CARBON-RICH GIANT PLANETS: ATMOSPHERIC CHEMISTRY, THERMAL INVERSIONS, SPECTRA, AND FORMATION CONDITIONS

We determine volume mixing ratios of the dominant molecular species using the approach described in Madhusudhan & Seager (2011). We consider hydrogen-dominated atmospheres with temperatures typical of known transiting extrasolar giant planets (T_eff ∼ 700–3000 K) and with different elemental abundances. We compute the molecular mixing ratios generally under the assumption of thermochemical equilibrium, considering the effects of non-equilibrium thermochemistry where desirable. We use the equilibrium chemistry code originally developed in Seager & Sasselov (2000), and subsequently updated and/or used in several recent works (Seager et al. 2005; Kuchner & Seager 2005; Miller-Ricci et al. 2009; Madhusudhan & Seager 2011). The code computes gas-phase molecular mixing ratios for 172 molecules, resulting from abundances of 23 atomic species, by minimizing the net Gibbs free energy of the system. The multi-dimensional Newton–Raphson method described in White et al. (1958) is used for the minimization. We adopt polynomial fits for the free energies of the molecules based on Sharp & Huebner (1990).

For the purpose of the present study, the major chemically and spectroscopically dominant species included in our model atmospheres are H₂O, CO, CH₄, and CO₂, in addition to H₂ which contributes H₂–H₂ collision-induced opacity. Equilibrium mixing ratios of these molecules are computed for different elemental abundances. Specifically, we compare the abundances of the major species H₂O, CO, and CH₄ for different C/O ratios, spanning oxygen-rich (C/O < 1) and carbon-rich (C/O > 1) compositions over a wide range in pressure–temperature (P–T) space. We then identify the temperature regimes in which C-rich atmospheres are chemically distinct, and identifiable in spectra, from O-rich atmospheres. A similar analysis for two C/O ratios (0.5 and 1.01) and for the specific case of Neptune-mass carbon planets was reported earlier by Kuchner & Seager (2005).

The concentrations of species, particularly CO and CH₄, can be driven out of equilibrium by transport rapid enough to move material out of a given pressure and temperature regime in a time short compared to the chemical equilibration time (Yung & Demore 1999). Vertical mixing is known to be prevalent in atmospheres of brown dwarfs (Noll et al. 1997; Saumon et al. 2003, 2006), giant planets in the solar system (Prinn & Bashay 1977; Yung et al. 1988; Fegley & Lodders 1994; Visscher et al. 2010), and in extrasolar giant planets (Stevenson et al. 2010; Madhusudhan & Seager 2011). The effects of non-equilibrium chemistry are by their nature more pronounced with decreasing temperature, implying that atmospheres of the cooler population of exoplanets (e.g., GJ 436b, GJ 1214b, HD 189733b, etc.) are more susceptible to CO–CH₄ disequilibrium (Zahnle et al. 2009; Line et al. 2010; Moses et al. 2011; Miller-Ricci Kempton et al. 2011). Consequently, in this study we investigate the effects of CO–CH₄ disequilibrium due to vertical mixing on the observable CO abundances in some special cases at low temperature.

The set of reactions governing chemical equilibrium between the dominant carbon-bearing species CH₄ and CO is summarized by

$$\begin{equation} {\rm CO + 3H_2 \rightleftharpoons CH_4 + H_2O}. \end{equation} \tag{ 1 }$$

At high pressures and temperatures in the deeper layers of the atmospheres (for P ≳ 1 bar), the reaction timescales are short and the species approach chemical equilibrium. Higher in the atmosphere where temperatures are lower, turbulent processes can cause timescales of vertical mixing to be shorter than the chemical timescales, thereby driving the atmosphere out of equilibrium. Our approach for computing the non-equilibrium concentrations of CO and CH₄ is described in Madhusudhan & Seager (2011; see also Visscher & Moses 2011).

