Following the Interstellar History of Carbon: From the Interiors of Stars to the Surfaces of Planets

2.1. Stellar nucleosynthesis of carbon and transport into the ISM

The origin of carbon lies in aging stars. Carbon-12 is predicted to be formed in the triple-α process, in which extremely high temperatures and pressures in stellar interiors cause the fusion of three α particles: 3⁴He → ¹²C (Timmes et al., 1995). This process occurs in very massive stars as they become Type II supernovae; ¹²C is formed in two separate shells in the layered structure of the stellar interior as the star explodes and is subsequently ejected into the ISM (Clayton, 2003). It is thought that about half the carbon-12 comes from Type II supernovae. The remaining half originates in the final phases of stellar evolution for low- to intermediate-mass stars (M _* ∼ 1–8 M _⊙), as does the less abundant isotope, ¹³C. The concluding stages of the life cycle of these stars are illustrated in a tool known as the Hertzsprung-Russell (HR) diagram, as presented in Fig. 3.

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FIG. 3.

Qualitative Hertzsprung-Russell (HR) diagram, showing the evolution of an intermediate-mass star through the red giant and asymptotic giant branch (AGB) phases and into the planetary nebula (PN) stage. Important steps in nucleosynthesis in the stellar interior and names of representative objects are shown on the diagram.

As shown in Fig. 3, low/intermediate-mass stars evolve off the so-called “main sequence” when they have converted all the hydrogen in their cores into helium. About 90% of all stars are currently on the main sequence, a stage characterized by core hydrogen burning. When the hydrogen is depleted, the stellar core consequently contracts, causing a rise in temperature that ignites hydrogen in a surrounding shell, the “H-Shell Burning” phase. During the contraction, the outer envelope expands and cools, creating a red giant star (e.g., Pagel, 1997). As the object ascends the so-called red giant branch (RGB) (see Fig. 3), convection becomes the dominant energy transport mechanism. The convective envelope eventually reaches down into the H-burning shell, bringing by-products of the CNO cycle to the surface, including ¹³C, in what is termed the “First Dredge-up.” Note that the CNO cycle converts hydrogen into helium with ¹²C as a catalyst. The First Dredge-up sweeps the ¹³C out of the interior and to the stellar surface before it is destroyed in the CNO cycle. It is then ejected into the ISM by mass loss, which starts to become significant on the RGB. Gravitational contraction of the core continues on the RGB, and finally the helium in the core ignites. For stars with mass >4 M _⊙, a “Second Dredge-up” can occur, and subsequent mass loss brings additional ¹³C into the ISM.

When the helium supply is exhausted, the star consists of a carbon core surrounded now by He- and H-burning shells, and the asymptotic giant branch (AGB) begins (e.g., Herwig, 2005). In the He-burning shell, the triple-α process forms copious amounts of ¹²C. As the star continues along the AGB, the star becomes thermally unstable and thermal pulses (TP) develop, which allows the convective zone to reach into the He-burning shell and mix significant amounts of ¹²C into the outer layers of the atmosphere. This mechanism, the so-called “Third Dredge-up,” creates a “carbon-rich” star, such as the object IRC + 10216. In such stars, the surface C/O ∼ 2:1 (Iben and Renzini, 1983; Mowlavi, 1999). AGB stars are in fact one of the few interstellar environments where carbon is more abundant than oxygen, because typically, C/O ∼ 1:2 in the general ISM. The carbon-enhanced circumstellar material is then expelled into the ISM through the increased mass loss that occurs in the AGB stage.

Not all AGB stars end up carbon-rich. It is speculated that more massive AGB stars (>4 M _⊙) can undergo so-called “hot-bottom burning” or HBB (e.g., Herwig, 2005). HBB occurs when the bottom of the convective layer reaches into the active H-burning shell, effectively mixing the by-products of the CNO cycle into the envelope. The main result of HBB is that the ¹²C needed to create a carbon-rich star is converted into nitrogen and ¹³C and effectively destroyed (Frost et al., 1998, and references therein). The AGB star thus reverts back into an O-rich object.

