of two extreme, low-temperature, evolved supergiants, IRC + 10216 and α Ori, showed that the outer limit of the molecular envelope is determined by the DIRF, which destroys the outermost molecular species by photodissociation, while the inner boundary is set by both the temperature and UV emission from the stellar chromospheres. In this respect, since the dynamics can be probed in exquisite detail for several of the nearer supergiants through molecular observations, and since the input abundances are known and atomic in nature, it is possible to use these stars as very well-conditioned laboratories for the study of the same processes which must be involved in at least some aspects of molecular cloud chemistry. For the densest envelopes, which are completely optically thick and hence very similar to molecular clouds, cosmic rays are significant in governing the ion fractions but can be neglected in thin envelopes (low mass loss rates).
The Hadean and Archaean Atmosphere and the Oldest Records of Life as Micro- or Chemofossils
Akio Makishima, in Origins of the Earth, Moon, and Life, 2017
8.2 The Perspective of Atmospheric Evolution From the Hadean to the Archean Earth
As the sun formed from its molecular cloud, it was accompanied by disk material that consisted of gas and small dust particles. Over several tens of millions of years, these dust particles formed the planets. This process occurred in several stages in the terrestrial planet zone, including moon-forming impacts on the proto-Earth (Canup and Asphaug, 2001).
A Hadean atmosphere containing N2 and CO2 and a Hadean ocean containing H2O seems to have formed as a natural consequence of planetary accretion in the terrestrial planet region. The atmosphere which had weak reducing potential with relatively high partial pressure of CO2 should have formed (Holland, 1984; Walker, 1985; Kasting, 1993; Ferus et al., 2015; Furukawa et al., 2015). In this condition the important biological precursor compounds for life were synthesized. It should be noted that high partial pressure of CO2 in the early Hadean atmosphere can be presumed. Such gas is like intestinal gas, so it is reasonable to suppose that life may have fermented.
Atmospheric O2 levels rose naturally and gradually, but not immediately, occurring by photosynthesis and organic carbon burial. At the same time, the concentrations of CO2 and other greenhouse gases did not compensate for the brightening sun. The Earth’s relatively stable climate was a result of the negative feedback between atmospheric CO2, surface temperature, and the weathering rate of silicate rocks (Kasting, 1993).
This atmospheric evolution implies that Earth is not a unique planet. If planets exist around other stars, some of them could reside in orbits where the illumination is similar to that received by the Earth. Planetary climates are buffered by the carbonate–silicate cycle. Therefore, the habitable zone around late F to mid K stars (see the top scale of Fig. Box 1.4; the Hertzsprung–Russel diagram) may be wider, and other habitable planets may exist. If the origin of life was not a fortuitous event, many of these planets could be inhabited and on some, intelligent life may even have evolved. Both of these speculations can be tested: the first by spectroscopic investigations from large, space-based telescopes; and the second by monitoring microwave and radio emissions from space (Kasting, 1993).
Morbidelli et al. (2000) suggested that the most plausible sources of the water accreted by the Earth were in the outer asteroid belt, in the giant planet regions, and in the Kuiper Belt. It is plausible that the Earth accreted water from the early phases when the solar nebula was still present to the late stages of gas-free scattered planetesimals. Asteroids and comets from the Jupiter-Saturn region were the first water deliverers, when the Earth was less than half its present mass. The bulk water presently on the Earth was carried by a few planetary embryos, originally formed in the outer asteroid belt and accreted by the Earth at the final stage of its formation. (see Fig. 2.10: The blue planetesimals are water rich, which could be a source of water for Earth).
Finally, a late veneer (this could be the same as or different from the late veneer for the highly siderophile elements discussed in Chapter 4), accounting for at most 10% of the present water mass, occurred due to comets from the Uranus-Neptune region and from the Kuiper Belt. The net result of accretion from these several reservoirs is that the D:H ratio of water on Earth is essentially the typical water condensed in the outer asteroid belt. This is in agreement with the observation that the D:H ratio in the oceans is very close to the mean value of the D:H ratio of water inclusions in carbonaceous chondrites.