A thermal inversion, like the Earth's stratosphere, is a region in the atmosphere where temperature increases outward (see Madhusudhan & Seager 2010 for a review). Observations of several hot Jupiters have been interpreted as containing evidence for thermal inversions in their dayside atmospheres (e.g., Burrows et al. 2008; Knutson et al. 2008, 2009; Machalek et al. 2009; Christiansen et al. 2010; Madhusudhan & Seager 2009, 2010). Thermal inversions in hot Jupiters have been hypothesized as caused by strong optical absorption due to TiO and VO at high altitudes (Hubeny et al. 2003; Fortney et al. 2008; but cf. Zahnle et al. 2009; Knutson et al. 2010).

Based on equilibrium abundances of TiO and VO, assuming solar elemental abundances, Fortney et al. (2008) suggested that hot Jupiters with high irradiation levels would likely be hot enough to host TiO and VO in gaseous form leading to thermal inversions. In contrast, Spiegel et al. (2009) showed that TiO/VO, being heavy molecules, are prone to gravitational settling, and would therefore be deficient at high altitudes in the atmosphere. The amount of TiO required to explain thermal inversions would have to result from vigorous vertical mixing in the atmosphere. Even then, the presence of possible cold traps can cause further depletion of TiO via condensation. VO is much less abundant than TiO. Consequently, TiO/VO might be able to cause inversions in only the extreme end members even amongst the very hot Jupiters of the pM class (Spiegel et al. 2009). However, the recent non-detection of a strong thermal inversion in one of the most highly irradiated exoplanets known, WASP-12b (Madhusudhan et al. 2011), poses a new challenge to the TiO hypothesis.

In this work, we investigate the cause behind the absence of a strong thermal inversion in WASP-12b. We compute abundances of TiO and VO in chemical equilibrium over a grid in temperature and pressure space representative of highly irradiated atmospheres. We investigate the influence of C/O ratio on the TiO and VO abundances by computing the TiO and VO abundances at different C/O ratios spanning oxygen-rich and carbon-rich regimes, i.e., solar C/O of 0.54 (as assumed by Fortney et al. 2008) and C/O ⩾1. By comparing the TiO and VO abundances obtained in the different regimes, we assess the likelihood of thermal inversions in high-temperature carbon-rich atmospheres, such as that of WASP-12b, versus those in oxygen-rich atmospheres.

We investigate observable signatures of CRG atmospheres over a representative range in effective temperature encompassing presently known transiting giant exoplanets. We report model spectra of dayside thermal emission, which can be observed for transiting planets at secondary eclipse. We compute the model spectra using a one-dimensional line-by-line radiative transfer code for exoplanetary atmospheres developed in Madhusudhan & Seager (2009). Given a P–T profile and the molecular abundance profiles, the code computes the emergent flux by solving the radiative transfer equation along with constraints of local thermodynamic equilibrium (LTE), hydrostatic equilibrium, and global energy balance. The model includes the major molecular and continuum opacity sources. We assume H₂-rich atmospheres, and we include opacities due to H₂O, CO, CH₄, CO₂, NH₃, TiO, VO, and H₂–H₂ collision-induced absorption. Our molecular line data are from Freedman et al. (2008), R. S. Freedman (2009, private communication), Rothman et al. (2005), Karkoschka & Tomasko (2010), and E. Karkoschka (2011, private communication). We obtain the H₂–H₂ collision-induced opacities from Borysow et al. (1997) and Borysow (2002). In order to compute the planet–star flux ratios, we divide the planetary spectrum by a Kurucz model of the stellar spectrum derived from Castelli & Kurucz (2004).