2.2. The circumstellar environment of AGB stars

Mass loss from evolving stars becomes quite significant on the TP-AGB, with typical rates of 10 ⁻⁵ to 10⁻⁴ M _⊙ per year (e.g., Zijlstra, 2006). This material flows from the star and creates a circumstellar envelope. As shown in the diagram in Fig. 4, these envelopes have considerable temperature and density gradients, which foster a rich chemical environment. Near the stellar photosphere (shown in orange), the circumstellar gas is hot (∼1000–2000 K) and dense (∼10¹⁰ cm⁻³), as it has just been driven from the star's surface (see Fig. 4). As this material flows from the star, it cools and becomes more diffuse. As the gas reaches the envelope edge, the temperature drops to ∼25 K, and the density is lowered to ∼10⁴ particles cm⁻³ (e.g., Glassgold, 1996; Crosas and Menten, 1997). Therefore, in the inner envelope (at ∼10 stellar radii or 10 R _*; see Fig. 4), three-body collisions occur that result in equilibrium (or “LTE” for local thermodynamic equilibrium) chemistry and the formation of stable molecules. Dust grains form in this region as well, which helps accelerate the shell material to typical expansion velocities of 10–20 km s⁻¹ as a result of radiation pressure on grain surfaces. In C-rich stars, the dust is thought to be principally silicon carbide in composition, but for more evolved objects, amorphous carbon dominates (Kwok, 2004). Circumstellar dust in carbon stars is about 0.3% of the total mass (Knapp, 1985). The envelope material therefore moves fairly quickly from the star such that much of the LTE chemical composition becomes “frozen” through the intermediate part (white in Fig. 4) of the shell. However, in the outer regions, photochemistry and radical-radical reactions become important, as UV radiation from nearby stars impinges on the envelope.

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FIG. 4.

Graphic representation of a cross section of an AGB circumstellar envelope, plotted as a function of radial distance (log scale) from the center of the star. The star is shown in green (or light gray), the dust acceleration zone in orange (or dark gray), and the remainder of the shell in white with a red (or black) border. Representative temperatures and densities are shown for the dust zone, in which chemistry occurs near thermodynamic equilibrium (“LTE” or “local thermodynamic equilibrium”), and the outer envelope, where photochemistry is dominant. (Color graphics available at www.liebertonline.com/ast)

As discussed, many AGB stars become carbon-rich, and this elemental composition is reflected in the circumstellar shell. A very distinctive gas-phase carbon chemistry occurs in these envelopes, as characterized by long, acetylenic chain–type species of the type HC _n N, C _n H, and C _n N, where n = 1, 2, 3, …. The most studied object of this type is IRC +10216, which lies about 150 pc from Earth and is thought to be representative of C-rich envelopes (see Cernicharo et al., 2000; Tenenbaum et al., 2010a). In this object, the series HCN, HC₃N, HC₅N, HC₇N, and HC₉N is found, as well as the allenic species HCCN and HC₄N (e.g., Guélin and Cernicharo, 1991; Cernicharo et al., 2000). Radicals and negative ions consisting of long carbon chains are also present, including the hydrocarbon series C₂H, C₃H, C₄H, C₅H, and so on, and species with nitrogen (CCN, C₃N, C₅N; C₃N⁻, C₅N⁻: e.g., Thaddeus et al., 2008; Anderson and Ziurys, 2014). Various other C-bearing compounds have been identified as well, as represented by HCN, HCCH, CH₃CN, and CH₃CCH, and some unusual species such as SiC₂, FeCN, and NaCN, containing carbon-silicon and carbon-metal (in the chemists' sense) bonds (e.g., Tenenbaum et al., 2010b; Zack et al., 2011). In fact, in the 1 mm region of the electromagnetic spectrum, the most common molecules are C₄H, SiC₂, and NaCN (Tenenbaum et al., 2010b). The complete list of C-containing molecules in IRC +10216 is given in Table 2. A sample spectrum of this object, showing some of the carbon-bearing compounds present, is shown in Fig. 5.

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FIG. 5.

Representative spectrum of the carbon-rich envelope of the star IRC+10216, showing emission lines originating from various carbon-bearing molecules, measured with the ARO SMT at the frequency of 240 GHz (from Tenenbaum et al., 2010b). “U” indicates an unidentified line.

The more stable molecules such as HCN and HCCH in C-rich shells are thought to be formed by equilibrium chemistry in the inner envelope and thus are “parent” species. Their abundances are not altered through most of the shell as the material expands outward (e.g., Glassgold, 1996). In the outer envelope near∼100–1000 R _*, however, photodissociation of parent species begins to take place, creating radical species that initiate a complex neutral-neutral chemistry that leads to carbon chains. Photodissociation is only important in the outer shell because it can be penetrated by UV radiation from the local ISM. Two major photodissociation processes involve the destruction of HCN and HCCH (Glassgold, 1996): \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{HCN}} + h {\upnu} \to { \rm{CN}} + { \rm{H}} \tag{1} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{HCCH}} + h {\upnu} \to { \rm{CCH}} + { \rm{H}} \tag{2} \end{align*} \end{document}

CN and CCH then react with other neutral species, building to multicarbon chains, for example, \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{CN}} + { \rm{HCCH}} \to { \rm{H}}{{ \rm{C}}_3}{ \rm{N}} + { \rm{H}} \tag{3} \end{align*} \end{document}