Stars, Massive
Steven N. Shore, in Encyclopedia of Physical Science and Technology (Third Edition), 2003
IV.B Induced Star Formation
Should the star formation begin within a molecular cloud, the winds and H II regions can either destroy the cloud by heating it up through radiative and mechanical processes or they can break free of the cloud. This latter type of star formation event is called a champagne flow, analogous to that commonly observed at weddings. The sudden release of pressure by the breakout of the shock from its environment causes a rapid outflow of material from the hot H II region surrounding the OB stars. This flow further drives a shock into the molecular cloud via momentum conservation and compresses the already dense material of the cloud core. Such an event may initiate collapse from gravitational instability, at least locally to the OB stars. Thus, the massive stars are not only capable of destoying the cloud environments in which they have formed, they can also serve as agents for propagating star formation through the cloud.
The formation of the massive stars is well traced by the radio and infrared flux emitted by the H II regions that surround them. Even the densest parts of molecular clouds are not optically thick longward of about 60 μm, so that far infrared (FIR) and radio photons freely escape the cloud. The argument that permits these to determine the rate of massive star formation is as follows. Every massive star that forms an H II region will be surrounded by a thermal, radio-emitting plasma with a temperature of about 104 K. The rate of radio emission depends on the rate of ionization, which in turn depends on the luminosity of the central stars. This in turn depends on the mass of the stars. Each radio photon is associated with the ionized medium, while each recombination eventually leads to a Lyα photon through radiative cascades. These photons, which are trapped in the optically thick H II regions, eventually collide with dust grains in the cloud and the H II region and are absorbed. The dust reradiates this energy at equilibrium in the FIR. Thus, there is an expected correlation between LFIR, the far infrared luminosity, Lradio, and ψ, the star formation rate per unit mass.
Such a correlation is observed in regions of active star formation in the Galaxy and nearby galaxies, although it is still not completely certain what the implied rates of star formation mean. For some galaxies, called starburst galaxies, the rate of massive star formation can reach more than 103 M⊙ yr−1. This enormous rate cannot continue for very long before molecular cloud material is consumed, but the triggering mechanism is not currently understood. Such systems may be the result of collisions between galaxies and are often seen in galactic systems undergoing merger.
UV astronomy and the investigation of the origin of life
Ana I. Gómez de Castro, Ada Canet, in Ultraviolet Astronomy and the Quest for the Origin of Life, 2021
2.3 UV spectral energy distribution and the photoelectric yield
In diffuse environments such as the envelopes of the molecular clouds, the atmosphere of young planetary disks or the interplanetary medium, UV photons are absorbed by materials in the dust grains (graphites, silicates and ferrites, mainly) yielding to the emission of photoelectrons that heat and ionize the environment and accelerate some chemical processes. Only UV radiation, with energy above ∼6 eV, is able to produce a significant photoelectric yield.
The heating rate depends on the density and energy of the ejected photoelectrons which, in turn, depends on the dust grains size (Watson, 1972), composition and charge state as well as on the spectral energy distribution (SED) of the radiation feeding the photoelectric flow (i.e. Pedersen and Gómez de Castro, 2011). The photoelectric yields are displayed in Fig. 2.4 (after Weingartner and Draine, 2001) for silicates and carbonaceous particles. Photons with energies above some 6 eV produce a significant yield from carbonaceous particles while the threshold goes up to 8 eV for silicates.
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Figure 2.4. Photoelectric yield of silicate and carbonaceous particles in space. The various curves correspond to different particle size. In general, the smallest the size the larger the yield because of the finite electron escape length.
After Weintgarner & Draine (2001).
The effectiveness of the process depends on the kinetic energy acquired by the photoelectrons to overcome the Coloumb field of the charged dust grain. Hence, there is a strong dependency between the final charge of the grains and the SED of the UV spectrum irradiating them. This effect shows in the temporal evolution of the charge of the grains (or charging profile) when submitted to a specific UV spectrum, as shown in Fig. 2.5. Two very different spectra are selected to illustrate this process: the very hot, but soft SED of an O-type star and the emission lines spectrum of a pre-main sequence (PMS) solar-like star. Above 6 eV, PMS stars only radiate significantly in very narrow bands associated to the main spectral lines at: ∼8 eV (the He II at 164 nm and CIV at 1550 nm lines), 8.9 eV (Si IV), 9.3 eV (CII), 9.5 eV (OI) and, of course, the Lyα hydrogen line at 10.2 eV. This is the radiation affecting the dust grains in young planetary disks. Dust in the HII regions and the bright reflection nebula are however, charged by the smooth UV spectrum of hot stars. Some examples are the B33 (Horsehead) nebula irradiated by a cluster of early B-type late O-type stars (Caballero, 2007) and the Rosette nebula by the cluster NGC 2244 of B8-9 type stars and 7 O4–O9 stars (Martins et al., 2012). As shown in the bottom panel of Fig. 2.6, dust charging profiles are step-wise in young planetary disks instead of the smooth rising curves from HII regions which are, otherwise consistent with the experimental observation (Sickafoose et al., 2000).