We compute thermal spectra over a nominal range of irradiation levels (or temperatures), planet types, and C/O ratios. We consider four representative irradiation and planetary regimes: a very hot Jupiter, a hot Jupiter, a hot Neptune, and a warm Neptune. The four chosen planet types have equilibrium temperatures (T_eq) of 2300 K, 1500 K, 1000 K, and 750 K, and we adopt P–T profiles motivated by those published in the literature using our present modeling approach (e.g., Madhusudhan & Seager 2009, 2010, 2011; Madhusudhan et al. 2011). For each case, we adopt the stellar and planetary properties of a known transiting exoplanetary system: hot Jupiter (HD 189733b; Bouchy et al. 2005), very hot Jupiter (HD 189733b with stellar irradiation scaled up), hot Neptune (HAT-P-11b; Bakos et al. 2010), warm Neptune (GJ 436b; Gillon et al. 2007). For each system, we compute mixing ratios of the major molecular species using the prescription described in Section 2.1, for C/O ratios of solar (0.54), 1, and 3. For the atomic abundances, we fix the O/H ratio at solar (Asplund et al. 2005), and we change the C/H ratio to obtain the corresponding C/O ratios. We then compute thermal spectra for each combination of P–T profile and composition using our radiative transfer model described above.

Close-in giant planets are thought to have originated in the cold outer regions of protoplanetary disks and migrated inward until they stopped at closer orbital radii (Goldreich & Tremaine 1980; Fogg & Nelson 2005, 2007). Here, we assume that the giant planet was formed via core accretion. We calculate the composition of the icy planetesimals accreted by the forming planet following the approach developed in Mousis et al. (2009b, 2011). In this scenario, building blocks accreted by the protoplanet may have formed all along its radial migration pathway in the protoplanetary disk. However, in this work, we assume that only the planetesimals produced beyond the snow line, i.e., those possessing a significant fraction of volatiles, materially affected the observed O and C abundances due to their vaporization when they entered the envelope of the planet (Mousis et al. 2009b, 2011). This hypothesis is supported by the work of Guillot & Gladman (2000) who showed that planetesimals delivered to a planet with a mass similar to or greater than that of Jupiter are ejected rather than accreted. This mechanism should then prevent further noticeable accretion of solids by the planet during its migration below the snow line.

We demonstrate our model for the specific case of WASP-12b. The composition of planetesimals accreted by WASP-12b can be inferred from the formation sequence and the composition of the different ices formed beyond the snow line of the protoplanetary disk. This model has been used to interpret the observed volatile enrichments in the atmospheres of Jupiter and Saturn in a way consistent with the heavy element content predicted by interior models (Mousis et al. 2009a) and also the apparent carbon deficiency observed in the hot Jupiter HD 189733b (Mousis et al. 2009b, 2011). It is based on a predefined initial gas-phase composition in which elemental abundances reflect those of the host star and describes the process by which volatiles are trapped in icy planetesimals formed in the protoplanetary disk. Oxygen, carbon, nitrogen, sulfur, and phosphorus are postulated to exist only in the form of H₂O, CO, CO₂, CH₃OH, CH₄, N₂, NH₃, H₂S, and PH₃. We set CO/CO₂/CH₃OH/CH₄ = 70/10/2/1 in the gas phase of the disk, values that are consistent with the interstellar medium (ISM) measurements considering the contributions of both gas and solid phases in the lines of sight (Frerking et al. 1982; Ohishi et al. 1992; Ehrenfreund & Schutte 2000; Gibb et al. 2000). In addition, S is assumed to exist in the form of H₂S, with H₂S/H₂ = 0.5 × (S/H₂)_☉, and other refractory sulfide components (Pasek et al. 2005). We also consider N₂/NH₃ = 1/1 in the disk's gas phase. This value is compatible with thermochemical calculations in the solar nebula that take into account catalytic effects of Fe grains on the kinetics of N₂ to NH₃ conversion (Fegley 2000). In the following, we adopt these mixing ratios in our model of the protoplanetary disk. Once the abundances of these molecules have been fixed, the remaining O gives the abundance of H₂O.