Observations have shown that the various chain species have narrow shell-like distributions in the outer envelope, each exhibiting peak abundances at different radii from the star (Bieging and Tafalla, 1993; Anderson and Ziurys, 2014). An example of this phenomenon is shown in Fig. 6. Here, the abundances of four carbon-chain molecules, along with HCN and CN, are plotted as a function of radial distance from the star (Anderson and Ziurys, 2014). The plots were derived by modeling observed line profiles of the molecules. It is clear that the abundance of HCN decreases as that of CN increases, and CN is the dominant outer envelope species. C₃N and HC₃N have similar shell-like distributions and comparable abundances, while CCP is created in the extreme outer shell. C-bearing molecules appear to originate in different, concentric rings, as also noted by Mauron and Huggins (1999) and predicted by theory (Glassgold, 1996).

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FIG. 6.

The outer shell abundances, relative to H₂, of related carbon-chain molecules and their precursors in the envelope of IRC+10216, plotted as a function of the radial distance from the star (log scale), beginning at r∼10¹⁶ cm (from Anderson and Ziurys, 2014). The distributions were obtained by modeling observed spectral line profiles obtained with the ARO SMT and 12 m. The chain molecules appear to be present in successive shells. (Color graphics available at www.liebertonline.com/ast)

Oxygen-rich envelopes of evolved stars support a less complex carbon synthesis. In these objects, carbon is typically represented by CO, HCN, CS, CN, HCO⁺, and HNC (e.g., Ziurys et al., 2009). O-rich stars also appear to have far fewer polyatomic species than their C-rich counterparts, with the most complicated compounds having only 3–4 atoms (NH₃, SO₂, HCN: see Tenenbaum et al., 2010a, 2010b). Dust grains are mostly oxidic in nature (e.g., Lodders and Amari, 2005). Therefore, very active molecular synthesis in an evolved stellar envelope appears to require C > O.

2.3. Into the planetary nebula phase

As shown in Fig. 3, eventually an AGB star expels almost all its initial mass through envelope loss and subsequently becomes a hot, UV-emitting white dwarf of approximately 0.5 M _⊙ (e.g., Kwok, 2000). Nucleosynthesis ceases at this juncture. The former circumstellar envelope continues to flow from the star and is partially ionized by the UV radiation, first creating a short-lived “protoplanetary” nebula and then a “planetary nebula”—a bright, colorful object characterized by high-energy atomic emission lines (with nothing to do with planets). CRL 2688 and the Helix are classic examples of these two types of objects (see Fig. 3). The common perception is that the molecular gas from the previous AGB envelope is mostly destroyed via photodissociation in planetary nebulae (PNe). Such molecular destruction is predicted by chemical models of these objects (e.g., Ali et al., 2001; Redman et al., 2003; Kimura et al., 2012). Molecular material in PNe should therefore be severely depleted as a function of time after leaving the AGB. Because PNe supply over 85% of the matter forming the ISM (e.g., Dorschner and Henning, 1995), the exact fate of the carbon-rich molecular material from C-rich envelopes is of interest.

However, very recent molecular studies of PNe show that, even in their very late stages, polyatomic molecular material survives, much of it containing carbon. Evidence for this new paradigm comes in part from a series of studies focusing on the Helix Nebula (NGC 7293), the oldest PN known at 12,000 years. This nebula has a large spatial extent on the sky. At first, various C-bearing molecules were found only at one position near the western edge of the nebula; the list of species present is composed of CO, HCN, CN, HCO⁺, HNC, H₂CO, CCH, and c-C₃H₂ (Bachiller et al., 1997; Tenenbaum et al., 2009). Subsequently, both HCO⁺ and H₂CO were detected at eight additional positions, sampled across the whole of the nebula (Zack and Ziurys, 2013), followed by HCN and CCH (Schmidt and Ziurys, 2016). The J = 1 → 0 transition of HCO⁺ was additionally mapped across the entire nebula, as traced by the optical image, and was found to be present virtually everywhere, as displayed in Fig. 7. The line profiles of this molecule varied with location as well, indicating a complex velocity structure within the nebula. Fractional abundances of HCO⁺ and H₂CO relative to H₂ were estimated to be f(HCO⁺/H₂) ∼0.3–7.3 × 10⁻⁸ and f(H₂CO/H₂) ∼0.3–2.1 × 10⁻⁷—several orders of magnitude higher than predicted by any chemical model (Zack and Ziurys, 2013).

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FIG. 7.

Spectra of the J = 1 → 0 transition of HCO⁺ observed across the Helix Nebula, the oldest known PN, superimposed over the optical image of the object (from Zeigler et al., 2013). The data show that HCO⁺ is present where bright, ionized atomic gas also exists.

