Context. Protoplanetary disks the birthplaces of planets commonly feature bright rings and dark gaps in both continuum and line emission maps. Accreting planets interact with the disk not only through gravity but also by changing the local irradiation and elemental abundances which are essential ingredients for disk chemistry. BR  Aims We propose that giant planet accretion can leave chemical footprints in the gas local to the planet which potentially leads to the spatial coincidence of molecular emissions with the planet in the ALMA observations. BR  Methods Through 2D multifluid hydrodynamical simulations in Athena with builtin sublimation we simulated the process of an accreting planet locally heating up its vicinity opening a gas gap in the disk and creating the conditions for Cphotochemistry. BR  Results An accreting planet located outside the methane snowline can render the surrounding gas hot enough to sublimate the Crich organics off pebbles before they are accreted by the planet. This locally elevates the disk gasphase CO ratio providing a potential explanation for the CSUB2SUBH lineemission rings observed with ALMA. In particular our findings provide an explanation for the MWC 480 disk where previous work identified a statistically significant spatial coincidence of lineemission rings inside a continuum gap. BR  Conclusions Our findings present a novel view of linking the gas accretion of giant planets and their natal disks through the chemistry signals. This model demonstrates that giant planets can actively shape their forming chemical environment moving beyond the traditional understanding of the direct mapping of primordial disk chemistry onto planets. P Movie associated to Fig. 4 is available at A hrefhttpswww.aanda.org10.105100046361202346637olmhttpswww.aanda.orgA
['astrochemistry', 'protoplanetary disks', 'planet-disk interactions', 'accretion', 'accretion disks', 'hydrodynamics', 'stars: variables: T Tauri', 'Herbig Ae/Be', 'Astrophysics - Earth and Planetary Astrophysics']
Chemical footprints of giant planet formation. Role of planet accretion in shaping the CO ratio of protoplanetary disks
{'Role of Planet Accretion in Shaping the C/O Ratio of Protoplanetary Disks': 'Haochang Jiang ( 蒋 昊昌 ) 1,2 , Yu Wang ( 王 雨 ) 2 , Chris W. Ormel 2 , Sebastiaan Krijt 3 , Ruobing Dong ( 董 若 冰 ) 4  \n- 1 European Southern Observatory, Karl-Schwarzschild-Str 2, 85748 Garching, Germany\n- 2 Department of Astronomy, Tsinghua University, 30 Shuangqing Rd, Haidian DS 100084, Beijing, China\n- 3 School of Physics and Astronomy, University of Exeter, Stocker Road, Exeter EX4 4QL, UK\n- 4 Department of Physics and Astronomy, University of Victoria, Victoria, BC V8P 5C2, Canada e-mail: hjiang@eso.org; wang-y21@mails.tsinghua.edu.cn  \nReceived February 1, 1997; accepted March 23, 1997', 'ABSTRACT': 'Context. Protoplanetary disks, the birthplaces of planets, commonly feature bright rings and dark gaps in both continuum and line emission maps. Accreting planets are interacting with the disk, not only through gravity, but also by changing the local irradiation and elemental abundances, which are essential ingredients for disk chemistry.  \nAims. We propose that giant planet accretion can leave chemical footprints in the gas local to the planet, which potentially leads to the spatial coincidence of molecular emissions with the planet in ALMA observation. Methods. Through 2D multi-fluid hydrodynamical simulations in Athena ++ with built-in sublimation, we simulate the process of an accreting planet locally heating up its vicinity, opening a gas gap in the disk, and creating the conditions for C-photochemistry. Results. An accreting planet located outside the methane snowline can render the surrounding gas hot enough to sublimate the C-rich organics o ff pebbles before they are accreted by the planet. This locally elevates the disk gas-phase C / O ratio, providing a potential explanation for the C2H line-emission rings observed with ALMA. In particular, our findings provide an explanation for the MWC 480 disk, where previous work has identified a statistically significant spatial coincidence of line-emission rings inside a continuum gap. Our findings present a novel view of linking the gas accretion of giant planets and their natal disks through the chemistry signals. This model demonstrates that giant planets can actively shape their forming chemical environment, moving beyond the  \nConclusions. traditional understanding of the direct mapping of primordial disk chemistry onto planets.  \nKey words. protoplanetary disk - planet formation - astrochemistry', '1. Introduction': "Protoplanetary disks are the cradle of planets. Imaging observations with high resolution and high contrast using instruments such as ALMA and Gemini / GPI have found that most large disks ( r > 50 au), if not all, exhibit substructures such as bright rings and dark gaps (e.g., Andrews et al. 2018; Long et al. 2018; van der Marel et al. 2019). It is widely speculated that planet-disk interactions are responsible for many of these substructures (e.g, Zhu et al. 2011; Dong et al. 2015; Rosotti et al. 2016; Zhang et al. 2018; Paardekooper et al. 2022). Nevertheless, detecting planets in these young systems is challenging (Xie et al. 2020; Huélamo et al. 2022; Follette et al. 2022). Besides PDS 70 (Keppler et al. 2018; Ha ff ert et al. 2019) and AB Aur (Currie et al. 2022; Zhou et al. 2022), most claims of planet detections are indirect (e.g., Pinte et al. 2019, 2020; Speedie et al. 2022; Izquierdo et al. 2022) and might be debatable (Speedie & Dong 2022; Fedele et al. 2023).  \nAmong the indirect planet detection method, the role of astrochemistry is gaining increasing prominence with the development of facilities such as ALMA and JWST (e.g., Öberg & Bergin 2021). The classical core accretion planet formation dictates that a Neptune-mass planet will quickly grow into a Jovianmass gas giant during a runaway growth phase, in which gas freely falls into its gravitational potential. During the runaway phase the planet's mass and accretion rate ˙ M p exponentially  \ngrow, making the protoplanet an irradiating source in the disk (e.g., Rafikov 2006; Mordasini et al. 2012).  \nLuminous young planets locally heat up their vicinity, resulting in e ffi cient thermal sublimation of molecular species that are originally frozen out. Therefore, localized submillimeter molecular emission might appear around the planet as proposed in Cleeves et al. (2015). This scenario is particularly compelling for the recently identified local emission enhancements of SO in HD100546 (Booth et al. 2023), and 13 CO in AS209 (Bae et al. 2022). Yet, the calculation of Cleeves et al. (2015) used a static gas density profile, where the di ff erential rotation of material around the planet is ignored. For a gap-opening planet at distant r 0, as the planet opens a wider gap with width ∆ w as it is accreting, the synodic timescale around the planet gap r 0 / ∆ w Ω K , where Ω K is the Keplerian frequency, gets shorter. If this timescale is shorter than the typical chemical timescale, materials can be sheared around the co-orbital region of the accreting planet, and the asymmetric planet-heated chemical emission might be elongated and become arc- or even ring-like.  \nThanks to the high spatial resolution and high sensitivity of ALMA, ring-like substructures are found to be ubiquitous in line-emission maps as well (e.g., Öberg et al. 2021). Among those species with line-emission detection, the radial C2H, one of the most well-studied hydrocarbons, is believed an indicate an active photo-chemistry, driven by high UV-fluxes and gas with a C / O ratio exceeding unity (Bergin et al. 2016; Kama et al. 2016;  \n∂ρ p ∂ t + ∇ · GLYPH<16> ρ p v p + F p , dif GLYPH<17> = 0 ,  \nGLYPH<16>  \nv  \n+  \nv  \nt  \np  \n∂  \nGLYPH<17>  \n∂ρ  \np  \n,  \np  \ndif  \n+  \n∇ ·  \nρ  \np  \np  \n,  \nGLYPH<16>  \nρ  \nsrc  \nArticle number, page 2 of 18  \np  \nv  \n+  \nf  \np  \nv  \nρ  \np  \np  \n+  \nΠ  \nv  \ng  \n-  \nv  \nt  \ns  \ndif  \np  \nGLYPH<17>  \n,  \n=  \nMiotello et al. 2019; Bergner et al. 2019; Alarcón et al. 2021; van der Marel et al. 2021; Bosman et al. 2021a). C2H can serve as a precursor for larger molecule formation through chemical reaction network (e.g., Bast et al. 2013; Öberg & Bergin 2021). Alternatively, C2H may also be a product of top-down photodissociation hydrocarbon chemistry (Guzmán et al. 2015).  \nDue to the specific formation conditions required for C2H high C / O ratio and high UV penetration- and its bridging role in the chemical reaction network, we aim to explore whether a young giant planet can be responsible for achieving these conditions, and discuss how the chemical footprint of the accretion of the planet will manifest in protoplanetary disks.  \nIn this study, we conduct 2D multi-fluid hydrodynamical simulations that include built-in sublimation to simulate an accreting planet in Athena ++ (Huang & Bai 2022; Wang et al. 2023). The planet is modeled as a sink particle in the simulation with self-consistent accretion prescription. As the planet accretes, the potential energy of the accreted gas is released as the planet's luminosity, locally heating the gas surrounding the planet and releasing C-rich volatile, whose abundance we follow in Athena ++ . We then post-process our simulation outputs with an empirical formula derived from chemical network simulations on C2H formation. In addition, we compare our simulation results with line-emission observations from ALMA.  \nThe plan of the paper is as follows. The theoretical background and numerical setup are described in Section 2. We present and explain the numerical results in Section 3. The results are further discussed in Section 4. We summarize the main results and conclusions in Section 5.", '2. Model': 'We simulate the planet and disk interaction with multi-fluid 2D hydrodynamics simulations in polar coordinates (radius r and azimuth ϕ ). The disk consists of three mass components non-condensable gas, ice-rich pebbles, and sublimated vapor. To model them, we introduce gas fluid (\'g\'), particle fluid (\'p\') and tracer fluid (\'tr\') respectively. The \'gas fluid\' is the mixture of non-condensable gaseous species (mostly H and He) and sublimated vapor. There is only one gas fluid. In addition, the density of each vapor species is followed by a \'tracer fluid\', which inherits the hydrodynamical properties (i.e., velocity field) of the gas fluid. Pebbles are modelled as compound of multiple \'particle fluids\', each of which represents a solid component (i.e., ice and refractory). With the recently developed multi-dust fluid module of Athena++ (Huang & Bai 2022), the gas and ice can co-evolve in a self-consistent manner. The governing equations read (see also Sec. 2.2 of Wang et al. (2023)),  \n∂ρ g ∂ t + ∇ · ( ρ g v g) = 0 , (1)  \n∂ ( ρ g v g) ∂ t + ∇ · ( ρ g v g v g + P g I + Π ν ) = ρ g f g , src , (2)  \n∂ρ tr ∂ t + ∇ · GLYPH<16> ρ tr v g + F tr , dif GLYPH<17> = 0 . (5)  \nIn the above, ρ is the density, v is the velocity, P g is the gas pressure, I is the identity tensor. The gas viscosity is described by the viscous stress tensor  \nΠ ν ≡ ρ g ν g " ∇ v g + GLYPH<16> ∇ v g GLYPH<17> T -2 3 GLYPH<16> ∇ · v g GLYPH<17> I # (6)  \nwhere the kinematic viscosity ν g is specifically expressed as a function of the dimensionless constant α (Shakura & Sunyaev 1973)  \nν g = α csH g . (7)  \nwhere H g is the gas scale height and cs is the local sound speed. With this, the particle concentration di ff usion flux is given by  \nF p , dif = -ρ g D p ∇ ρ p ρ g ! ≡ ρ p v p , dif , (8)  \nwhich meanwhile defines the e ff ective di ff usive velocity v p , dif. And the di ff usivity for dust can be prescribed as (Youdin & Lithwick 2007)  \nD p ≡ ν g 1 + St 2 (9)  \nwith the Stokes number St ≡ t s Ω K . Π dif is the momentum di ff usion flux tensor which, combined with v p , dif, ensures the Galilean invariance (Huang & Bai 2022) in Eq. (4). Similar to Eq. (8), the tracer concentration di ff usion flux is defined as,  \nF tr , dif = -ρ g D tr ∇ ρ p ρ g ! (10)  \nwith the tracer di ff usivity D tr = ν g, where we implicitly assume that the Schmidt number is 1. Without loss of generality, we assume the gas viscosity coe ffi cient α to be 5 × 10 -4 , and initialize the pebbles with constant Stokes number St = 0 . 