The process of volatile trapping in planetesimals, illustrated in Figure 1, is calculated using the equilibrium curves of hydrates, clathrates, and pure condensates, and the ensemble of thermodynamic paths detailing the evolution of temperature and pressure in the 5–20 AU range of the protoplanetary disk. We refer the reader to the works of Papaloizou & Terquem (1999) and Alibert et al. (2005) for a full description of the turbulent model of accretion disk used here. This model postulates that viscous heating is the predominant heating source, assuming that the outer parts of the disk are protected from solar irradiation by the shadowing effect of the inner disk parts. In these conditions, the temperature in the planet-forming region can decrease down to very low values (∼20 K; Mousis et al. 2009a).

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For each ice considered in Figure 1, the domain of stability is the region located below its corresponding equilibrium curve. The clathration process stops when no more crystalline water ice is available to trap the volatile species. The equilibrium curves of hydrates and clathrates derive from the compilation of published experimental work by Lunine & Stevenson (1985), in which data are available at relatively low temperatures and pressures. On the other hand, the equilibrium curves of pure condensates used in our calculations derive from the compilation of laboratory data given in the CRC Handbook of Chemistry and Physics (Lide 2002). For each ice considered, the domain of stability is the region located below its corresponding equilibrium curve. The clathration process stops when no more crystalline water ice is available to trap the volatile species. Note that, in the pressure conditions of the solar nebula, CO₂ is the only species that crystallizes at a higher temperature than its associated clathrate. We then assume that solid CO₂ is the only existing condensed form of CO₂ in this environment. In addition, we have considered only the formation of pure ice of CH₃OH in our calculations since, to our best knowledge, no experimental data concerning the equilibrium curve of its associated clathrate have been reported in the literature. The intersection of a thermodynamic path at a given distance from the star with the equilibrium curves of the different ices allows determination of the amount of volatiles that are condensed or trapped in clathrates at this location in the disk. Indeed, the volatile, i, to water mass ratio in these planetesimals is determined by the relation (Mousis & Gautier 2004; Mousis & Alibert 2006)

$$\begin{equation} {m_i = \frac{X_i}{X_{{\rm H_2O}}} \frac{\Sigma (r; T_i, P_i)}{\Sigma (r; T_{{\rm H_2O}}, P_{{\rm H_2O}})}}, \end{equation} \tag{ 2 }$$

where X_i and $X_{{\rm H_2O}}$ are the mass mixing ratios of the volatile i and H₂O with respect to H₂ in the nebula, respectively. Σ(R; T_i, P_i) and $\Sigma (R; T_{{\rm H_2O}}, P_{{\rm H_2O}})$ are the surface density of the nebula at a distance r from the Sun at the epoch of hydration or clathration of the species i, and at the epoch of condensation of water, respectively. From m_i, it is possible to determine the mass fraction M_i of species i with respect to all the other volatile species taking part in the formation of an icy planetesimal via the following relation:

$$\begin{equation} {M_i = \frac{m_i}{\sum _{j=1,n} m_j}}, \end{equation} \tag{ 3 }$$

with ∑_{i = 1, n}M_i = 1.

Assuming that, once condensed, the ices add to the composition of planetesimals accreted by the growing planet along its migration pathway, this allows us to reproduce the volatile abundances by adjusting the mass of planetesimals that vaporized when entering the envelope. Note that, because the migration path followed by the forming WASP-12b is unknown, the 5–20 AU distance range of the ensemble of thermodynamic paths has been arbitrarily chosen to determine the composition of the accreted ices. The adoption of any other distance range for the planet's path beyond the snow line would not affect the composition of the ices (and thus WASP-12b's global volatile abundances) because it remains almost identical irrespective of (1) their formation distance and (2) the input parameters of the disk, provided that the initial gas-phase composition is homogeneous (Marboeuf et al. 2008).