Subsequent observations have shown the presence of carbon-bearing molecules such as HCN, H₂CO, HCO⁺, and CCH in numerous other nebulae, including K4-47, M2-48, NGC 6537, K3-58, M1-7, M4-14, M3-28, M3-55, NGC 6720, and NGC 6853 (e.g., Edwards et al., 2014; Schmidt and Ziurys, 2016). Sample spectra of the CCH radical in PNe are shown in Fig. 8 (Schmidt and Ziurys, 2017). Clearly, carbon chemistry remains quite active in PNe. Observational evidence suggests that the C-bearing compounds arise from remnant AGB shell molecular material and in situ nebular chemistry (e.g., Edwards et al., 2014). The density of the molecular gas in PNe is typically ∼10⁵ cm⁻³ and therefore high enough to foster chemical reactions (Zack and Ziurys, 2013; Edwards et al., 2014; Schmidt and Ziurys, 2016). The molecules are also likely mixed with dust, which aids in their preservation by providing shielding from UV radiation.

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FIG. 8.

Spectra of the N = 3 → 2 transition of the CCH radical observed with the ARO SMT toward four different PNe spanning the age range 500–10,000 years. The spin-rotation doublet structure is apparent in the spectra (from Schmidt and Ziurys, 2017).

Current chemical models of PNe chiefly are concerned with molecule destruction, not production. However, photoionization does initiate ion-molecule reactions that can lead to some polyatomic species. For example, Kimura et al. (2012) suggested that HCO⁺ and HCN could be synthesized in these nebulae via the following pathways: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{O}}^ + } + {{ \rm{H}}_2} \to { \rm{HC}}{{ \rm{O}}^ + } + { \rm{H}} \tag{4} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{HCN}}{{ \rm{H}}^ + } + {{ \rm{e}}^ - } \to { \rm{HCN}} + { \rm{H}} \tag{5} \end{align*} \end{document}

However, the HCN abundance these authors derive is a factor of 100–1000 less than those observed, although that for HCO⁺ is comparable. The gas-phase carbon chemistry of PNe is a topic that clearly needs further exploration, both observationally and theoretically.

Carbon also appears to be present in some PNe in what could be termed macromolecules. During the early PN phase, features appear in the IR spectra of these objects between 3 and 14 microns in wavelength, which are thought to arise from both carbon aromatic and aliphatic stretching and bending vibrations. Some are characteristic of aliphatic side groups attached to aromatic rings (Kwok, 2004). These spectra suggest the existence of large carbonaceous macromolecular material in PNe, perhaps with a mixed aromatic/aliphatic structure as shown in Fig. 9, from Kwok and Zhang (2011), or large ionized polyaromatic hydrocarbon (PAH) species (e.g., Jones, 2009). It is also interesting to note that the fullerene species C₆₀ and C₇₀ have now been detected in a few younger PNe (e.g., Zhang and Kwok, 2013).

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FIG. 9.

A proposed carrier of the unidentified infrared bands (UIBs), which are found in the spectra of young PNe (from Kwok and Zhang, 2011). The UIBs contain both aromatic and aliphatic carbon features apparently arising from macromolecular material, as suggested in the figure.

3.1. From planetary nebulae to diffuse clouds

The matter ejected from PNe becomes part of the diffuse ISM and ultimately is incorporated into “diffuse clouds.” These clouds are categorized by low densities (∼10–100 cm⁻³) and warm temperatures (T ∼ 30–100 K); the ambient stellar UV radiation field is somewhat attenuated so that the H₂ content is >0.1 of the total available hydrogen, although the gas-phase carbon is predominantly in the form of C⁺ (e.g., Snow and McCall, 2006). Diffuse clouds were for decades thought to be relatively devoid of molecules, except for simple diatomic species. However, this picture has changed over the past 10 years with the discovery of numerous polyatomic species, including carbon-bearing compounds, as summarized in Table 3 (Liszt et al., 2006, 2008; Snow and McCall, 2006). As shown in the table, carbon is represented in C₃, H₂CO, CCH, HCN, HNC, HCO⁺, and c-C₃H₂, as well as diatomic species. Very recently, C₆₀ ⁺ has been identified in diffuse gas as well (Campbell et al., 2015).