01 throughout the domain. By assuming spherical pebbles, the particles size  \nsp = 2 π St Σ g ρ · (11)  \nis ∼ 0.5 mm at the planet orbit, with internal density ρ · = 1 . 67g cm -3 (Birnstiel et al. 2018). Those values match the latest constraint approaches based on ALMA observational results (see the recent review by Rosotti 2023, and references therein). In our simulations setup, the particle size sp is kept constant throughout the simulation. The final term in Eq. (4) represents the aerodynamic drag experienced by particles, which follow the Epstein regime throughout our simulation, with the backreaction on gas omitted in our formulation. External source terms are denoted by f src, including stellar and planetary gravity (see Eq. (18)).', '2.1. Disk model': '(3) We assume that the initial gas and pebble surface densities follow power-law distributions  \n(4) Σ g = Σ g , 0 × r r 0 ! -γ Σ p = Σ p , 0 × r r 0 ! -γ (12)  \nwhere r is the distance from the star and Σ g , 0 is the gas surface density at the reference location r 0. The gas disk is parameterized based on the observational constraints of MWC 480, a typical Herbig Ae / Be star with a clear annular gap in its continuum at ∼ 73 au (Long et al. 2018; Liu et al. 2019). From thermo-chemical modeling based on line emission of CO isotopes, Zhang et al. (2021) suggest a gas surface density profile of Σ g , 0 = 11g cm -2 , γ = 1 at r 0 = 73 au. The densities of the ice-rich pebble at r 0 is Σ p , 0 = f m Σ g , 0, where the mass fraction f m depends on the volatiles species of interest in each simulation (see Table 1 for the values of each species). The stellar masses have been dynamically determined to be ∼ 2 . 1 M ⊙ as described in (Teague et al. 2021).  \nWeassume that the disk is isothermal in the vertical direction and employ a locally-isothermal equation of state. The ambient temperature is determined by the disk background, which is influenced by the irradiation of the host star, and the heating from the accreting planet. Specifically, the disk background temperature is described as  \nT d = T 0 × r r 0 ! -q (13)  \nwhere T 0 = 20 K , and q = 0 . 45 according to the power-law fit of the dust temperature at the midplane (Sierra et al. 2021), based on the 2D (R-Z) dust density and temperature profile in the themo-chemical models in Zhang et al. (2021). The aspect ratio of the gas disk, therefore, follows  \nh g ≡ H g r ≡ cs Ω Kr = h 0 × r r 0 ! 0 . 5 -q / 2 (14)  \nwith h 0 = 0 . 05 based on the temperature profile. We note this value is inconsistent with the best-fit aspect ratio based on the SED fitting. Zhang et al. (2021) fit the vertical temperature profile at each radius directly, by matching the emission surface of CO isotope lines. Since our model focuses on the ice sublimation at the midplane, we opt for the midplane temperature in our simulations.  \nAs the planet is accreting, the lost gravitational energy of the accreting materials can be translated into accretion luminosity  \nL acc = GM p ˙ M R p . (15)  \nwhere M p is the planet mass, and R p is the radius of the planet accretion shock front, and ˙ M is the mass accretion rate of the planet. We fix R p = 2 R Jupiter in all of our simulations, which is the typical shock front radius of an accreting planet in the disk (1.54 R Jupiter, see e.g., Marleau et al. 2022, and references therein). A detailed explanation for the prescription of planet mass M p and accretion rate ˙ M is given in Section 2.3.  \nWe utilize the accretion luminosity to prescribe the temperature surrounding the planet. We consider the radiation near the planet to be optically thin at infrared wavelength, which is typically the case in the planet-opened gaps at the outer regions of the protoplanetary disk. The temperature is thus prescribed as (e.g., Rafikov 2006)  \nT = GLYPH<18> L acc 16 πσ d 2 + T 4 d GLYPH<19> 0 . 25 (16)  \nwhere d is the distance to the planet.  \nTable 1. Abundances n (X) / n (H) for di ff erent molecules X of interest. Similar values are used in recent modeling (e.g., Bosman et al. 2021a; Schneider & Bitsch 2021), and consistent with cometary ices in forming protoplanetary disk midplanes (Drozdovskaya et al. 2016).', '2.2. Volatile sublimation': 'At each grid point, volatile ice can be sublimated into vapor via a recently-developed phase change module implemented in Athena++ (Wang et al. 2023). This module solves mass transfer due to ice sublimation after each hydrodynamic step, thus enabling the self-consistent simulation of volatile transport in both ice and vapor states. The volatile follows the saturated (or equilibrium) vapor pressure, which is given by the ClausiusClapeyron equation  \nP eq = P eq , 0 e -Ta / T (17)  \nwhere Ta and P eq , 0 are constants depending on the molecular species. We take Ta = 1110 K and P eq , 0 = 3 . 67 × 10 9 g cm -1 s -2 for methane-the major C-carrier icy volatile (Fray & Schmitt 2009, where we neglect higher order correction terms). When the partial pressure of the volatile vapor is lower than the saturated vapor pressure, ice is assumed to be instantly converted into vapor. In reality, the sublimation timescale of the volatile ∝ d s / v th, where d s is the depth of the volatile-rich surface layer of the pebble and v th is the thermal velocity of vapor particles (e.g., Schoonenberg & Ormel 2017). Thus the instant sublimation of ice-rich pebbles is validated for the volatiles of our interests, which are assumed covered on the surface of the pebble. As the released vapor will mainly stay inside the gap but outside the co-orbital region, where there are few pebbles for vapor to condense on, we, for simplicity, ignore condensation in the simulations (see Sect. 4.1 for further discussion and Fig. 10).  \nIn each simulation, pebbles are assumed to be composed of refractories, which dominate by mass, and a single ice species, which can sublimate from ice to vapor phase. Therefore, we assume that the size of pebbles is not a ff ected by sublimation. The abundances of all species included in our simulation are listed in Table 1. The values are marginally the same as those taken in Bosman et al. (2021a), and consistent with cometary ices in forming protoplanetary disk midplanes (Drozdovskaya et al. 2016). CH4 are ice-rich pebbles that are initially in the ice phase, in other words, the planet locates beyond their snowlines. Other materials, including CO, are assumed to be non-condensable gases throughout the simulation, which means that the planet is located inside their snowlines. For instance, as suggested in Zhang et al. (2021), the midplane CO snowline for MWC 480 is around 100 au.', '2.3. Growth of the planet': 'In our 2D simulation, we modeled the planet as a sink particle with a fixed circular orbit at a radius of r 0. The planet mass is initialized as 10 M ⊕ , a typical critical core mass for a planet starting the runaway gas accretion (e.g., Perri & Cameron 1974; Rafikov 2006), but also see recent works considering atmosphere pollution suggesting even lower values (Hori & Ikoma  \nFig. 1. Sketch of the 2D mesh grid used in our simulations. The circle represents the location of the planet (located at r 0 = 73 au) and the star symbol is the central star. The radial domain ranges from 30 au to 150 au. The static mesh refinement level is 2, which means the smallest mesh is four times smaller than the unrefined mesh at the same location.  \n<!-- image -->  \n2011; Ormel et al. 2021). The static mesh refinement was employed throughout the simulation to better capture the gas flow and ice sublimation in the vicinity of the planet, as depicted in Fig. 1. The simulation domain extended radially from 30 au to 150 au, and the smallest mesh size was d x min = 0 . 1 R Hill, where R Hill ≡ ( M p / 3 M ⋆ ) 1 / 3 r 0 is the Hill radius of the planet. We used the value of R Hill for a planet mass of M p = M 0 ≡ 2 . 3 M J, as suggested by previous studies reproducing the continuum profile of MWC480 (Liu et al. 2019). We tried smaller d x min for a convergent test, while the simulation results have no changes. A softening parameter ϵ = 0 . 1 R Hill is chosen for the gravitational potential  \nΦ = -GM p √ d 2 + ϵ 2 . (18)  \nThe choice of ϵ = 0 . 1 R Hill ensures that the accretion rate on the planet is high, which feature is essential to capture with our model as it heats the vicinity of the planet, allowing pebbles to sublimate. In the literature, however, a larger softening parameter is favored for 2D simulations in order to correctly capture the wave excitation and gap opening by the planet (e.g., Dong et al. 2011; Müller et al. 2012). Our choice for ϵ in 2D simulations could therefore expedite the gap-opening process somewhat. However, conducting a run with ϵ = 0 . 6 Hg , we found little di ff erence in the gap profile and therefore do not expect that the choice of ϵ would much a ff ect our conclusions.  \nStarting with the initial planet mass of 10 M ⊕ , we run the simulation with the fixed-initial-mass planet for t inin = 300 orbits, which is about 10 5 yr at the distance of the planet. The time is su ffi ciently long compared with a total disk lifetime, and enough for the disk to adapt itself to the planet embedded in the disk. We tried di ff erent t inin, with which our major conclusion is una ff ected. After t inin, we assign a mass fraction of the ice-rich pebbles (Table 1), then switch on mass accretion in the simulations. We remove all pebbles entering 0 . 9 R Hill from  \nthe simulation. The gas accretion prescription follows Crida & Bitsch (2017) and Bergez-Casalou et al. (2020), and the accretion rate is based on the 3D isothermal shearing box simulations results from Machida et al. (2010), which is determined by the local gas density within 0 . 9 R Hill of the planet  \n˙ M p = f A D Σ g E 0 . 9 R Hill H 2 g Ω K × min         0 . 14 , 0 . 83 × R Hill H g ! 9 2         . (19)  \nwhere H g is the gas scale height. The prefactor f A equals unity in the default setup, while we also try di ff erent values in Section 3.2. At each timestep, the accreted mass of gas is removed from the domain and added to the planet accordingly. Following Bergez-Casalou et al. (2020), the distribution of the mass removing rate ˙ Σ g , acc( d , t ) reads  \n˙ M p = " d < 0 . 9 R Hill ˙ Σ g , acc( d , t ) d A = " d < 0 . 9 R Hill f red( d ) ˙ Σ g( t ) d A (20)  \nwhere the smooth distribution function f red reads  \nf red =          1 d < 0 . 45 R Hill cos 4 π d R Hill -0 . 