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Figure 2.5. Top: STIS spectra of the Sun (in red, from Meftah et al., 2020), the PMS binary AK Sco (blue, from Gómez de Castro et al., 2020b) and the O-type star HD46223 (yellow, data from the Hubble archive). AK Sco and HD46223 fluxes are scaled so the integrated flux is the same as the solar flux received from the Earth at 1 AU in the STIS wavelength range. Bottom: Charging profile of a silicate dust grain (a = 30 nm) submitted to the radiation fields of an O-type and a PMS star (Pedersen & Gómez de Castro, 2011).
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Figure 2.6. Annular bright field (ABF) images of a Highly Ordered Pyrolytic Graphite (HOPG) sample after irradiation with a Deuterium UV lamp. The high-resolution scanning transmission electron microscopy (STEM) images resolve the separation between the individual graphene layers in the HOPG. The image evidences the dislocation of the superficial graphitic layers after irradiation and the formation of bumps. Electron Energy Loss spectroscopy (EELs) of the bumps show that the sp2 to sp3 hybridation transition is a surface effect that is maximum at the irradiated surface. High sensitivity X-ray diffraction experiments confirm the formation of diamond within the bumps. On the right, the separation between the graphene layers is plotted. There is a transitional area where separation between the layers increases from 3.53 Å to 3.87 Å.
After Gómez de Castro et al., (2020a).
As interstellar and interplanetary dust grains are not in isolation, the final grain charge will also depend on the temperature of the surrounding gas that acts as a regulator since dust grain’s charge is modified in the gas-dust collisions (Spitzer, 1978).
Formation of the Proto-Earth in the Solar Nebula
Akio Makishima, in Origins of the Earth, Moon, and Life, 2017
2.1 Evolution of Molecular Clouds to the Solar Nebula
The evolution of an early star and its surrounding molecular cloud (nebula) are controlled by four components: (1) the centrifugal force; (2) mass ejection (mass and angular momentum loss as ejection jet); (3) mass accretion (mass gain from the disk surrounding the star); (4) a magnetic field (the mass ejection and the accretion follow the magnetic field). These four components are schematically depicted in Fig. 2.1 (Montmerle et al., 2006).
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Figure 2.1. The four components for evolution of the early star and surrounding disk (each component is not in scale).
The star has two strong magnetic fields (B). One magnetic field extends vertically along the rotation axis of the star. Another magnetic field stretches horizontally. In the accretion-ejection model, these two magnetic fields meet along the X-ring (or at the X-point if plotted two dimensionally). The accreting materials from the horizontal direction are ejected at the X-ring.
Our solar system is considered to have been formed from cosmic gas clouds as shown in Fig. 2.2. The Fig. 2.2 was obtained by the Hubble Space Telescope (HST) (see Box 2.1). These gas clouds are made of hydrogen, helium, and heavy elements that are the remnants of a supernova (see Chapter 1). The shockwave from a supernova produced inhomogeneity of the gas cloud near the location where our Sun would appear. The gas cloud about 3 light years began to gather and form a dense gas disk (103 ∼ 104 AU; AU is the present distance between the Sun and Earth). At the center of the gas disk was a huge gas ball (protostar) mainly composed of hydrogen.
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Figure 2.2. The close-up of cosmic clouds and stellar winds at the Orion Nebula.
Image credit: NASA, European Space Agency, and the Hubble Heritage Team www.nasa.gov/sites/default/files/images/724329main_hubble_feature_full_full.jpg.
Box 2.1
The Hubble Space Telescope (HST)
The HST was launched into Earth orbit in 1990 (see Fig. Box 2.1A and B). Its 2.4-m mirror and main instruments can observe the near ultraviolet, visible, and near infrared spectra. The telescope was named after astronomer Edwin Hubble. The HST can take high-resolution images without atmospheric distortion and the background light of the atmosphere. The HST was built by the National Aeronautics and Space Administration (NASA, the U.S. space agency), with contributions from the European Space Agency.
The various stellar phases with timescales and sizes are shown in Fig. 2.3 after Feigelson and Montmerle (1999). These models were established due to the advancement of observation techniques using infrared and X-ray telescopes. The initial stage, which is previously explained, is called the infalling protostar stage (see the top figure of Fig. 2.3). The accretion and ejection of the molecular gases occur. Although the gas ball is very large, it cannot be seen from outside because of the surrounding clouds. This stage lasts fewer than 104 years.