Atmospheric chemistry in hydrogen-dominated atmospheres is critically influenced by the ratio of carbon to oxygen. The major species resulting from carbon–oxygen chemistry in chemical equilibrium are water vapor (H₂O), carbon monoxide (CO), and methane (CH₄). For solar abundances (C/O = 0.54), CH₄ is the dominant carrier of carbon at low temperatures (T), and CO is the dominant carbon carrier at high T, as shown in Figure 2. The temperature of transition ($T_{\rm CO - CH_4}$) depends on the pressure (P); at a representative P ∼ 1 bar, $T_{\rm CO - CH_4} \sim$ 1200–1400 K. On the other hand, H₂O is a major carrier of oxygen at all temperatures. At low T, almost the entire O is contained in H₂O, whereas at high T, a significant amount of O is also contained in the abundant CO. As CO overtakes CH₄ at $T \sim T_{\rm CO - CH_4}$, O is present in nearly equal proportions between H₂O and CO. This distribution of H₂O, CO, and CH₄ for solar abundances is substantially perturbed when the C/O ratio is increased.

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At high T, hydrogen-dominated carbon-rich (with C/O ⩾ 1) atmospheres bear manifestly different concentrations of the major species from those in solar abundance atmospheres, as shown in Figure 2, and previously suggested by Kuchner & Seager (2005). One of the most apparent changes is a general enhancement in CH₄ abundance and depletion in H₂O abundance over a wide region in P–T space. Most of the O is occupied by CO over a wide temperature range. The degree of H₂O depletion over solar abundance concentrations increases with T. Considering the C/O = 1 case in Figure 2, for instance, at a nominal P ∼ 1 bar, H₂O can be depleted by a factor of 10–100 for T between 1400 and 2000 K. At lower pressures, such as P ∼ 0.1 bar, which are typically probed by the observations (Madhusudhan et al. 2011), the depletion is even greater. In contrast, over the same temperature and pressure range, the CH₄ abundance exceeds that in solar abundance by factors as high as 10³.

The depletion in H₂O and enhancement in CH₄, at high T, increase further as the C/O ratio increases beyond 1, as shown for C/O = 3 in Figure 2. Thus, whereas solar abundances typically suggest H₂O mixing ratios of ∼5 × 10⁻⁴ at all temperatures, high C/O ratios can cause H₂O mixing ratios as low as ∼10⁻⁷–10⁻⁶ in the observable atmospheres (P ≲ 0.1–1 bar) at high T (≳  1400 K). Similarly, while CH₄ is generally presumed to be scarce at high T, high C/O ratios can yield observable CH₄ abundances up to ∼10⁻⁵–10⁻⁴.

The strong influence of the C/O ratio on atmospheric composition at high T indicates that highly irradiated transiting hot Jupiters and hot Neptunes are ideal probes of C/O ratios and form good candidates to search for CRPs. At lower temperatures (T ≲ 1200 K at P ∼ 1 bar), the distinction between the C-rich and O-rich regimes is less apparent, as shown in Figure 2. In this regime, irrespective of the C/O ratio, most of the carbon is carried by CH₄, CO is scarce, and H₂O is the primary carrier of oxygen. Thus, for C/O = 1, CH₄ and H₂O are equally abundant, with CH₄ being only a factor of two more abundant than in a solar abundance model (C/O = 0.54). Such small deviations of less than a factor of ∼10 in molecular abundances are well below the uncertainties in current molecular abundance determinations for exoplanetary atmospheres (e.g., Swain et al. 2008; Madhusudhan & Seager 2009, 2011; Madhusudhan et al. 2011). Consequently, based on equilibrium chemistry alone, reliable determination of C/O ratios of cooler giant planets such as GJ 436b, with T ≲ 1200 K at P ∼ 1 bar, will need high-resolution spectra and high-precision retrieval of CH₄ and H₂O abundances along with those of other species.