45 ! 0 . 45 R Hill < d < 0 . 9 R Hill (21)  \nfollowing Robert et al. (2018). The fraction of accreted mass increases as it gets closer to the planet, but turns flat within the innermost 0 . 45 R Hill. The gas removal rate within 0 . 45 R Hill, ˙ Σ g( t ), is obtained by equating Eqs. (19) and (20). In other words, the mass accretion rate of the planet is calculated through Eq. (19), and the accreted gas is removed from the surrounding protoplanetary disk according to Eq. (20).', '2.4. Sketch of the chemical evolution': "We propose a chemical model to interpret a potential spatial correlation between the chemical ring and dust continuum profile caused by a planet exterior to the methane snowline. This model hinges on the local increase in the C / O ratio, which is attributed to localized volatile sublimation caused by the luminosity of the accreting planet. A sketch of the model can be found in Fig. 2.  \nHigh concentrations of organics like C2H and HCN indicate an active photo-chemistry, driven by high UV-fluxes and gas with a C / O ratio exceeding unity (Bergin et al. 2016; Bosman et al. 2021a). In our model, a young giant planet is responsible for achieving these conditions. First, pebbles in the planet's orbital region are accreted by the growing protoplanet. Since the accreting planet is hot, carbon-rich volatile ices (e.g., methane) with sublimation temperatures around 25 K are readily liberated from the ice, elevating the C / O ratio of the gas in the vicinity of the planet's orbit above unity. On the other hand, the sublimation of CO2 (H2O) - the major oxygen carrier - requires a temperature in excess of ∼ 55 K (120 K). Such a high temperature at 73 au can only be reached in the immediate vicinity of the planet ( < R Hill), where the water and carbon-dioxide vapor can no longer escape the planet's gravity (e.g., Johansen et al. 2021; Barnett & Ciesla 2022). In other words, the sublimation temperature di ff erence between C-rich ices (e.g., CH4, C2H2)and O-rich (e.g., CO2, H2O) renders the gas-phase C / O > 1. Then, as the gap region becomes optically thinner due to the removal of material, stellar UV photons penetrate the gap more deeply and trigger photo-chemistry reactions. These processes may produce  \nFig. 2. Schematic of how a gap-opening planet triggers the C2H ring associated with the continuum gap. (a) Before runaway growth, a small planet is embedded inside the gaseous disk and carbon-rich ice is frozen out on pebbles. (b) During runaway gas accretion the planet grows rapidly, consuming the pebbles in its vicinity and opening a gas gap in the disk. Since the accreting planet is hot, carbon-rich volatile ice (e.g., methane) with lower sublimation temperature is readily liberated from the ice. On the other hand, the sublimation of H2O or CO2 - the major oxygen carriers - can only be reached in the immediate vicinity of the planet ( < R Hill), by which point the material is lost to the planet. The discrepancy in sublimation temperature renders the gas C / O > 1. In addition, stellar UV photons penetrate the gas gap. The combination of C / O > 1 and high UV-fluxes trigger photo-chemistry reactions to form the C2H ring; (c) After the gap opens, the supply of sublimated pebbles terminates due to the gap opening. The C / O ratio is smoothed out across the disk through radial mixing by di ff usion.  \n<!-- image -->  \nring-like structures in C2H (Bergin et al. 2016), one of the typical carbon-bearing radicals. And the more complex gas-phase formation of organic molecules compounds (Bast et al. 2013; Calahan et al. 2023).", '2.5. C2H chemistry': "We focus on C2H rings in this work. As a radical, C2H can act as a precursor for the formation of larger molecules, through photochemical reactions (e.g., Bast et al. 2013; Öberg & Bergin 2021). On the other hand, it is also proposed that C2H can be a product of top-down photo-dissociation hydrocarbon chemistry (Guzmán et al. 2015). From the observational side, Bergner et al. (2019) has suggested that C2H and HCN seem to share a common driver. More research is needed to fully understand the connection between C2H and the formation of complex organics in these environments, where a close link between a high C / Oratio and molecular formation appears.  \nTo estimate the column density of C2H, we adopt an empirical approach instead of relying on computationally expensive full chemical network modeling methods (e.g., Bruderer et al. 2012; Du & Bergin 2014). Specifically, we utilize the DALI models presented in Bosman et al. (2021a), which were calibrated to reproduce the observed C2H emission levels of AS209, HD163296, and MWC480 using source-specific gas-grain thermochemical models (see their Fig. 2 and Fig. 5). Based on the model results of n (C2H) in the three sources and the input parameters in Bosman et al. (2021a), we summarize an empirical relation (see Appendix B)  \nlog 10 n (C2H) = FC -log 10 n (H) + 2 log 2 (C / O) (22)  \nwhere FC = 37 ± 1 is a fitting constant. Equation (22) allows us to estimate the C2H column density based on our simulation outputs. The dependence of n (H) is the result of two e ff ects: (a) lower CO column density, therefore less self-shielding, allowing more dissociation of CO; and (b) the depletion of the small dust, which scales with the total gas density, increasing the UV penetration. The third term on the RHS of the formula directly  \nFig. 3. (a) Time evolution of the planet mass. The horizontal gray line indicates the planet mass M 0 = 2 . 3 M J inferred by Liu et al. (2019), which we use as a reference mass. (b) Time evolution of the planet's gas accretion. (c) Time evolution of the planet's accretion luminosity (Eq. (15)). (d) Time evolution of the disk temperature at the Hill radius. The sublimation temperature of CH4 (green) and CO2 (orange) are marked by horizontal lines for reference, see discussion in Sect. 4.2. In each panel, the vertical dashed lines indicate the time when the mass of the planet reaches M 0.  \n<!-- image -->  \nreflects the strong dependence of C2H on the carbon-to-oxygen ratio C / O, as shown in Fig. 2 of Bosman et al. (2021a). Nevertheless, more comprehensive chemical modeling may rely on a series of additional assumptions and prescriptions, see our discussion in Section 4.1.  \nFinally, we remark that although we have chosen the fiducial values of our model parameters based on the observational constraint on MWC 480, the model we propose here is in principle applicable to any planet located exterior to its natal disk methane snowlines, see discussion in Sect. 4.2.", '3. Results': 'In this section, the results of the 2D hydrodynamical simulations that implement the model introduced in Fig. 2 are presented. The fiducial model is described in Sect. 3.1, and a parameter variation of the gas accretion rate of the planet is conducted in Sect. 3.2.', '3.1. Fiducial model': 'In our fiducial model, we include methane as the only ice-rich pebble. The time-dependent changes in both the mass and the accretion rate of the planet following the initiation of gas accretion are presented in panels (a) and (b) of Fig. 3.  \nFor the first 20,000 years ( ≈ 50 orbits) after the onset of runaway gas accretion, the accretion rate shows an exponential increase. During this period, the gas surface density at the location of the planet does not significantly decrease. The reason is that the timescale for the opening of the gas gap by the planet, which is the viscous timescale over the gap width (Kanagawa', 'et al. 2017)': "t gap = 2 . × 10 3 2 π Ω K M p / M ⋆ 10 -4 ! h 0 . 05 ! -3 . 5 GLYPH<18> α 5 × 10 -4 GLYPH<19> -1 . 5 (23)  \nis longer than the orbital timescale. The slope of the accretion rate as a function of time slightly changes around t = 20 , 000 yr, because of the transition from Hill radius scaling towards constant prefactor 0.14 in Eq. (19), causing the accretion rate to flatten. Since the planet has reached a mass above 1 M J at t ≃ 30 , 000 yr, the planet starts to sculpt the gas gap more strongly. Therefore, as the gas available around the planet decreases with the deeper and wider gap, the accretion rate of the planet goes down rapidly with a scaling of ˙ M p ∝ t -2 . A dotdashed line is presented in Fig. 3(b) for reference to show the trend.  \nAs depicted in panels (c) and (d) in Fig. 3, by using the planet's mass and gas accretion rate, we calculate the corresponding luminosity of the planet (Eq. (15)) and the temperature at its Hill radius  \nT Hill =       L acc 16 πσ R 2 Hill + T 4 d       1 4 =        G (3 M ⋆ ) 2 / 3 M 1 / 3 p ˙ M p 16 πσ R in r 2 + T 4 d        1 4 (24)  \n= 43 K M ⋆ M ⊙ ! 1 6 M p M J ! 1 12 ˙ M p 10 -4 M J yr -1 ! 1 4 R in 2 R J ! 1 4 GLYPH<18> r 73 au GLYPH<19> -1 2  \nwhere R Hill is the Hill radius of the planet. The second equation assumes that Td can be neglected. Due to the high accretion rate, the temperature inside the Hill radius is notably elevated compared to the background disk temperature, rising from 20 K up to approximately 45 K at the Hill sphere (see also Fig. 4(e)). Between t = 20 , 000 yr ( ≈ 50 orbits) and t =  \n<!-- image -->  \n<!-- image -->  \n<!-- image -->  \n<!-- image -->  \n<!-- image -->  \nFig. 4. Snapshot of the fiducial model at t = 75 751 yr (176 orbits), showing di ff erent components from the simulation. (a) Surface density of the gas (including vapor). Arrows indicate the velocity of the gas. (b) CH4 ice surface density. (c) Surface density ratio between CH4 vapor and total gas. (d) CH4 sublimated vapor surface density. (e) Disk midplane temperature. (f) C / O ratio in the gas phase. In each panel, the red dashed circle indicates the Hill radius of the planet. A video showing the time evolution can be found in the online material.  \n<!-- image -->  \n150 , 000 yr ( ≈ 350 orbits), balancing the increased planet mass and the reducing amount of available gas, the accretion rate of the planet remains in the range of 10 -5 to 10 -4 M J yr -1 during its growth, which roughly corresponds to the temperature range  \nwhere T Hill above the sublimation temperature of CH4 (the horizontal green line in Fig. 3(d), see Sect. 4.2 and Fig. 11), while is still below the sublimation temperature of CO2 (the horizontal orange line in Fig. 3(d), see also Sect. 4.2 and Appendix A). In  \nFig. 6. The radial profile of gas-phase C / O ratio as a function of time in the fiducial model. The horizontal white dashed line marks t = 176 orbits. The orange lines represent r 0 ± R Hill( t ), where R Hill( t ) the Hill radius of the planet at each time.  \n<!-- image -->  \nFig. 5. Azimuthally averaged radial profiles of the fiducial model at t = 75 751 yr (176 orbits). (a) Radial profile of total gas density (blue), CH4 ice (green dashed), and CH4 vapor (green solid) at 176 orbits. The gas phase C / O ratio (black) is calculated based on the abundance assumed in Table 1. (b) The expected C2H column density based on Eq. (22). The shaded region indicates the one-magnitude uncertainty of the fitting constant FC in Eq. (22). We mark the typical C2H peak density (orange) from the MAPS data constraint as a reference value for illustration.  \n<!-- image -->  \naddition, we show the temperature purely contributed from the planet by taking T d = 0 in Eq. (24), which is indicated by the red line in Fig. 3(d). As the figure shows, in the region where T > 25 K, the temperature at Hill radius T Hill is primarily determined by the planet's luminosity, rendering the results insensitive to the background temperature once it is below the CH4 sublimation temperature.  \nThe simulation snapshot of the surface densities of all the components, the temperature map, and the C / O ratio map in the gas phase at t = 64 , 560 yr (176 orbits) are displayed in Fig. 4, along with the corresponding azimuthally averaged radial density profiles shown in Fig. 5. At this time, the planet has reached a mass of 2.3 M J, consistent with the planet mass suggested by Liu et al. (2019), which we utilized as an educated estimate and point of reference. The Hill radius of the planet is marked by red circles in each panel. The top panel of Fig. 4 shows the azimuthally averaged radial profile of di ff erent components. A moderate gas gap has been opened by the planet (upper left panel of Fig. 4). And pebbles have followed the pressure gradient, resulting in dust trapping at the outer edge of the gas gap (upper right panel of Fig. 4). The dust gap is quickly emptied, while a small fraction of pebbles remain inside the co-orbital region of the planet, where a faint arc appears. This amount of dust is also shown in the radial profile Fig. 5. The opened gap allows for increased UV penetration, therefore, possibly higher temperature (Broome et al. 2022), which may help with the sublimation of the ice trapped inside the co-orbital region. Both of which is not included in our simulation, but see Section 4.1 for more discussion.  \nThe CH4 ice sublimates and transfers to vapor phase close to the planet location, where the temperature is higher compared with the disk background temperature. The majority of CH4 vapor is transported away from the planet through the gas spiral arm, as shown in panels (a) and (c) of Fig. 4. However, a significant fraction of vapor is transported from the vicinity of the planet into the gap as well, though not exactly into the co-orbital region. When considering the CH4 vapor mass fraction over the total gas density, the vapor fraction reaches a relatively high value of ≈ 10 -5 as shown in Fig. 4(c). Although the amount of CH4 vapor that ends up in the planet gap region is rather low, it is nevertheless su ffi cient to raise the C / O ratio above unity. Using the assumed constant CO abundance of 8 . 9 × 10 -6 and the CH4 vapor density, we calculate the C / O ratio in the gas phase in Fig. 4(f). The azimuthally averaged C / O ratio is indicated by the black solid line in the top panel of Fig. 5. Initially, the disk is cold and there is no methane released by the planet. Thus, both carbon and oxygen in the gas phase are determined by the non-condensable CO in our model. Therefore, the gasphase C / O ratio is unity throughout the entire simulation box at the beginning. As the CH4 -the major C-carrier volatile in our simulation- is released from ice-rich pebbles into the gas, the C / O ratio increases accordingly.  \nSignificant vapor release happens over a relatively short timescale at the beginning where T Hill is above the sublimation temperature of CH4, as shown in Fig. 3(d), and the gap remains shallow. In Fig. 6, we show the time evolution of the radial profiles of the C / Oin the gas phase. As the planet rapidly grows and gradually opens the gap, ice-rich pebbles can no longer reach the planet, and the sublimated vapor will be smeared out by radial mixing as shown in Fig. 6. In reality, the CH4 vapor may eventually be destroyed and converted into other species (e.g., Eistrup et al. 2016), or condense back to the ice. Both processes are not included in the simulations (see Sect. 4.1).  \nBoth high UV radiation and high C / O ratio are beneficial for C2H formation. With the simulation outputs, we estimate C2H  \ncolumn densities via Eq. (22) in the bottom panel of Fig. 5. A clear ring profile peaks inside the dust gap. We mark the peak C2H column densities of AS 209, HD 163296, and MWC 480 estimated from fitting the C2H lines hyperfine structure for reference (Guzmán et al. 2021), where the observational results are from the Molecules with ALMA at Planet-forming Scales (MAPS) large program (Öberg et al. 2021). For all three systems, the column density of C2H typically peaks ∼ 10 15 cm -2 at the local maximum. Due to the simplification of the empirical formula Eq. (22), we also show the 1 magnitude uncertainty of the fitting constant FC by the gray dash region. Our model result predicts a peak C2H column density value that is by-and-large consistent with the MAPS observation results. Moreover, the scenario described by the model is testable in the high spatial resolution result of ALMA, as a significant spatial correlation between the C2H rings and the continuum gap is expected. We discuss this in Sect. 4.4 and Sect. 4.5.", '3.2. Influence of the gas accretion rate': 'The precise gas accretion rates onto planets remain an open question. Studies di ff er on the accretion rates attained during the runaway phase with values ranging from ∼ 5 × 10 -9 to ∼ 3 × 10 -4 M J yr -1 (Machida et al. 2010; Szulágyi et al. 2014; Tanigawa et al. 2014; Tanigawa & Tanaka 2016; Homma et al. 2020; Szulágyi et al. 2022). To investigate how the accretion rate changes our results, we vary the gas accretion formula (Eq. (19)) by adopting di ff erent prefactor f A = 0 . 05 , 0 . 1 , 0 . 2 , 0 . 5 , 1 , 2 , 5 in the equation.  \nThe planet masses and accretion rates are shown in Fig. 7(a) and (b) separately, with di ff erent colors indicating di ff erent prefactors f A. We also show the gap depth (the gas surface density at the planet-opened gap Σ gap over the unperturbed surface density Σ g , 0) as functions of time in Fig. 7(c). At early times, when the gap remains shallow ( Σ gap / Σ g , 0 ∼ 1), all simulations show an exponentially growing accretion rate. The accretion rate increases with larger f A, resulting in a more massive planet within the same period after the onset of gas accretion. At the same time, the simulations with higher accretion rates open the gas gap faster, through a combination of increased accretion and stronger torques from the more massive planet pushing the gas away. As a result, it reaches the maximum accretion rate at an earlier turning point. In Fig. 7(d) and (e), the corresponding planet luminosity and temperature at the Hill radius T Hill at each run are calculated. The increased mass and accretion rate leads to a higher temperature around the planet, which reaches 50 K in simulation with f A = 2 and 60 K in simulation with f A = 5, causing more sublimation of ice-rich pebbles in the planet\'s vicinity within the same time. Since the accretion rates decay in a similar fashion, ˙ M p ∝ t -2 , after passing the peak, a higher accretion rate peak means that the temperature around the Hill radius of the planet remains higher than the CH4 sublimation temperature for a longer time.  \nWe present the maximum gas-phase C / Oratio versus time in Fig. 7(f). Initially, when the planet is accreting slowly, carbon in the gas phase is dominated by CO, the C / O ratio is unity, and UV photons cannot penetrate the shallow gap. These conditions are unfavorable for the formation of C2H rings. In simulations with faster accretion rates (more massive planets), the C / O in the gas phase rises faster and earlier. Yet, the maximum C / O ratio in the gas phase is limited by the (fixed) amount of carbon ice in the pebble reservoir. By virtue of the assumption that the total (gas + dust) CO and CH4 abundances are comparable, the maximum C / O ratio becomes ∼ 2 when methane ice fully sub-  \nmates into vapor. Additionally, a more massive planet opens a deeper gap, allowing more UV radiation penetration and promoting chemistry, which both facilitate C2H formation. Hence, a higher gas accretion rate is preferred to generate C2H rings.  \nConversely, in simulations with f A ≤ 0 . 2, the required growth timescale of the planet is much longer. The planet\'s mass grows slowly, resulting in a relatively low accreting luminosity, even at the accretion rate peak (Fig. 7(e)). Consequently, only a fraction of methane ice sublimates to vapor, leading to a peak C / O ratio of only 1.6 in the run with f A = 0 . 2, and even lower, barely exceeding unity, in the runs with f A = 0 . 1 and f A = 0 . 05. As the gap becomes deeper, pebbles are pushed away from the planet and trapped at the outer edge of the gap. In all simulations, the gas-phase C / O ratio around the continuum gap is then gradually smoothed out due to the viscous evolution of the gas, as shown in Fig. 6, and the peak value decreases.  \nWe then compare simulations with di ff erent prefactors f A ≥ 0 . 5 at various time steps: t = 40, 88, 176, and 402. These timesteps correspond to the point when the mass of the planet has just reached 2 . 3 M J in each run. Figure 8 displays the radial profiles of gas, ice, and vapor separately, as well as the correspondingly estimated C2H column density profiles. The gas profile and C / O ratio exhibit significant di ff erences among the simulations, but no apparent di ff erence in C2H production is observed in simulations with f A > 0 . 5. Planets with higher accretion rates reach M 0 more quickly and therefore have shallower gaps. Shallower gas gaps mean less UV penetration, which disfavors C2H formation. However, the higher luminosity and shallower gap allow for more methane to be released from ice-rich pebbles via sublimation, resulting in a higher C / O ratio. These two e ff ects balance each other, leading to similar C2H column densities across different accretion rate setups. As a result, inferring a planet\'s mass or luminosity from the radial profile of C2H can be challenging.  \nFor the other low accretion rate runs with f A ≤ 0 . 2, the planet mass has not reached 2 . 3 M J at the end of the simulation time of 1000 orbits, the result at the end of the simulation are shown in Fig. 9, where the planet masses are 1 . 66 M J, 0 . 98 M J and 0 . 47 M J in the simulation with f A = 0 . 2, 0.1, 0.05, respectively. Once the planet opens a su ffi cient gap, the accretion rate passes its peak turning point, as shown in Fig. 7(b), and the temperature at the Hill radius follows suit, as seen in Fig. 7(e). As the temperature around the planet cools down and the gas gap prevents further pebble supply, running the simulation further does not aid in methane release after the accretion rate peak is passed, as shown in Fig. 7(f). All of our simulations have passed the peak accretion rate / temperature point.  \nAs discussed above and shown in Fig. 9, if the accretion rate drops below f A ≤ 0 . 2, the results will be qualitatively di ff erent. Specifically, if the peak luminosity of the planet is too small, the disk will remain too cold to release any methane throughout the gas accretion period (Fig. 7(f)). Although the increased UV penetration might be helpful with increasing the C2H density inside the gap, the absolute column density of C2H would still be insu ffi cient to account for the observed values in ALMA observations, unless a very deep gap is opened (gap depth Σ gap / Σ g , 0 ≪ 0 . 01), as seen in Fig. 9 for the simulation with f A = 0 . 