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Figure 2.3. Various stellar phases with timescales and sizes (AU = Sun-Earth distance = 1.5 × 108 km; each stage cartoon is not in scale).
The next stage is the evolved protostar stage. The accretion and ejection of the gases are still vigorously occurring. However, the size of the gas ball becomes 10 times smaller. This stage lasts for 0.1 Myr. The gas disk also glows with accretion.
Then the classic T Tauri stage begins. This stage is also the ejection stage. The gas clouds are vigorously ejected, and the ejection jet is clearly observed. This stage lasts for 1–10 Myr, and the thick disk shrinks to about 100 AU. In this stage, the planetary system would be formed.
The next stage is the weak-lined T Tauri stage. “Weak-lined” means that a “weak-Hα” line of ionized hydrogen is observed in the spectrum of the (proto)star.
Finally, the clouds dissipate and the Sun can be seen. The Sun becomes a main sequence star in the H-R diagram (see Fig. Box 1.4).
The evolution model described here is for the evolution of the Sun (star) based on actual observation of stars and theories. However, when the masses of the presolar nebula or the planetary system are too small compared to the Sun, the evolution of the presolar nebula requires different views and theories.
Planets, Asteriods, Comets and The Solar System
T.H. Burbine, in Treatise on Geochemistry (Second Edition), 2014
2.14.7.3 Formation of Material in the Solar Nebula
Mineralogical and isotopic studies of meteorites, observations of molecular clouds and young stellar objects, and theoretical models give insight on the history of the early Solar System (e.g., Alexander et al., 2001). Grossman (1972) calculated the order (Table 1) that minerals would condense out of the solar nebula with decreasing temperature. Later studies (e.g., Ebel, 2006) confirmed this basic sequence. The first minerals that are predicted to condense are oxides and members of the melilite groups, which are common constituents of CAIs. CAIs are the oldest components of meteorites and are commonly found in chondritic meteorites. The oldest known CAIs have formation ages of 4.567–4.568 Ga (Amelin et al., 2010; Bouvier and Wadhwa, 2010) with uncertainties of hundreds of thousands of years. After CAIs, metallic iron and a number of FeO-free minerals (diopside, forsterite, anorthite, enstatite) would condense. Enstatite chondrites, which are predominately enstatite and metallic iron, are examples of such material. As the temperature decreases, more alkali feldspars and FeO-bearing olivines and pyroxenes would condense. Then, troilite will condense at ~ 700 K. At temperatures below ~ 490 K (Lewis, 1997), phyllosilicates will form from the reaction of water vapor with Fe-bearing olivines and pyroxenes. Magnetite will condense at ~ 405 K.
Table 1. Condensation sequence for different minerals from Grossman (1972) for a gas of solar composition at 10−3 atm
Temperature (K) Mineral
1758 (1513) Corundum (oxide), Al2O3
1647 (1393) Perovskite (oxide), CaTiO3
1625 (1450) Melilite group, Ca2Al2SiO7–Ca2MgSi2O7
1513 (1362) Spinel (oxide), MgAl2O4
1471 FeNi metal
1450 Diopside (pyroxene), CaMgSi2O6
1444 Forsterite (olivine), Mg2SiO4
1362 Anorthite (feldspar), CaAl2Si2O8
1349 Enstatite (pyroxene), Mg2Si2O6
< 1000 Alkali-bearing feldspar, (Na,K)AlSi3O8–CaAl2Si2O8
< 1000 Ferrous olivines, (Mg,Fe)2SiO4; ferrous pyroxenes, (Mg,Fe)2Si2O6
700 Troilite (sulfide), FeS
405 Magnetite (oxide), Fe3O4
The temperatures in parentheses are the temperatures where the condensate disappears.
Water ice will condense at temperatures of ~ 160–170 K in the solar nebula (e.g., Lunine, 2006). Organics can form in a variety of ways. One way to form organics is through Fischer–Tropsch reactions (e.g., Kress and Tielens, 2001) where CO and H2 are converted to hydrocarbons through heating with a transition metal catalyst. Simple organics can also form (e.g., Ehrenfreund and Charnley, 2000; Nuevo et al., 2011) in extremely cold environments through irradiation by ultraviolet irradiation and cosmic-ray bombardment of H2O, CO, CO2, CH3OH, and NH3 ices on the surfaces of silicate and carbonaceous dust grains.