The origin of the polyatomic molecules is unknown, because diffuse clouds are too low in density to produce them in situ. Molecule synthesis is again driven by photochemistry in these objects, with ion-molecule and neutral-neutral reactions playing important roles. CO and HCO⁺ are thought to be created from C⁺ (Snow and McCall, 2006): \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}^ + } + { \rm{OH}} \to { \rm{C}}{{ \rm{O}}^ + } + { \rm{H}} \tag{6} \end{align*} \end{document}

CO⁺ then leads to HCO⁺ via Reaction 4, and CO forms from the dissociative electron recombination of HCO⁺. The synthesis of molecules such as H₂CO and HCN remains a mystery. However, the widespread distribution of HCO⁺, HCN, CCH, and H₂CO in the Helix Nebula, the oldest PNe known, and the presence of C₆₀ in PNe, suggest that carbon-containing molecular material is ejected into the diffuse ISM. In fact, there is a remarkable similarity between the molecules observed in PNe relative to diffuse clouds, as shown in Table 4; and the relative abundances are not incompatible, particularly for C-bearing species. Typically, molecular abundances in diffuse clouds are a factor of 10–100 less than those in older PNe, as shown in the table. Furthermore, molecular ratios appear to be similar. The HCN/HNC ratio is about 1–2 in PNe and about 5 in diffuse clouds. The CO/HCO⁺ ratio is typically ∼10³ in both types of sources, while CN/HCN ∼ 4–9 in PNe and CN/HCN ∼ 7 in diffuse clouds. These results are consistent with the scenario of molecular material being injected into the diffuse ISM from PNe, with some fraction surviving as the matter collapses into diffuse clouds. Calculations suggest that a clump with n ∼ 3 × 10⁵ cm⁻³ will dissipate into diffuse material with A_v  = 1 in about 10⁷ years (e.g., Tenenbaum et al., 2009).

Significantly larger organic species may also be present in diffuse clouds, perhaps containing tens to hundreds of atoms. The recent discovery of C₆₀ ⁺ in these regions supports this hypothesis (Campbell et al., 2015). Evidence for these as-yet-unidentified molecules comes from the observation of the same 3–14 micron IR bands observed in PNe, as discussed, as well as broad spectra measured in the optical region, the “diffuse interstellar bands,” or DIBS. Some theories to the origin of both types of features include PAHs, long carbon chains, and fullerenes (e.g., Snow and McCall, 2006). With the exception of C₆₀ ⁺, the exact identities of both the IR and optical features remain a mystery. They likely are related to the macromolecules found in PNe. Based on elemental depletions measured in diffuse clouds, about 50% of the carbon is in the gas phase as atoms; the remainder is in molecules and dust (e.g., Palme and Jones, 2005).

3.2. The transition to dense clouds

Dense clouds, characterized by gas densities near 10⁴ to 10⁶ cm⁻³, are thought to be created from diffuse clouds, either by collapse or by accretion caused by shock waves (e.g., Inoue and Inutsuka, 2012). These objects are typically colder than their diffuse counterparts, with T ∼ 10–50 K. Because they are dense, these sources are rich in molecular material and are usually referred to as “molecular clouds.” Dense clouds contain a large fraction of the mass in our galaxy, as much as half the total gas mass in the inner 8–9 kpc. Most of the clouds are present in an annulus between 3.5 and 7.5 kpc from the Galactic Center, the so-called “molecular ring” (e.g., Ferrière, 2001). Dense clouds have a wide range of masses. The largest, “giant molecular clouds,” or GMCs, contain between 10³ and 10⁶ M _⊙ and are 10–100 pc in size. These objects are typically the sites of both high- and lower-mass star formation, as well as planetary system creation. There are also less massive dense clouds, such as Bok globules, which have 2–50 M _⊙ and are no larger than ∼0.3 pc across. Very cold clouds (T ≤ 10 K) are sometimes referred to as “dark clouds.” Evidence of star formation has been seen in all types of dense clouds (e.g., Stutz et al., 2010).

The chemistry of carbon reaches a new degree of complexity in these objects, in particular for the GMCs. Oxygen plays a major role in the synthesis, producing compounds that are typical of the organic laboratory. As shown in Fig. 10, alcohols (CH₃OH, CH₃CH₃OH), carboxylic acids (HCOOH, CH₃COOH), aldehydes (CH₃CHO), ketones (CH₃COCH₃), ethers (CH₃OCH₃), and even esters (CH₃OCHO) are represented in the chemical inventory of GMCs. There are also a few non-aromatic ring species, c-C₃H₂ and CH₂OCH₂, and a simple sugar, glycolaldehyde, CHOCH₂OH. Such large molecules and their D and ¹³C isotopologues help generate the very complex spectra observed in GMCs. An example spectrum is shown in Fig. 11, which displays data at the frequency of 142.6 GHz, observed toward Sgr B2(N), a very large GMC in the center of the Galaxy. The density of spectral lines is so high in GMC data that the “confusion limit” is now being reached; that is, there are continuous spectral features and no obvious baseline. The plethora of lines brings enhanced difficulty to the identification of new molecules.

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FIG. 10.

Common organic functionality observed in interstellar molecules found in dense clouds. Alcohols, ketones, carboxylic acids, etc. are commonly observed in interstellar gas. (Color graphics available at www.liebertonline.com/ast)

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FIG. 11.