2. We acknowledge that this result is sensitive to the empirical formula used in Eq. (22), and high C / O ratios are essential for reproducing the observed value in the literature with more detailed chemical networks (Bosman et al. 2021b,a). The possibility that C2H formation is mainly due to high UV penetration is proposed by recent work (e.g., Calahan et al. 2023), which should be particularly interesting in transitional disks and some deep gas gaps. Since the majority of gas gaps found in, e.g., MAPS samples are  \nFig. 7. (a) Evolution of planet mass with di ff erent gas accretion rates ( f A factor in Eq. (19). The horizontal gray line indicates the planet mass M 0 = 2 . 3 M J inferred by Liu et al. (2019), which we use as a reference mass. (b) Planet\'s gas accretion. (c) Gas surface density at the gap over the unperturbed surface density Σ g , 0 as functions of time. (d) Planet luminosity (Eq. (15)). (e) Disk temperature at the Hill radius. (f) Peak gas-phase C / O ratio. The prefactors f A used in Eq. (19) are listed in the legend. The vertical dashed line indicates 176 orbits, when the planet mass has reached 0 . 04 , 0 . 05 , 0 . 12 , 1 . 2 , 2 . 3 , 3 . 8 , 6 . 8 M J in simulation with f A = 0 . 05 , 0 . 1 , 0 . 2 , 0 . 5 , 1 , 2 , 5 respectively.  \n<!-- image -->  \nrelatively shallow (gaps depth Σ gap / Σ g , 0 ≳ 0 . 1, Table 4 in Zhang et al. 2021), this routine is not the primary focus of this work.  \nThus, a dichotomy exists between "hot" and "cold" planets, determined by the maximum accretion rate the planet can reach. For the disk setup used in this study, a planet accretion rate above ∼ 10 -5 M J yr -1 is necessary to create a prominent C2H ring whose peak intensity is comparable to those detected in MAPS (Guzmán et al. 2021). This threshold is only weakly dependent on the planet mass M p, as Eq. (24) shows that T Hill ∝ M 1 / 12 p ˙ M 1 / 4 p . Besides, planet mass is always around M J at the accretion rate peak. Therefore, the maximum accretion rate that a planet can achieve will primarily determine the evolution of volatiles in its vicinity. Conversely, the observed spatial correlation between a C2H ring and a continuum ring can, in principle, constrain the accretion history of the putative planet.', '4. Discussion': 'Our findings suggest that the formation of a giant planet could have a substantial impact on the chemistry of protoplanetary disks by locally altering the elemental abundance in the gas and  \nice phases. In the following, we list the main caveats of this work and we broadly discuss some applications and implications of our model.', '4.1. Limitations': 'One major limitation of our study is that the hydrodynamical simulation is constrained within a 2D framework. To fully capture the complexity of planet accretion, it necessitates the use of high-resolution, 3D hydrodynamical simulations. Previous studies have demonstrated that the accretion rate and flow pattern may vary significantly in di ff erent model setups for both gas (e.g., Machida et al. 2010; Tanigawa et al. 2012; Fung et al. 2019) and solid materials (e.g., Tanigawa et al. 2014; Homma et al. 2020; Szulágyi et al. 2022). In particular, some major differences can be caused by the lack of meridional flow, which can only be probed in 3D simulation (e.g., Szulágyi et al. 2014; Fung &Chiang 2016; Teague et al. 2019). The meridional flow would trigger strong vertical and radial mixing, further complicating the elemental mixing and chemistry. Nevertheless, the computational cost of 3D high-resolution hydrodynamical simulations  \nFig. 9. Azimuthally averaged radial profiles of simulations with di ff erent accretion rates f A ≤ 0 . 2 at t = 1000 orbits. Panels are the same as Fig. 8. The amount of the released carbon is relatively low compared with the high accretion rate runs. High C2H column density might still be reached when the gap is very deep.  \n<!-- image -->  \nFig. 8. Azimuthally averaged radial profiles of simulations with different accretion rates f A ≥ 0 . 5 at M p = 2 . 3 M J. (a) Total gas density (solid), CH4 vapor density (dashed) at the indicated planet mass over the unperturbed density gas and pebble profile Eq. (12). (b) C / Oratio in the gas (midplane) in our simulations based on the abundance assumed in Table 1. (c) Expected C2H column density.  \n<!-- image -->  \nof gas accretion is substantial, hindering our ability to study the evolution of these systems using current computing resources.  \nIn addition, the lack of a vertical dimension in the simulations limits our ability to account for the stellar UV radiation, and renders the related disk irradiation temperature and UVphoto-chemistry uncertain. As recent works have shown, the planet-opened gap may also become warmer due to higher irradiation from the star (Booth et al. 2021; Pirovano et al. 2022; Broome et al. 2022). The parameterized temperature used in our simulation might only be a reference lower limit. Moreover, the growth and settling of pebbles can a ff ect the vertical transport of ice reservoirs and the depletion of small dust in the upper layer of the disk, which can also potentially elevate the C / O ratio of the gas and increase the UV penetration, as demonstrated by Van Clepper et al. (2022). Due to the above reasons, we opted for an empirical C2H column density estimation in this work (see Section 2.4 and Appendix B). The absence of chemical reactions does not impact our conclusion regarding C / O elevation as long as the reaction product remains in the gas phase. To better reproduce the ALMA observational results and understand the astrochemical evolution of protoplanetary disks, future follow-up  \ntests of our model featuring simultaneous hydrodynamic simulation and chemical network are worthwhile (e.g., see some pioneer e ff orts, Gong et al. 2017, 2023; Ilee et al. 2017; Booth & Ilee 2019; Krijt et al. 2020; Bergner & Ciesla 2021; Hu et al. 2023).  \nAnother caveat is that we do not include condensation of vapor in our simulation. Based on the simulation output in our fiducial model, we calculate the condensation timescales for methane vapor back to particles with size of ∼ 0.5 mm, the pebble size which is used in the hydrodynamical simulations (see Section 2). The condensation timescale reads (e.g., Schoonenberg & Ormel 2017)  \nτ con ≡ Σ vapor ˙ Σ vapor → ice = 1 8 r µ Z kBT mp s 2 p H g Σ pebble (25)  \nwhere µ Z = 16 m H is the mean molecular weight of methane vapor, and m H is the proton mass, Σ pebble is the surface density of the pebbles and H g is the gas scale height, and the typical particle mass follows mp = (4 / 3) π s 3 p ρ · . A map of the condensation timescale based on our simulation outputs is shown in Fig. 10 (a). In regions of high pebble density, outside the  \ngap region, the condensation timescale is shorter than the orbital timescale. In addition, the co-orbital region of the planet, where a moderate amount of pebbles linger, is characterized by short condensation timescales. Overall, however, due to the depletion of solids, the condensation timescale in the gap region is much longer than the orbital period, with maximum values of ∼ 10 5 yr. In particular, the C / O ratio is higher in the region with long condensation timescales, as illustrated by the green contour in Fig. 10, which indicates C / O = 1.5. Regions with condensation timescales exceeding 250 orbits (white) are associated with C / O > 1.5. Therefore, we conclude that accounting for the condensation of methane in the simulations will not prevent high C / O ratios in the gap, and therefore our major conclusions.  \nSimilarly, the sublimation timescale reads (see a similar formula for the desorption timescale in Piso et al. (2015))  \nτ sub ≡ Σ ice ˙ Σ ice → vapor = 1 8 √ 2 π s kBT µ Z mp s 2 p 1 P eq (26)  \nwhich strongly depends on the equilibrium pressure P eq (Eq. (17)). In Fig. 10 (b), the sublimation timescale displays a sharp transition around the planet Hill radius, within which the sublimation timescale becomes ≪ 1 orbit, much shorter than the local spreading time, validating our assumption of "instantaneous" sublimation. The quick sublimation around the planet and the long condensation timescale in the gap indicate that the methane vapor will first spread through the co-orbital region (which happens on a synodical timescale of several orbital periods) before the vapor has a chance to condense back the "cold" pebble at the midplane.  \nIn addition, the reduced dust optical depth in the gap enables radiation to penetrate deeper into the disk, more e ff ectively warming the gap (e.g., Calahan et al. 2021; Broome et al. 2022). Finally, through strong vertical transport such as meridional flows) (Tanigawa et al. 2012; Morbidelli et al. 2014), the released vapor may reach the upper layer of the disk, where temperatures are known to be higher (e.g., D\'Alessio et al. 1998). Therefore, the recondensation of CH4 vapor onto settled pebbles is likely to be ine ff ective in the gap region. It is worth noting that our results are not sensitive to vapor condensation onto grains of size ≪ mm. Assuming condensation timescale are short and grains perfectly coupled with the gas 1 the ice re-condensed on the surface of these small grains would nevertheless be exposed to su ffi ciently high doses of UV photons at the surface of the disk , releasing the carbon back into the gas.  \nOn the other hand, condensation is bound to happen in the colder regions exterior to the continuum gap, where densities are high and temperatures low. As we do not include condensation processes in our simulation, the methane vapor seen beyond 100 au and the corresponding plateau in the C / O ratio are artificial (see, e.g., Fig. 5). Hence, the peak in CH4 vapor at 100 au seen in Fig. 5a and the corresponding shoulder seen in Fig. 5(b) will, in reality, be absent. We have verified that the conclusions of this study, in particular, the peak C2H are una ff ected by its presence, however.  \nFinally, the model neglects any potential migration of the planet (see reviews of Kley & Nelson 2012; Paardekooper et al.  \n2022). Given the short chemical (Bosman et al. 2021b) and sublimation (desorption) timescale (Piso et al. 2015) of 10 4 yr in the disk region of our interests, as well the short gas accretion windows shown in our simulation, our results should be una ff ected by the planet migration happening over a timescale longer. Once the planet opens the gas gap, the migration speed of the planet is slower (e.g., Kanagawa et al. 2018), which further validates our assumption of a slow migration.', '4.2. Location of the planet relative to different snowlines': "In our scenario, the ability for an accreting planet to elevate the C / O ratio rest on two conditions: (1) the planet should be located outside the disk CH4 snowline; (2) the planet's accretion luminosity should be such to sublimate the CH4, but not O-carriers like CO2 or H2O. To understand how general these conditions are, we plot in Fig. 11 the saturation profiles of said molecules for the disk model adopted in this work. In Fig. 11 the solid curves show the gas-phase abundance ( ni / n (H)) one obtains by assuming complete saturation, i.e., when the partial pressure equals the equilibrium pressure of the vapor, PZ = P eq, where the equilibrium pressure for di ff erent species are taken from Fray & Schmitt (2009) as described in Appendix C. In other words, the curves give the maximum vapor abundance for each species at a given radius or corresponding temperature. It shows that at the initial background temperature of 20 K (73 au), all CO is in the gas phase, whereas CO2 sublimation remains insignificant even at the point where the highest T Hill is reached in the fiducial model.  \nIn particular, a local C / O > 1 gas enhancement will always appear whenever the planet is located beyond the methane snowline ( ≈ 40 au). As was shown in Eq. (24) and Fig. 3(d), the peak temperature is determined by the accretion luminosity of the planet, independent of the background disk temperature. Furthermore, The presence of the CO snowline does not influence the outcome of a C / O > 1 gas.  \nOn the other hand, accreting planets located between snowlines of species with higher sublimation temperatures may leave a di ff erent chemical footprint. For instance, an accreting planet located between the methane and carbon dioxide snowlines may, in turn, lower the C / O ratio in its vicinity by unleashing more oxygen than carbon through the local sublimation of carbon dioxide The basic tenet of this paper - an accreting planet can locally alter the gas chemistry by selective sublimation of ices applies generally.", '4.3. Exoplanet atmosphere and disk chemistry link': 'The relationship between the astrochemistry in protoplanetary disks and planetary atmospheres was first highlighted by Öberg et al. (2011). The authors propose that examining the C / Oratio in planetary atmospheres may allow us to trace the location where the majority of the planetary gas accretion happened, as the C / O of both gas and solid phases varies as ladders with increasing radius in the protoplanetary disk. This picture, in which giant planets inherited their atmosphere directly from the gaseous protoplanetary disk, has been actively developed in the past decade from both theoretical and observational aspects (e.g., Eistrup et al. 2018; Öberg & Wordsworth 2019; Booth & Ilee 2019; Cridland et al. 2019, 2020; GRAVITY Collaboration et al. 2020; Mollière et al. 2022).  \nConversely, our results show that giant planet formation may have significant local impacts on the astrochemistry of the disks.  \n<!-- image -->  \nFig. 10. Maps of (a) condensation timescale and (b) sublimation timescale of methane ice in the default simulation at t = 75751 yr (176 orbits), expressed in units of the orbital period. With the exception of the co-orbital region, the condensation timescale of methane in the gap region is much longer than the synodical timescale ∼ 10 orbits. This allows the gas to spread azimuthally. Since moderate amounts of pebbles may remain inside the co-orbital region, the condensation timescale could still be low there. Yet, the C / O ratio is higher in the region with long condensation timescales as shown in Fig. 4 (f), see also the green contour which indicates the C / O = 1.5 boundary. Regions with condensation timescales exceeding 250 orbits (white) are associated with C / O > 1.5. The red dashed circles indicate the Hill radius of the planet.  \n<!-- image -->  \nFig. 11. Location of snowlines at midplane for molecules of interest in this paper. Based on the midplane gas density at each radius, equating the partial pressure with Eq. (17), one can derive the saturation abundance of molecules (solid lines). The CO and CH4 abundance used in this work (8 . 9 × 10 -4 , gray horizontal line) is given as a reference.  \n<!-- image -->  \nThe gap opening and accretion of the planet can locally alter the chemical composition of the disks. Specifically, with a planet located outside the methane snowline, the water and carbon dioxide ice component of the accreted pebbles end up in the atmosphere of the planet, which would therefore become O-rich. In contrast, CH4 sublimates from pebbles to vapor, which renders the gap region C-rich. Our results indicate that giant planets actively shape their environment not only by opening gaps but also by changing the chemical inventory of disks, which extends the conventional understanding that the initial chemistry of natal disks is directly inherited by gas giants, as previously suggested by Öberg et al. (2011).', '4.4. Chemical rings in MWC 480': 'Among the five targets of MAPS, MWC 480 stands out as a particularly interesting object for our research. In the continuum, MWC480 features an annular gap, D76, at 76 au, potentially indicative of a planet (Liu et al. 2019). What makes the D76 gap special is that it is spatially coincident with almost all detected emission lines (14 out of the 18 species investigated in Law et al. (2021), see Fig. 12). Jiang et al. (2022) have found no statistically significant spatial correlations between line emission and continuum in other sources, except MWC 480. Intriguingly, such concentration of line-emission rings in a dust gap (cavity) configuration is also found in PDS 70 (Facchini et al. 2021).  \nFor these reasons, we have tailored the simulation parameters based on the observational constraints of MWC 480. Our findings can therefore be directly compared to the observations. In Fig. 12(c), we show the radial profiles of the ALMA continuum emission (Sierra et al. 2021) and the line emission map of C2H 3-2 (Law et al. 2021). By applying our model to MWC 480, the normalized C2H column density profile at t = 176 orbits in the fiducial model is shown in orange for comparison. We convolve the profile with a 0.15 arcsec beam, which is used in the MAPS data reduction process. Our results are therefore consistent with the observational data. It indicates that an accreting planet can be responsible for the formation of the observed C2H ring, in the D76 continuum gap of MWC 480 and possibly other molecular rings. Further investigation to solidify these findings, e.g., by including a chemical network in our simulations, is worthwhile to confirm these preliminary results.', '4.5. C2H rings in other systems': 'Apart from MWC480, ring-like C2H line-emission substructures are identified and studied in a number of other sources (Bergin et al. 2016; Facchini et al. 2021; Miotello et al. 2019; Bergner et al. 2021; Law et al. 2021). In these, dust rings and gaps are spatially resolved in the ALMA continuum as well. This raises an interesting question: can the model proposed in this paper also explain other systems?', '257 GHz Continuum (a)': '<!-- image -->  \nFig. 12. ALMA Observations of the MWC 480 disk. (a): 257 GHz mm dust continuum emission (Sierra et al. 2021). (b): Line emission moment 0 maps (Law et al. 2021) of six molecular transitions (selected for illustrative clarity; in total over 10 lines feature similar ring-like structures). Names of lines are labeled above each panel. The dashed ellipse marks the D76 continuum gap, and the solid ellipse marks the B98 continuum ring on the outside in all panels.  \n<!-- image -->  \nFig. 13. Azimuthally-averaged continuum emission profile (gray, normalized to unity at B98) and intensity profiles of C2H 3-2 (blue, normalized to unity at the peak). The normalized pebble surface densities and C2H column densities profile at t = 176 orbits in the default run are shown in orange and green for comparison  \n<!-- image -->  \n, which are convolved by 0.1 arcsec and 0.15 arcsec beams used in MAPS observation separately. The dashed and solid vertical lines coincide with the D76 gap and the B98 ring in the continuum, respectively.  \nWithin the MAPS samples, HD 163296 is the only other system where both C2H 3-2 and 2-1 lines present a sharp local maximum inside the D49 continuum gap. Besides, a C2H 3-2 ring peaks inside the D117 continuum gap of IM Lup, where Pinte et al. (2020) suggest a planet candidate depending on its localized kinematics perturbation (kink). And C2H 1-0 emission radially peaks inside the D100 continuum gap of AS 209, where a planet candidate is hinted by its kinematic signals Fedele et al. (2023) Therefore, our model is possibly applicable to these systems as well.  \nBy statistically studying the existence of C2H emission in disks and the radial profile of their continuum counterpart in 26 disks, van der Marel et al. (2021) has suggested a possible correlation between the presence of dust rings outside the CO snowline and the detection of C2H. It should be noted that MWC 480 is an outlier to the correlation, as there is no pebble ring outside its CO snowline, which is approximately at 100 au.  \nUnder the scenario proposed in this work, the accreting planet may sculpt a pebble ring outside its orbit, at the outer edge of the gap it opened. This ring is therefore located exte-  \nrior to the methane snowline. Moreover, it is likely to be situated near or beyond the CO snowline as the scenario proposed in this study works at any location beyond the CH4 snowline. Therefore, our model, which examines the e ff ect of an accreting planet outside the methane snowline on the chemistry of the disk, could potentially provide an explanation for and further complete the tentative correlation identified by van der Marel et al. (2021) between the presence of dust rings outside the CO snowline and the detection of C2H.  \nThe sample size of disks with detected C2H remains small, and the available spatial resolution for testing the correlation is limited, which is only possible in MAPS samples. Future observations, such as the ongoing ALMA Disk-Exoplanet C / Onnection (DECO) large program, will shed more light on this issue by assisting sample selection. For instance, in the low spatial resolution date of Miotello et al. (2019), two sources Sz 71 (GWLup) and Sz129, exhibit tentative features in their C2H 32 emission. In both cases, the emission plausibly radially peaks at the D74 and D64 continuum gaps respectively. Interestingly, localized kinematic perturbations consistent with the presence  \nof planet candidates have been detected within these gaps (Pinte et al. 2020). Deeper observation with a higher spatial resolution is required to further confirm the link and will be useful for further study.', '4.6. Detectablity of the hypothesized planet': 'Observed rings and gaps in protoplanetary disks are often interpreted as evidence of planet formation (e.g., Pérez et al. 2019; Toci et al. 2020). Correspondingly, a growing number of embedded planets are being claimed within these disks as well (e.g., Hammond et al. 2023; Pinte et al. 2023). Therefore, it is important to ask the question where the planet is located? A compelling feature of the planet proposed in our model is that, in order to heat up the surrounding area, it is currently or very recently in the process of accreting.  \nFor instance, MWC 480 provides an apt illustration of this challenge. Although the Strategic Exploration of Exoplanets and Disks with Subaru (SEEDS) survey has succeeded in directly imaging this object (Uyama et al. 2017), its relatively bright Herbig Ae / Be star status - with an H band magnitude of 6.26 - means that detection is still an arduous task. Moreover, identifying any planets in close proximity to the dust ring around MWC480 is made all the more di ffi cult due to the possibility of confusion with disk features that may be influenced by advanced image analysis techniques (Rameau et al. 2017; Ha ff ert et al. 2019).  \nDetecting planetary accretion signals provides another direct evidence of planets. If a planet accretes gas from the disk, H α line emission originates from the accretion front at the planetary surface. According to the ESO archive, VLT / MUSE (Bacon et al. 2010) observations have been carried out for more than 50 disks in the past 5 years to search for H α emission (e.g., Xie et al. 2020); yet, so far, PDS 70 is the only one in which accreting planets have been identified in MUSE (Ha ff ert et al. 2019). A simple explanation, perhaps, is that many planets in disks are accreting too slowly. Therefore, they do not produce detectable H α line emission (Marleau et al. 2022). We might have been fortunate with PDS 70 b,c, in which the two planets are just above the detection limit. Alternatively, the formation of rings and gaps structure can be independent of the existence of planets (e.g, Bai & Stone 2014; Flock et al. 2015; Okuzumi et al. 2016; Owen & Kollmeier 2019; Ohashi et al. 2021; Jiang & Ormel 2021; Kuznetsova et al. 2022; Tominaga et al. 2022), making the problems more complex.  \nAs discussed in Section Section 3.2, our rough calculation may face challenges in distinguishing di ff erent accretion rate histories above 10 -5 M J yr -1 . Additionally, a series of caveats need to be acknowledged, as outlined in Sect. 4.1. Despite these limitations, making an educated guess for further investigation is worthwhile. By selecting a mass of 2 . 3 M J and the corresponding accretion rate of 3 × 10 -5 M J yr -1 in our fiducial model, we estimate an H α flux from the planet of f H α = 1 . 2 × 10 -15 erg s -1 cm -2 . This estimate is based on the relationship between H α emission and accretion rates presented in Figure 10 of Marleau et al. (2022) and is reasonably above the MUSE detection limit. Our models, therefore, proposed a novel testable approach for searching for planets in disks through their chemistry signals. A searching e ff ort for the potential accreting planets in the future might be also been intriguing in other instruments like MagAO-X (Males et al. 2018) SPHERE / ZIMPOL (Cugno et al. 2019), SCExAO / VAMPIRES (Uyama et al. 2020), and HST (Zhou et al. 2021).', '5. Conclusions': "In this work, with the newly-developed phase change module (Wang et al. 2023) of the Athena++ code, we perform global multi-fluid hydrodynamic simulations accounting for the sublimation process of the volatile components of (pebble-)accreting planets in a typical protoplanetary disk. The gas accretion of the planet is parameterized with the Machida et al. (2010) formula. Our simulations demonstrate how an accreting planet locally heats its surroundings, creates a gap in the disk, and establishes conditions conducive to C-photochemistry. Starting with a 10 M ⊕ core, the planet grows into a Jovian mass planet with accretion rate ∼ 10 -5 M J yr -1 in our fiducial setup. A su ffi ciently deep gas gap with a gap depth ∼ 30 % of the unperturbed disk is opened at the moment when the planet reaches 2 . 3 M J (Fig. 5). Such a bright planet may locally heat its vicinity, elevating the temperature at its Hill radius to ∼ 45 K. A sketch of our model is shown in Fig. 2. Our main findings are:  \n- 1. The enhanced temperature sublimates the major C-carrier ice, methane, surrounding the planet but is not high enough for water and carbon dioxide ice to sublimate. The released methane locally raises the carbon-to-oxygen ratio inside the gap from unity (CO-dominant gas) to ∼ 2. The opened gap allows more UV penetration towards the planet's neighborhood. Both high C / O ratio and high UV radiation assist with C-bearing radical (C2H) formation and subsequent assembly of organic compounds.\n- 2. As a result, a C2H ring is expected to emerge at the dust / gas gap location. The results o ff er an explanation for the emission of organic molecular lines observed by ALMA in the MAPS large program, where the majority of detected line emission (including C2H, see Fig. 12) peaks inside the D76 continuum gap of MWC 480. This model may also apply to the C2H ring peak at the D49 continuum gap of HD 163296.\n- 3. The outlined model operates when an accreting planet is located beyond the disk CH4 snowline. It predicts the formation of a C2H ring at the planet's location and a pebble ring exterior to it. As the planet may be located beyond the CO snowline as well, the model can potentially account for the correlation between the presence of dust rings outside the CO snowline and the detection of C2H, as observed by van der Marel et al. (2021), see Section 4.5.\n- 4. This work demonstrates that the accretion of a giant planet can significantly impact its surrounding astrochemistry environment by locally altering the chemical composition of disks. Due to the di ff erent sublimation temperatures, the hot accreting planets can release some volatiles (e.g. methane) while incorporating others (e.g. water), leading to di ff erent chemical impacts depending on where the planet is located. This study extends the conventional understanding that the initial chemistry of disks is simply inherited by gas giants, by postulating that planets actively shape their chemical environment of formation.\n- 5. The finding that a high C / O ratio is the consequence of a gap-opening planet can be inverted to give a rough constraint on the accreting luminosity of the planet. We found that an accretion rate above 10 -5 M J yr -1 is necessary to ensure that the planet is hot enough to release methane, elevate the C / O ratio, and create a prominent C2H ring.\n- 6. For reference, in the fiducial model, the planet mass reaches 2 . 3 M J mass when the accreting rate is 3 × 10 -5 M J yr -1 . The H α flux from the planet is around f H α = 1 . 2 × 10 -15 erg s -1 cm -2 , which is potentially detectable  \nvia instruments such as VLT / MUSE, SCExAO / VAMPIRES, SPHERE / ZIMPOL, and MagAO-X.  \nThe link between accreting planets and disk astrochemistry proposed by this work is testable with current optical (see Sect. 4.6), IR (JWST), and (sub)mm (ALMA) wavelength facilities. Conducting planet-hunting campaigns in disks with pronounced chemical rings and continuum gap correlation might therefore be a new avenue to hunt and confirm (proto)planets. At the same time, further investigation into both the spatially resolved disk chemistry and characterization of the atmosphere of the protoplanet in confirmed planet-forming (e.g., for PDS 70, Facchini et al. 2021; Cridland et al. 2023) is of importance to understand how the composition of the planet is built upon, how it influences the disk chemistry in reality, and possibly revealing the formation history of the planet.  \nSimultaneously, we plan to extend the model to 3D and apply it towards more sources. We also aim to couple our hydroand thermo-dynamical code with a (reduced) on-the-fly chemical network to obtain more quantitative constraints on chemical interactions between the planet and disk. Further investigation will focus on how the accretion of the planet shapes their atmospheric chemistry, how the e ff ect of planet migration influences the chemical signatures, as well as the impact of multiple planets. Overall, our proposed model opens up exciting new avenues for studying the co-evolution of the disks and protoplanets in planet-forming disks.  \nAcknowledgements. We thank the anonymous referee for their thoughtful and constructive report that improved the initial manuscript. H.J. and Y.W. would like to thank Pinghui Huang for fruitful discussions about Athena++ code. H.J. thanks Bin Ren and Chen Xie for insightful discussions. H.J. highly appreciates valuable comments from Ilse Cleeves, Alex Cridland, Stefano Facchini, Enrique Macías, Anna Miotello, Giovanni Rosotti, and Claudia Toci. C.W.O. acknowledges support by the National Natural Science Foundation of China (grants no. 12250610189 and 12233004). The authors acknowledge the Tsinghua Astrophysics High-Performance Computing platform at Tsinghua University for providing computational and data storage resources that have contributed to the research results reported within this paper. This work has used Astropy (Astropy Collaboration et al. 2013), Matplotlib (Hunter 2007), Numpy (Harris et al. 2020), Scipy (Virtanen et al. 2020) software packages.", 'References': "Alarcón, F., Bosman, A. D., Bergin, E. A., et al. 2021, ApJS, 257, 8  \nAndrews, S. M., Huang, J., Pérez, L. M., et al. 2018, ApJ, 869, L41  \n- Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33\n- Bacon, R., Accardo, M., Adjali, L., et al. 2010, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 7735, Ground-based and Airborne Instrumentation for Astronomy III, ed. I. S. McLean, S. K. Ramsay, & H. Takami, 773508\n- Bae, J., Teague, R., Andrews, S. M., et al. 2022, ApJ, 934, L20\n- Bai, X.-N. & Stone, J. M. 2014, ApJ, 796, 31\n- Barnett, M. N. & Ciesla, F. J. 2022, ApJ, 925, 141\n- Bast, J. E., Lahuis, F., van Dishoeck, E. F., & Tielens, A. G. G. M. 2013, A&A, 551, A118\n- Bergez-Casalou, C., Bitsch, B., Pierens, A., Crida, A., & Raymond, S. N. 2020, A&A, 643, A133\n- Bergin, E. A., Du, F., Cleeves, L. I., et al. 2016, ApJ, 831, 101\n- Bergner, J. B. & Ciesla, F. 2021, ApJ, 919, 45\n- Bergner, J. B., Öberg, K. I., Bergin, E. A., et al. 2019, ApJ, 876, 25\n- Bergner, J. B., Öberg, K. I., Guzmán, V. V., et al. 2021, ApJS, 257, 11\n- Birnstiel, T., Dullemond, C. P., Zhu, Z., et al. 2018, ApJ, 869, L45\n- Booth, A. S., Ilee, J. D., Walsh, C., et al. 2023, A&A, 669, A53\n- Booth, A. S., Walsh, C., Terwisscha van Scheltinga, J., et al. 2021, Nature Astronomy, 5, 684\n- Booth, R. A. & Ilee, J. D. 2019, MNRAS, 487, 3998\n- Bosman, A. D., Alarcón, F., Bergin, E. A., et al. 2021a, ApJS, 257, 7\n- Bosman, A. D., Alarcón, F., Zhang, K., & Bergin, E. A. 2021b, ApJ, 910, 3\n- Broome, M., Shorttle, O., Kama, M., & Booth, R. A. 2022, arXiv e-prints, arXiv:2204.11555  \nArticle number, page 16 of 18  \n- Bruderer, S., van Dishoeck, E. F., Doty, S. D., & Herczeg, G. J. 2012, A&A, 541, A91\n- Calahan, J. K., Bergin, E., Zhang, K., et al. 2021, ApJ, 908, 8\n- Calahan, J. K., Bergin, E. A., Bosman, A. D., et al. 2023, Nature Astronomy, 7, 49\n- Cleeves, L. I., Bergin, E. A., & Harries, T. J. 2015, ApJ, 807, 2  \nCrida, A. & Bitsch, B. 2017, Icarus, 285, 145  \n- Cridland, A. J., Facchini, S., van Dishoeck, E. F., & Benisty, M. 2023, arXiv e-prints, arXiv:2303.17899\n- Cridland, A. J., van Dishoeck, E. F., Alessi, M., & Pudritz, R. E. 2019, A&A, 632, A63\n- Cridland, A. J., van Dishoeck, E. F., Alessi, M., & Pudritz, R. E. 2020, A&A, 642, A229\n- Cugno, G., Quanz, S. P., Hunziker, S., et al. 2019, A&A, 622, A156  \nCurrie, T., Lawson, K., Schneider, G., et al. 2022, Nature Astronomy, 6, 751 D'Alessio, P., Cantö, J., Calvet, N., & Lizano, S. 1998, ApJ, 500, 411 Dong, R., Rafikov, R. R., Stone, J. M., & Petrovich, C. 2011, ApJ, 741, 56 Dong, R., Zhu, Z., Rafikov, R. R., & Stone, J. M. 2015, ApJ, 809, L5 Drozdovskaya, M. N., Walsh, C., van Dishoeck, E. F., et al. 2016, MNRAS, 462,  \n- 977\n- Du, F. & Bergin, E. A. 2014, ApJ, 792, 2  \nEistrup, C., Walsh, C., & van Dishoeck, E. F. 2016, A&A, 595, A83 Eistrup, C., Walsh, C., & van Dishoeck, E. F. 2018, A&A, 613, A14 Facchini, S., Teague, R., Bae, J., et al. 2021, AJ, 162, 99 Fedele, D., Bollati, F., & Lodato, G. 2023, arXiv e-prints, arXiv:2302.11028 Flock, M., Ruge, J. P., Dzyurkevich, N., et al. 2015, A&A, 574, A68 Follette,  \n- K. B., Close, L. M., Males, J. R., et al. 2022, arXiv e-prints, arXiv:2211.02109  \nFray, N. & Schmitt, B. 2009, Planet. Space Sci., 57, 2053 Fung, J. & Chiang, E. 2016, ApJ, 832, 105 Fung, J., Zhu, Z., & Chiang, E. 2019, ApJ, 887, 152 Gong, M., Ho, K.