Even though hydrated silicates should have formed in the solar nebula, the hydrated silicates in meteorites appear to have formed through aqueous alteration (e.g., Brearley, 2006; Tomeoka, 1990). For aqueous alteration to take place, the body must have liquid water and be exposed to a heat source (e.g., Grimm and McSween, 1989; McSween et al., 2002). Temperatures of at least ~ 300 K (e.g., Jones and Brearley, 2006) are thought to be needed to aqueously alter asteroids.
The temperature of the disk decreases with increasing distance from the Sun (e.g., Boss, 1998; Chapter 2.3) with the slope of the temperature decrease being very model-dependent. CAIs would condense closest to the Sun but have been shown to be able to dynamically diffuse outward (e.g., Boss, 2012; Ciesla, 2010; Cuzzi et al., 2003) so they can be incorporated into various chondritic meteorite classes (e.g., MacPherson and Boss, 2011; Scott and Krot, 2005) and comet-forming regions (e.g., Simon et al., 2008). X-winds (high speed bipolar collimated jets) around young stellar objects have also been proposed (e.g., Hu, 2010; Shu et al., 1996, 2001) to be the mechanisms for producing CAIs (and chondrules) and transporting them outward from the Sun; however, some researchers (e.g., Desch et al., 2010) do not believe the X-wind model is viable for producing and transporting CAIs.
Partially melted and fully differentiated parent bodies would form from the melting of chondritic planetesimals with the heat source usually assumed to be due to the decay of 26Al. Dating of differentiated and chondritic material find that the accretion of differentiated planetesimals occurred at an earlier time than the accretion of undifferentiated planetesimals (e.g., Baker et al., 2005; Qin et al., 2008). Aubrites are commonly assumed to be the result of melting of enstatite chondrite-like material (e.g., Keil, 1989, 2010). Experiments have shown (Jurewicz et al., 1991, 1993) partial melting of carbonaceous chondrites produce angrite- or eucrite-like depending on the oxygen fugacity. Ford et al. (2008) found that partial melting of ordinary chondrites does not appear to form primitive achondrite material (e.g., acapulcoites/lodranites) and argue that the precursor material for these meteorites was not as FeO-rich and oxidized as ordinary chondrites.
Ebel and Alexander (2011) predict that enstatite chondrites would have formed under similar conditions that formed Mercury at 0.4 AU, which also argues that enstatite chondrites formed in the inner solar system. From MESSENGER's x-ray measurements (Nittler et al., 2011; Weider et al., 2012) of its surface, Mercury's elemental abundance ratios are best matched by a partially melted enstatite chondrite mineralogy (Burbine et al., 2002b). Wasson and Wetherill (1979) previously proposed that enstatite-rich meteorites formed near or interior to 1 AU. Depending on the model, water ice would start condensing between ~2 and ~5 AU (e.g., Lunine, 2006). FeO-bearing silicates would condense at intermediate heliocentric distances between the regions where enstatite and water ice condense.
Meteorites, Comets, and Planets
J.E. Chambers, in Treatise on Geochemistry, 2007
1.17.2.1 Protoplanetary Disks
The solar system probably formed from the collapse of a “molecular cloud core,” a cold, dense portion of the interstellar medium containing gas and dust with a temperature of 10–20 K (see Chapter 1.04). Collapse may have occurred spontaneously or may have been triggered externally, for example by a supernova (Cameron, 1996). As the cloud core collapsed, most of its mass fell to the center to become a T Tauri star, while the remaining material formed a rotationally supported disk or protoplanetary nebula. Such disk-shaped regions are clearly seen around young stars silhouetted against brighter background material in the Orion nebula (O'Dell and Wen, 1994).
Many T Tauri stars give off unexpectedly large amounts of infrared radiation, which comes from fine grains of dust orbiting the star in an optically thick disk (Lada et al., 2006). This dust is typically larger than dust in the interstellar medium and shows signs of thermal processing (Kessler-Silacci et al., 2006). The fraction of stars with “infrared excesses” declines with age. Few stars older than 6 Myr have detectable amounts of hot dust orbiting within 0.1 AU (astronomical units) of the star (Haisch et al., 2001). Cooler dust, orbiting at larger distances, disappears on the same timescale as the hot dust (Andrews and Williams, 2005). The spectra of many T Tauri stars have ultraviolet and visible emission lines caused by hot gas from the disk accreting onto the star’s surface. The inferred accretion rates are 10−6–10−9 solar masses per year (Muzerolle et al., 2001).