The concentrations of CO and CH₄ are also influenced by non-equilibrium chemistry in the form of vertical mixing, as discussed in Section 2.1, especially at low T. In particular, enhanced CO abundances over that possible in chemical equilibrium can be transported from the lower hotter regions of the atmosphere to upper observable regions of the atmosphere. As shown in Madhusudhan & Seager (2010) for GJ 436b, mixing ratios of order 10⁻³ are attainable for high metallicities, corresponding to ∼10–30 × solar. However, in the present context the question is whether vertical mixing is sensitive to the overall atmospheric C/O value in the low-T regime. Figure 3 shows the range of CO abundances possible in the observable atmosphere of GJ 436b due to non-equilibrium chemistry with different C/O ratios. We find that the differences between the CO abundances resulting from the variation in C/O ratios merely reflect the higher C/H abundances in the various cases. Note that we fix the O/H and vary the C/H to obtain different C/O ratios. As such, the CO abundance from our C/O = 3 case, for example, can also be obtained by a C/O = 0.5 model with ∼6 times higher bulk metallicity. We therefore find that even with the effect of vertical mixing, high and low C/O atmospheres would be challenging to distinguish observationally. In principle, other mechanisms of non-equilibrium chemistry such as photochemistry (e.g., Line et al. 2010; Moses et al. 2011) might give rise to effects that could distinguish between the different C/O ratios. For example, if methane were destroyed and CO enhanced by non-equilibrium processes (Madhusudhan & Seager 2011), the observable concentrations of H₂O and CO, along with other photochemical products, might be used to estimate the C/O ratio of the atmosphere.

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Our results indicate that carbon-rich atmospheres are less likely to host thermal inversions due to TiO and VO. Figure 4 shows molecular mixing ratios of TiO and VO for different C/O ratios over a range of temperatures characteristic of highly irradiated hot Jupiters. WASP-12b, which is one of the most extremely irradiated hot Jupiters, is known to have a lower atmosphere consistent with an isotherm of ∼3000 K (Madhusudhan et al. 2011). We therefore consider three isothermal temperature profiles at 2650 K, 3000 K, and 3250 K spanning a representative range for very hot Jupiters. For each temperature, we determine TiO and VO abundances in chemical equilibrium for C/O = 0.54, i.e., solar (solid curves) and C/O = 1 (dotted curves), as shown in Figure 4. We find that a C/O = 1 leads to TiO and VO abundances of ∼100 times lower than the corresponding abundances for solar C/O, at the P ∼ 1 bar level. The depletion is even more enhanced at the 0.1 bar level, where the planetary "photosphere" typically lies.

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The depletion of TiO and VO due to an increased C/O ratio offers a natural explanation for the lack of a thermal inversion in the atmosphere of WASP-12b which is known to have a C/O ⩾1 (Madhusudhan et al. 2011). The ∼100 times depletion in TiO and VO for a C/O = 1 as discussed above may well be a lower limit. Considering higher C/O ratios, as are possible in WASP-12b, would lead to even greater TiO/VO depletion. Considering non-equilibrium gravitational settling, as suggested to be important by Spiegel et al. (2009), could contribute additional depletion. Spiegel et al. (2009) further suggest that a depletion in TiO below solar abundances by a factor of ∼10 is enough to rule out a thermal inversion in the atmosphere. Thus, considering all the effects discussed above, we conclude that the degree of TiO/VO depletion caused by a C/O ⩾1 is adequate to rule out a thermal inversion due to TiO and VO even in the most highly irradiated hot Jupiters, such as WASP-12b. It remains to be seen if other compounds might exist in highly irradiated carbon-rich atmospheres that could provide visible or ultraviolet absorption comparable to TiO/VO, such as ZrO and/or YO found in carbon-rich stars (e.g., Jørgensen 1994) or higher order hydrocarbons. Finally, the abundances of TiO/VO and other inversion causing compounds may also be correlated with the stellar activity (Knutson et al. 2010).