A typical spectrum exhibited by dense molecular clouds with star formation, in this case the Sgr B2(N) cloud, observed with the ARO 12 m. The center frequency is 142.6 GHz. The molecular carriers of the spectral features are labeled, with those that are unidentified marked as “U.” A variety of organic-type compounds are present in this spectrum, including CH₃CH₂CN, NH₂CHO, CH₂CHCN, and HCOOCH₃. (Color graphics available at www.liebertonline.com/ast)

There are numerous models of the chemistry of carbon-bearing molecules and related species, usually involving thousands of chemical reactions; for example, see the work of Garrod and Herbst (2006) and references therein. The gas-phase chemistry of carbon is chiefly thought to involve ion-molecule reactions, along with neutral-neutral processes and radiative association pathways (e.g., Herbst and Millar, 2008). Radiative association is particularly important for carbon, as the simplest ion-molecule reaction of this element, C⁺ + H₂, is endothermic. Thus, synthesis involving carbon is initiated by the association reactions: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}^ + } + {{ \rm{H}}_2} \to { \rm{C}}{{ \rm{H}}_2}^ + + h {\upnu} \tag{7} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}} + {{ \rm{H}}_2} \to { \rm{C}}{{ \rm{H}}_2} + h {\upnu} \tag{8} \end{align*} \end{document}

The rate of the first process has been calculated to be 2 × 10⁻¹⁶(T/300)^−1.3 cm³ s⁻¹ in the range 10–300 K, but few studies have been conducted for the second pathway. A third initiation reaction is also thought to be important: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}} + {{ \rm{H}}_3}^ + \to { \rm{C}}{{ \rm{H}}^ + } + {{ \rm{H}}_2} \tag{9} \end{align*} \end{document}

These three pathways serve to bring carbon into molecular form (Herbst and Millar, 2008). The larger species are then produced via carbon “insertion,” condensation, and radiative association of heavier fragments. The insertion and condensation processes can occur either with ion-molecule or neutral-neutral reactions. Insertion can take place with either C⁺ or C, building larger hydrocarbons: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}^ + } + { \rm{C}}{{ \rm{H}}_4} \to {{ \rm{C}}_2}{{ \rm{H}}_3} ^+ + { \rm{H \ or }} \ {{ \rm{C}}_2}{{ \rm{H}}_2} ^ + + {{ \rm{H}}_2} \tag{10} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}} + {{ \rm{C}}_2}{{ \rm{H}}_2} \to c \hbox{-} {{ \rm{C}}_3}{ \rm{H}} + { \rm{H \ or }} \ {{ \rm{C}}_3} + {{ \rm{H}}_2} \tag{11} \end{align*} \end{document}

Note that such reactions have more than one product channel and that the branching ratios between these pathways are often not known. Condensation also builds larger organic species, including those with a heteroatom (also see Eq. 3): \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3} ^ + + { \rm{C}}{{ \rm{H}}_4} \to {{ \rm{C}}_2}{{ \rm{H}}_5} ^ + + {{ \rm{H}}_2} \tag{12} \end{align*} \end{document}

Radiative association can lead to larger polyatomic species as well: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3} ^ + + { \rm{N}}{{ \rm{H}}_3} \to { \rm{C}}{{ \rm{H}}_3}{ \rm{N}}{{ \rm{H}}_3} ^ + + h {\upnu} \tag{13} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3} ^ + + { \rm{CO}} \to { \rm{C}}{{ \rm{H}}_3}{ \rm{C}}{{ \rm{O}}^ + } + h {\upnu} \tag{14} \end{align*} \end{document}

The product ions such as CH₃NH₃ ⁺ and CH₃CO⁺ then lead to the neutral species by dissociative electron recombination, in this case creating CH₃NH₂ and CH₂CO.

Reactions on grains can also involve carbon, but such processes are even less well understood and more difficult to accurately model (e.g., Herbst and van Dishoeck, 2009). There is evidence that CO, once frozen onto grain surfaces, can become hydrogenated, leading to methanol. Hydrogenation of HCN is also speculated to lead to CH₂NH and CH₃NH₂ (e.g., Watanabe and Kouchi, 2002). However, observations suggest that these two molecules may be created by neutral-neutral processes in the gas phase (Halfen et al., 2013). Most recent models include both gas-phase and surface chemistry, as well as variations such as grain “freeze-out” and subsequent evaporation (e.g., see Herbst and Millar, 2008, and references therein and below).

Carbon is of course also contained in dust grains in dense clouds, primarily in amorphous form (e.g., Jones and Nuth, 2011). There seems to be an evolution from H-rich hydrocarbon solids to more aromatic material. Such carbon grains are easily destroyed in shocks, however, such that toward jets and winds, most of the carbon is in the gas phase (e.g., Jones, 2009). It is thought that the carbonaceous grains readily reform through accretion onto surviving grains and silicates. Dust grains in dense clouds also readily acquire icy mantles with water, CO, CO₂, and methanol, and grain-grain collisions may lead to grain growth (e.g., Jones, 2001).