The Sun’s protoplanetary disk must have contained at least 0.01 solar masses of material. This minimum-mass nebula is obtained by adding up the amount of rocky and icy material in the planets, and adding enough hydrogen and helium to give a nebula with the same composition as the Sun (Weidenschilling, 1977a). Theoretical models suggest that planet formation is actually an inefficient process, so the protoplanetary nebula was almost certainly more massive than this.
Most stars form in dense clusters containing hundreds or thousands of stars, similar to the Trapezium cluster in Orion. Large clusters generally contain at least one massive, OB star. These stars produce large amounts of ultraviolet radiation that can erode the outer parts of nearby protoplanetary disks by photoevaporation. Massive stars also release short-lived radioactive isotopes into the interstellar medium around them. There is evidence that many of these isotopes existed in the early solar system (see Section 1.17.3.2; Chapter 1.16), which suggests the Sun formed in a large stellar cluster.
The Precambrian Earth
P.G. Eriksson, ... O. Catuneanu, in Developments in Precambrian Geology, 2004
Formation of Our Solar System
Due at least in part to the lower mass of its precursor molecular cloud, the development of our solar system was fortunately very different from that outlined above for supernovae. Smaller stars of up to a few solar masses have longer pre-main sequence histories, and substantially longer lifetimes, than more massive stars. The interstellar cloud from which our solar system formed was derived from the ejecta of a range of stellar sources, including red giants and supergiants, AGB, nova, supernova and possibly also Wolf–Rayet stars (massive, high-temperature stars with extremely high mass-loss rates). A nearby supernova event injected newly synthesised, short-lived nuclides into the interstellar cloud and triggered its collapse to form a proto-Sun with radius about 5 times that of the present Sun over a period of < 105 years (see Cameron, 1995). Collapse will have occurred progressively from the inner to the outer part of the cloud, with conservation of angular momentum causing the collapsing cloud to spin faster. Collisions of dust and gas particles orbiting the proto-Sun in the same direction caused the loss of their energy, resulting in the flattening of the cloud, particularly near the centre where the densities are highest. Rotation of both the disk and the proto-Sun around a common centre of mass generated spiral density waves in the surrounding nebula. Within the evolving nebula, gravitational energy can be converted to heat during collapse and can initially be radiated away, so temperatures initially decrease with increasing distance from the cloud core. However, as the density of the cloud increased, heat could not be lost efficiently. At some time within 105 years from the onset of collapse of the cloud core, the proto-Sun commenced the violent early H-burning (or T-Tauri) phase of its evolution (Cameron, 1995). The infall rate of material from the rotating accretion disk increased episodically during this phase, although a significant proportion of the mass of the Sun was lost or recycled back into the disk via energetic bipolar outflows emitted along the axis of rotation. Magnetic field instabilities emanating from the Sun and vigorous flares, violent eruptions and strong stellar winds will have caused turbulence and mixing of ionised gas and dust within the accretion disk.
Temperatures within the accretion disk will have changed dynamically as the disk evolved, with parts shielded from the increasing temperatures of the inner nebula by the increasing density of the accumulating dust and gas closer to the cloud core and near the mid-plane of the disk (see Meibom et al., 2000). Where temperatures dropped below 1500 K, high-temperature refractory elements, such as Ca, Ti (titanium) and Al, condensed to form fragile “fluffy” sub-micron sized grains. These dust grains collided and accreted to form precursors of the Ca-and Al-rich inclusions (CAIs) of primitive meteorites. These inclusions are composed of a variety of high condensation temperature minerals, such as corundum (Al2O3), perovskite ([Ca, Na, Fe, Ce]TiO3), melilite ([Ca, Na]2 [Mg, Fe, Al, Si]3O7), hibonite ([Ca, Ce][Al, Ti, Mg]12O19) and spinel ([Mg, Fe, Ni, Cr] Al2O4). The more abundant Fe, Ni (nickel), and silicate-rich components condensed within parts of the nebula at lower temperatures, whereas volatile components such as water, ammonia, and methane ices, condensed only in the cold outer regions of the accretion disk. Volatile components in the inner solar system may have been carried to the outer regions of the solar system as ionised gas and dust by the solar wind (Shu et al., 1994). The Sun would have settled into the hydrogen-burning “main sequence” phase within 3–30 My of the onset of collapse of the cloud core (Strom et al., 1993; Cameron, 1995).
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