The chemical compositions of CRGs lead to observable spectroscopic signatures over a wide range of atmospheric temperature profiles. The key differences between spectra of CRGs and those of giant planets with solar C/O ratio result from their differing concentrations of the major molecular absorbers, H₂O, CO, and CH₄, depending on the temperature regime. As discussed in Section 3.1, at high temperatures (T ≳ 1400 K, for P < 1 bar), giant planet atmospheres with solar C/O ratio (i.e., oxygen rich) are expected to be dominated by CO and H₂O as the major carbon- and oxygen-bearing species, while CH₄ is scarce. On the other hand, atmospheres of CRPs (C/O ⩾1) at the same high temperatures are characterized by a substantial enhancement of CH₄ and depletion of H₂O, compared to those with solar C/O; the amount of CO remains comparable. At low temperatures, however, the distinction between the different C/O regimes is less pronounced. Atmospheres at lower T_eq for both C/O regimes have comparable amounts of both CH₄ and H₂O, and much less CO. These compositional characteristics of high and low C/O atmospheres in the different temperature regimes are manifested in the corresponding spectra.

As discussed in Section 2.3, we construct dayside thermal spectra of giant exoplanetary atmospheres over a representative range of irradiation levels (or temperatures) and C/O ratios. Model spectra for the different cases are shown in Figure 5. As shown in the top-left panel, for an extremely irradiated hot Jupiter (with T_eq ∼ 2300 K), the spectrum of a carbon-rich atmosphere is distinctly different from one with a solar C/O, across a broad spectral range in the near- to mid-infrared. In particular, the CRG spectrum contains high absorption due to the strong CH₄ bands around 3.5 μm and 8 μm, along with milder features in the K band (around 2.2 μm), while predicting high flux (or low absorption) at wavelengths associated with strong H₂O absorption (e.g., at wavelengths 1.8 μm, 2.8 μm, 5–7 μm, and >15 μm). In contrast, spectra with solar C/O ratios are saturated with absorption features due to H₂O, and display minimal CH₄ absorption. In both high and low C/O cases, however, the spectra show comparable absorption in the CO bands, around 4.5 μm. The top-right panel of Figure 5 shows thermal spectra of a typical hot Jupiter with T_eq ∼ 1500 K and with different C/O ratios. The differences between the solar and high C/O cases are still evident across the CH₄ and H₂O bands. The bottom panels of Figure 5 show two cases of giant planets at the cooler end of known transiting planets: a hot Neptune with T_eq ∼ 1000 K (bottom-left panel) and a warm Neptune with T_eq ∼ 750 K (bottom-right panel). In both cases, the differences between spectra of models with solar C/O and those with C/O = 1 are much less discernible compared to the hotter cases discussed above. However, spectra of CRGs with T_eq ∼ 1000 K are still distinguishable from solar-composition models for CRG atmospheres with C/O ≳ 3. On the other hand, CRGs with T_eq ∼ 750 K are nearly indistinguishable from solar models even for high C/O ratios.

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Currently known transiting giant exoplanets provide a rich sample for future observational campaigns in search of CRGs. The higher temperature hot Jupiters and hot Neptunes (with T_eq ≳ 1000 K) provide the most favorable conditions to detect CRGs. Such atmospheres show clear spectroscopic differences between solar and carbon-rich compositions, with a signal clearly detectable with existing or forthcoming instruments and analysis techniques (e.g., Madhusudhan et al. 2011). For instance, in a CRG atmosphere observed with the equivalent of four channels of Spitzer IRAC photometry, the thermal spectrum would display a high flux in the 5.8 μm channel (due to low water absorption) and low 3.6 μm and 8 μm fluxes (due to strong methane absorption). On the other hand, giant planets with low irradiation levels, T_eq ≲ 800 K, are not readily distinguishable between the low and high C/O cases, unless high-resolution spectroscopy is available allowing high-precision retrieval of H₂O and CH₄ abundances.