The majority of dense cloud models begin with atoms as the initial constituents. However, if diffuse clouds already contain significant amounts of molecular material, then a more realistic starting point might be simple polyatomic molecules. Certainly this approach is more consistent with observations and could lead to the creation of ever more complex species over the 10⁶ to 10⁷ lifetime of a GMC. This avenue of modeling is a subject of future work.

4.1. Prestellar and protostellar cores

Formation of stars usually starts in compact cores inside dense clouds, which have molecular densities in the range 10⁴ to 10⁶ cm⁻³. Typical core masses are found to lie between 0.5 and 5 M _⊙. Other star formation sites are the small (diameter < several parsecs; mass < 50 M _⊙), dark, spherical Bok globules. Molecular hydrogen is the dominant species in all these environments. In such cores, carbon-containing molecules like CO, CH₃OH, and CH₄ freeze out on interstellar grains. This process reduces the chemical abundance of these species in the gas phase. When star formation begins, the increase in temperature caused by the embedded young stellar object leads to re-sublimation of the frozen-out molecules, which triggers a “second generation” chemistry that leads to the formation of fairly complex carbon-containing molecules (Brown et al., 1988). For example, larger organic molecules like acetylene can be formed from methane through carbon insertion (see Eq. 10). The C₂H₂ ⁺ ions formed from this process can further react with molecular hydrogen through radiative association: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}_2}{{ \rm{H}}_2} ^ + + {{ \rm{H}}_2} \to {{ \rm{C}}_2}{{ \rm{H}}_4}^ + + h {\upnu} \tag{15} \end{align*} \end{document}

Subsequent dissociative recombination of C₂H₄ ⁺ then leads to acetylene; the C₂H₃ ⁺ also created in Eq. 10 can generate HCCH or CCH (Kalhori et al., 2002; Ehlerding et al., 2004). The intermediate ions C₂H₂ ⁺ and C₂H₃ ⁺ formed in these processes can further react with methane to form larger ions, for example, \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}_2}{{ \rm{H}}_2} ^ + + { \rm{C}}{{ \rm{H}}_4} \to {{ \rm{C}}_3}{{ \rm{H}}_5}^ + + { \rm{H}} \tag{16} \end{align*} \end{document}

which by dissociative recombination can yield cyclic species such as c-C₃H₂ (Mookerjea et al., 2012). In certain instances, it is thought that sublimated methane enhances carbon-chain chemistry in a warmer environment (e.g., Sakai and Yamamoto, 2013).

Very often radiative association leads to protonated species, followed by dissociative recombination, which have been invoked to explain the formation of complex organic compounds like alcohols, esters, and ethers encountered in star-forming regions, for example, \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3} ^ + + {{ \rm{H}}_2}{ \rm{O}} \to { \rm{C}}{{ \rm{H}}_3}{ \rm{O}}{{ \rm{H}}_2} ^ + + h {\upnu} \tag{17} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3}{ \rm{O}}{{ \rm{H}}_2} ^ + + {{ \rm{e}}^ - } \to { \rm{C}}{{ \rm{H}}_3}{ \rm{OH}} + { \rm{H}} \tag{18} \end{align*} \end{document}

However, recent experiments have demonstrated that only small branching ratios occur for the production of alcohols and ethers upon dissociative recombination. Fragmentation into small radicals and fracturing of C-C and C-O bonds are very common (Geppert et al., 2006; Hamberg et al., 2010). In some cases (e.g., dimethyl ether and formic acid), organic molecules can be formed through ion-molecule reactions followed by dissociative recombination (Garrod and Herbst, 2006): \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3}{ \rm{O}}{{ \rm{H}}_2} ^ + + { \rm{C}}{{ \rm{H}}_3}{ \rm{OH}} \to { \rm{C}}{{ \rm{H}}_3}{ \rm{OHC}}{{ \rm{H}}_3} ^ + + {{ \rm{H}}_2}{ \rm{O}} \tag{19} \end{align*} \end{document} \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{C}}{{ \rm{H}}_3}{ \rm{OHC}}{{ \rm{H}}_3} ^ + + {{ \rm{e}}^ - } \to { \rm{C}}{{ \rm{H}}_3}{ \rm{OC}}{{ \rm{H}}_3} + { \rm{H}} \tag{20} \end{align*} \end{document}

Such reaction sequences, however, appear to be unsuccessful in producing certain other molecules in their observed abundances, such as methanol. Processes on grain surfaces, such as the successive hydrogenation of CO, as mentioned, or neutral-neutral reactions in the gas phase, may be responsible for the formation of methanol. One major challenge for surface production of CH₃OH is the mechanism for grain desorption, which is difficult at the low temperatures characteristic of dark clouds (Geppert et al., 2006).