-W., Stone, J. M., et al. 2023, arXiv e-prints, arXiv:2305.04965 Gong, M., Ostriker, E. C., & Wolfire, M. G. 2017, ApJ, 843, 38 GRAVITY Collaboration, Nowak, M., Lacour, S., et al. 2020, A&A, 633, A110 Guzmán, V. V., Bergner, J. B., Law, C. J., et al. 2021, ApJS, 257, 6 Guzmán, V. V., Pety, J., Goicoechea, J. R., et al. 2015, ApJ, 800, L33 Ha ff ert, S. Y., Bohn, A. J., de Boer, J., et al. 2019, Nature Astronomy, 3, 749 Hammond, I., Christiaens, V., Price, D. J., et al. 2023, MN-  \nRAS[  \narXiv:2302.11302  \n]  \nHarris, C. R., Millman, K. J., van der Walt, S. J., et al. 2020, Nature, 585, 357  \nHomma, T., Ohtsuki, K., Maeda, N., et al. 2020, ApJ, 903, 98  \nHori, Y. & Ikoma, M. 2011, MNRAS, 416, 1419  \n- Hu, X., Li, Z.-Y., Wang, L., Zhu, Z., & Bae, J. 2023, MNRAS, 523, 4883 Huang, P. & Bai, X.-N. 2022, ApJS, 262, 11\n- Huélamo, N., Chauvin, G., Mendigutía, I., et al. 2022, A&A, 668, A138\n- Hunter, J. D. 2007, Computing in Science and Engineering, 9, 90\n- Ilee, J. D., Forgan, D. H., Evans, M. G., et al. 2017, MNRAS, 472, 189\n- Izquierdo, A. F., Facchini, S., Rosotti, G. P., van Dishoeck, E. F., & Testi, L. 2022, ApJ, 928, 2\n- Jiang, H. & Ormel, C. W. 2021, MNRAS, 505, 1162\n- Jiang, H., Zhu, W., & Ormel, C. W. 2022, ApJ, 924, L31\n- Johansen, A., Ronnet, T., Bizzarro, M., et al. 2021, Science Advances, 7, eabc0444\n- Kama, M., Bruderer, S., van Dishoeck, E. F., et al. 2016, A&A, 592, A83\n- Kanagawa, K. D., Tanaka, H., Muto, T., & Tanigawa, T. 2017, PASJ, 69, 97  \nKanagawa, K. D., Tanaka, H., & Szuszkiewicz, E. 2018, ApJ, 861, 140  \n- Keppler, M., Benisty, M., Müller, A., et al. 2018, A&A, 617, A44\n- Kley, W. & Nelson, R. P. 2012, ARA&A, 50, 211\n- Krijt, S., Bosman, A. D., Zhang, K., et al. 2020, ApJ, 899, 134  \nKuznetsova, A., Bae, J., Hartmann, L., & Low, M.-M. M. 2022, ApJ, 928, 92  \nLaw, C. J., Loomis, R. A., Teague, R., et al. 2021, ApJS, 257, 3  \nLiu, Y., Dipierro, G., Ragusa, E., et al. 2019, A&A, 622, A75  \nLong, F., Pinilla, P., Herczeg, G. J., et al. 2018, ApJ, 869, 17  \n- Machida, M. N., Kokubo, E., Inutsuka, S.-I., & Matsumoto, T. 2010, MNRAS, 405, 1227\n- Males, J. R., Close, L. M., Miller, K., et al. 2018, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 10703, Adaptive Optics Systems VI, ed. L. M. Close, L. Schreiber, & D. Schmidt, 1070309 Marleau, G. D., Aoyama, Y., Kuiper, R., et al. 2022, A&A, 657, A38 Miotello, A., Facchini, S., van Dishoeck, E. F., et al. 2019, A&A, 631, A69 Mollière, P., Molyarova, T., Bitsch, B., et al. 2022, ApJ, 934, 74 Morbidelli, A., Szulágyi, J., Crida, A., et al. 2014, Icarus, 232, 266 Mordasini, C., Alibert, Y., Georgy, C., et al. 2012, A&A, 547, A112 Müller, T. W. A., Kley, W., & Meru, F. 2012, A&A, 541, A123 Öberg, K. I. & Bergin, E. A. 2021, Phys. Rep., 893, 1 Öberg, K. I., Guzmán, V. V., Walsh, C., et al. 2021, ApJS, 257, 1 Öberg, K. I., Murray-Clay, R., & Bergin, E. A. 2011, ApJ, 743, L16 Öberg, K. I. & Wordsworth, R. 2019, AJ, 158, 194\n- Ohashi, S., Kobayashi, H., Nakatani, R., et al. 2021, ApJ, 907, 80  \n- Okuzumi, S., Momose, M., Sirono, S.-i., Kobayashi, H., & Tanaka, H. 2016, ApJ, 821, 82\n- Ormel, C. W., Vazan, A., & Brouwers, M. G. 2021, A&A, 647, A175\n- Owen, J. E. & Kollmeier, J. A. 2019, MNRAS, 487, 3702\n- Paardekooper, S.-J., Dong, R., Du ff ell, P., et al. 2022, arXiv e-prints, arXiv:2203.09595  \nPérez, S., Casassus, S., Baruteau, C., et al. 2019, AJ, 158, 15  \n- Perri, F. & Cameron, A. G. W. 1974, Icarus, 22, 416\n- Pinte, C., Hammond, I., Price, D. J., et al. 2023, MNRAS[ arXiv:2301.08759 ]\n- Pinte, C., Price, D. J., Ménard, F., et al. 2020, ApJ, 890, L9\n- Pinte, C., van der Plas, G., Ménard, F., et al. 2019, Nature Astronomy, 3, 1109\n- Pirovano, L. M., Fedele, D., van Dishoeck, E. F., et al. 2022, A&A, 665, A45\n- Piso, A.-M. A., Öberg, K. I., Birnstiel, T., & Murray-Clay, R. A. 2015, ApJ, 815, 109\n- Rafikov, R. R. 2006, ApJ, 648, 666\n- Rameau, J., Follette, K. B., Pueyo, L., et al. 2017, AJ, 153, 244\n- Robert, C. M. T., Crida, A., Lega, E., Méheut, H., & Morbidelli, A. 2018, A&A,", '617, A98': '- Rosotti, G. P. 2023, New A Rev., 96, 101674\n- Rosotti, G. P., Juhasz, A., Booth, R. A., & Clarke, C. J. 2016, MNRAS, 459, 2790  \nSchneider, A. D. & Bitsch, B. 2021, A&A, 654, A71  \nSchoonenberg, D. & Ormel, C. W. 2017, A&A, 602, A21  \n- Shakura, N. I. & Sunyaev, R. A. 1973, A&A, 500, 33\n- Sierra, A., Pérez, L. M., Zhang, K., et al. 2021, ApJS, 257, 14\n- Speedie, J., Booth, R. A., & Dong, R. 2022, ApJ, 930, 40\n- Speedie, J. & Dong, R. 2022, ApJ, 940, L43\n- Szulágyi, J., Binkert, F., & Surville, C. 2022, ApJ, 924, 1\n- Szulágyi, J., Morbidelli, A., Crida, A., & Masset, F. 2014, ApJ, 782, 65\n- Tanigawa, T., Maruta, A., & Machida, M. N. 2014, ApJ, 784, 109\n- Tanigawa, T., Ohtsuki, K., & Machida, M. N. 2012, ApJ, 747, 47\n- Tanigawa, T. & Tanaka, H. 2016, ApJ, 823, 48\n- Teague, R., Bae, J., Aikawa, Y., et al. 2021, ApJS, 257, 18\n- Teague, R., Bae, J., & Bergin, E. A. 2019, Nature, 574, 378\n- Toci, C., Lodato, G., Fedele, D., Testi, L., & Pinte, C. 2020, ApJ, 888, L4\n- Tominaga, R. T., Tanaka, H., Kobayashi, H., & Inutsuka, S.-i. 2022, ApJ, 940, 152\n- Uyama, T., Hashimoto, J., Kuzuhara, M., et al. 2017, AJ, 153, 106\n- Uyama, T., Norris, B., Jovanovic, N., et al. 2020, Journal of Astronomical Telescopes, Instruments, and Systems, 6, 045004\n- Van Clepper, E., Bergner, J. B., Bosman, A. D., Bergin, E., & Ciesla, F. J. 2022, ApJ, 927, 206\n- van der Marel, N., Bosman, A. D., Krijt, S., Mulders, G. D., & Bergner, J. B. 2021, A&A, 653, L9\n- van der Marel, N., Dong, R., di Francesco, J., Williams, J. P., & Tobin, J. 2019, ApJ, 872, 112\n- Virtanen, P., Gommers, R., Oliphant, T. E., et al. 2020, Nature Methods, 17, 261 Wang, Y., Ormel, C. W., Huang, P., & Kuiper, R. 2023, arXiv e-prints, arXiv:2306.06169\n- Xie, C., Ha ff ert, S. Y., de Boer, J., et al. 2020, A&A, 644, A149 Youdin, A. N. & Lithwick, Y. 2007, Icarus, 192, 588\n- Zhang, K., Booth, A. S., Law, C. J., et al. 2021, ApJS, 257, 5\n- Zhang, S., Zhu, Z., Huang, J., et al. 2018, ApJ, 869, L47\n- Zhou, Y., Bowler, B. P., Wagner, K. R., et al. 2021, AJ, 161, 244\n- Zhou, Y., Sanghi, A., Bowler, B. P., et al. 2022, ApJ, 934, L13\n- Zhu, Z., Nelson, R. P., Hartmann, L., Espaillat, C., & Calvet, N. 2011, ApJ, 729, 47', 'Appendix A: Impact of O-carrier volatiles': 'The key of our work relies on the local elevation of the C / O ratio triggered by the accretion luminosity of the planet. In our default setup, we only consider methane in the volatile component. However, oxygen is present in pebbles as well in the form of CO2 and H2O ice. Consequentially, one would expect that oxygen-bearing carriers may also be released from the volatile to the gas phase, which would in contrast lower the local C / O ratio.  \nAlthough it is unlikely to happen due to the relatively high sublimation temperature, see Sect. 4.2, we conducted a comparison run by assuming that all icy volatiles in our model is CO2 instead while keeping the other setups the same as the fiducial model. The equilibrium pressure Eq. (17) of CO2 follows Ta = 2571 K and P eq , 0 = 2 . 57 × 10 12 g cm -1 s -2 (Fray & Schmitt 2009) in the simulation. The outcome is straightforward. The typical sublimation temperature of the CO2 ice is ≈ 55 K (the orange line in Fig. 3(d) and see Fig. 11), significantly higher than methane. These temperatures can only be reached in the inner regions of the Hill sphere, but never at the Hill radius. Therefore, as T Hill < 50 K, CO2 is almost entirely accreted by the protoplanet and cannot be unleashed into the disk.  \nSince water ice has even higher sublimation temperatures (Fig. 11) H2O will also stay on the pebbles in ice form, and be lost to the planet before it sublimates. As a result, we can confidently assert that our key findings remain unchanged regardless of the water content of the pebble, or indeed any other Ocarrier with a significantly higher sublimation temperature than methane.', 'Appendix B: The C 2 H column densities formula': 'In the study by Bosman et al. (2021a), a gas-grain chemical network was employed to calculate the 2D ( r -z ) chemical makeup of three MAPS disks. However, due to the lack of a vertical dimension in our simulation, it is challenging to directly compare our results with those of Bosman et al. (2021a), where the C2H abundance is dominant by photo-chemistry in the upper layer of the disk (e.g., see their Appendix C). Nevertheless, we outline the key features of the Bosman et al. (2021a) model outputs and motivate an empirical fit formula that links the C / O ratio to the C2H abundance, Eq. (22).  \nBosman et al. (2021a) concluded that the resolved observations of high but localized C2H column densities suggest the need for locally higher C / O ratios at C2H ring location. Additionally, a correlation is observed between the C / O ratio and CO abundance, indicating that a higher CO abundance corresponds to a lower C / O ratio, vice verse.  \nHere, we provide a description of how the empirical formula Eq. (22) is determined:  \nlog 10 n (C2H) = FC -log 10 n (H) + 2 log 2 (C / O)  \nThe dependency of n (C2H) on n (H), and therefore log 10 n (CO) ≡ log 10 n (H) + 5, is motivated by the anti-correlation observed in Figure 1, Figure 2, and Figure 5 of Bosman et al. (2021a). These figures demonstrate an anti-correlation between CO column densities and C2H column densities. In particular, Figure 5 shows that the CO column density at the gap location in AS 209 ( ∼ 10 18 cm -2 ) is lower than in HD 163296 and MWC480( ∼ 10 20 cm -2 ). However, in Figure 2, the C2H column density in AS 209 ( ∼ 10 14 cm -2 ) is higher than in HD 163296 and MWC 480 ( ∼ 10 12 cm -2 ) when C / O = 1. Thus, a value of FC = 37 is obtained as the sum of these contributions.  \nThe dependency of n (C2H) on log 2 (C / O) is estimated based on Figure 2 of Bosman et al. (2021a), which shows that as the C / O ratio increases (or decreases) from 1.0 to 2.0 (or 0.47), the C2H column densities roughly rise (or drop) by approximately two orders of magnitude.  \nConsidering the moderate C2H column densities variation observed in Figure 2 when taking di ff erent small dust depletion factors, and the fitting values were visually identified, we leave an uncertainty of approximately one order of magnitude in FC .', 'Appendix C: Thermodynamic properties of different molecules': 'The Clausius-Clapeyron equation is derived from the equality of the chemical potentials between a pure gas-phase compound and its solid phase, which occurs when they are in equilibrium. Taking the perfect gas approximation and neglecting the molar volume of solids, the Clausius-Clapeyron equation can be written as (Fray & Schmitt 2009)  \nln( P eq) = ln( P 0) + Z T T 0 ∆ H eq( T ) RT 2 d T (C.1)  \nwhere P 0 is the equilibrium (sublimation) pressure at temperature T 0 and  \nd ∆ H eq( T ) d T = ∆ Cp ( T ) = ∆ Cp , gas( T ) -∆ Cp , ice( T ) (C.2)  \nis the equilibrium (sublimation) enthalpy which depends on the heat capacities of gas Cp , gas and ice Cp , ice. If the equilibrium enthalpy is supposed to be independent of temperature, then the equilibrium pressure follows the simple relation  \nln( P eq) = A 0 + A 1 / T (C.3)  \nwhich can be linked to Eq. (17) by P eq , 0 = e A 0 and Ta = -A 1. We, therefore, take the value of A 0 and A 1 from Table 5 of Fray & Schmitt (2009) to calculate the P eq , 0 and Ta used in our simulation. The original empirical extrapolations formula in Fray & Schmitt (2009) follows ln( P eq) = P n i = 0 AiT -i , yet we only take the first two terms to keep the formula still physically meaningful.'}