Gas-phase radical-neutral and atom-neutral reactions are also involved in the formation of carbon compounds in star-forming regions, as represented by Eq. 10. Another example is the formation of HC₃N via the radical-neutral process: \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{ \rm{CN }} + {{ \rm{C}}_2}{{ \rm{H}}_2} \to { \rm{H}}{{ \rm{C}}_3}{ \rm{N}} + { \rm{H}} \tag{21} \end{align*} \end{document}

Such reactions often have considerable rates (∼10⁻¹⁰ to 10⁻¹¹ cm³ s⁻¹), because one of the reactants is an unstable radical species.

A universal (grain surface or gas-phase) mechanism for generating complex molecules in star-forming regions cannot really be formulated. The formation pathway of each species must be explored individually. However, mapping the formation routes of complex species in star-forming regions has important consequences for biochemical evolution.

4.2. Collapse to protostellar disks

Young stars are usually surrounded by rotating protoplanetary disks, in which dust grains coagulate into planetesimals and then further assemble into planets and asteroids. Protoplanetary disks are transient features with lifetimes of 3,000,000–10,000,000 years, composed of mostly of gas and dust. From a simplistic viewpoint, they consist of a warm inner zone (the site of planet formation) and a cold outer zone (e.g., Henning and Semenov, 2013). In the perpendicular direction, they can be divided into a dense midplane (where planets form and molecules are frozen out onto grains, obscured from observations), the surface (where chemistry is dominated by photons), and the intermediate (warm molecular) zone in between. Several carbon-containing molecules and ions have so far been identified toward protoplanetary disks: H₂CO, CS, C₂H, c-C₃H₂, HCN, HNC, CN, DCN, HC₃N, HCO⁺, and DCO⁺.

Ultraviolet radiation from the star and the interstellar background penetrates the surface zone and photodissociates/photoionizes most molecules. Gas-phase molecular abundances thus tend to peak in the intermediate zone, where molecules are desorbed from the grains but not yet photolyzed (Aikawa and Herbst, 1999). Models of protoplanetary disks predict that C⁺ is the prevalent ion close to the disk surface, whereas over the whole disk HCO⁺ is the most abundant ionic species. In the outer, cooler regions of protoplanetary disks, the chemistry is driven by high-energy radiation and cosmic rays (Harada et al., 2010) and is dominated by ion-molecule reactions followed by dissociative recombination as the terminal step leading to stable neutrals, among other by-products. This type of reaction sequence also leads to carbon-chain molecules (C _n H), which are probably formed by the dissociative recombination of C _n H₂ ⁺ ions (Markwick et al., 2002): \documentclass{aastex}\usepackage{amsbsy}\usepackage{amsfonts}\usepackage{amssymb}\usepackage{bm}\usepackage{mathrsfs}\usepackage{pifont}\usepackage{stmaryrd}\usepackage{textcomp}\usepackage{portland, xspace}\usepackage{amsmath, amsxtra}\pagestyle{empty}\DeclareMathSizes{10}{9}{7}{6}\begin{document} \begin{align*}{{ \rm{C}}_n}{{ \rm{H}}_2}^ + + {{ \rm{e}}^ - } \to {{ \rm{C}}_n}{ \rm{H}} + { \rm{H}} \tag{22} \end{align*} \end{document}

Deep in the interior of the outer disk, however, most molecular species freeze out onto grain surfaces. In the densest parts of the inner disk, where high temperatures (100–5000 K) and high densities (up to 10¹² cm⁻³) prevail and ionizing radiation is attenuated, neutral-neutral reactions with barriers and 3-body processes contribute to the chemistry (Willacy et al., 1998; Furuya et al., 2013). Also, surface reactions play a decisive role in the synthesis of molecules in protoplanetary disks (e.g., Furuya and Aikawa, 2014; Walsh et al., 2014). HNCO is thought to be produced on grain surfaces from OCN and H, and methanol from CH₂OH and H. Observations have also shown that grain growth and dust settling occurs in protoplanetary disks, with the gas-to-dust ratio increasing by at least a factor of 10 relative to the ISM (e.g., Horne et al., 2012). Such grain growth likely leads to planetesimals. Although crystalline silicates are readily identified in disks via their IR signatures, carbonaceous dust is harder to identify. Its presence in these objects is inferred from modeling (e.g., Apai and Lauretta, 2010). The rich chemistry in protoplanetary disks can exert a great influence on the composition of planets, their atmospheres, asteroids, and comets (Aikawa and Herbst, 1999; Henning and Semenov, 2013). Further studies with the Atacama Large Millimeter/submillimeter Array (ALMA) will help probe the connection between disk and planetesimal chemistry.