Open Access
Issue
A&A
Volume 635, March 2020
Article Number A94
Number of page(s) 16
Section Interstellar and circumstellar matter
DOI https://doi.org/10.1051/0004-6361/201732243
Published online 17 March 2020

© E. Di Folco et al. 2020

Licence Creative Commons
Open Access article, published by EDP Sciences, under the terms of the Creative Commons Attribution License (http://creativecommons.org/licenses/by/4.0), which permits unrestricted use, distribution, and reproduction in any medium, provided the original work is properly cited.

1 Introduction

The evolution of circumstellar discs from the initial reservoir of the interstellar medium (ISM)–dust and gas mixture towards planetary systems is governed by the growth and coagulation of dust grains, viscous spreading and accretion onto the central star (or on proto-planets), and dissipation of the primordial gas (Williams & Cieza 2011). Although a large amount of high-resolution images and spectroscopic data have been collected during recent decades, the various stages of disc evolution remain poorly constrained by observations (Wyatt et al. 2015). The main accretion phase corresponds to the Class II phases for T Tauri (TTS) and Herbig Ae (HAe) stars. The final phase is represented by debris discs, when the primordial reservoir has been fully dissipated and/or incorporated into planetesimals and planets, and the detected emission from second-generation dust comes from collisions between the leftovers of planet formation or comet evaporation (Wyatt 2008; Hughes et al. 2018). Evidence for structural changes has been collected for a large number of pre-main sequence stars at various stages of the disc evolution (e.g. Andrews et al. 2018; Avenhaus et al. 2018; Garufi et al. 2018). This includes the presence of spirals, rings and gaps, central cavities, or azimuthal asymmetries, where different processes are at play such as the protoplanet(s)–disc gravitational or hydrodynamical interaction, viscous evolution, grain growth and opacity change, gas-grain hydrodynamical interaction, the development of vortices and magneto-rotational instabilities, and the photo-evaporation of gas. Transitional discs with large dust cavities (e.g. Espaillat et al. 2014)were found to represent almost 30% of the millimetre(mm)-bright discs in Taurus and Ophiuchus regions (Andrews et al. 2011). These systems show a lack of continuum mm emission with dust depletion factors in the rangeof 100–1000 in the disc central regions, while a variety of gas depletion factors have been reported within thesecavities (e.g. van der Marel et al. 2016, 2018). They may witness gravitational interaction between gas and sub-stellar/planetary companion(s) in the disc cavity, and drag forces between dust grains and gas may cause the mm grains to be trapped by pressure bumps at large radial distances (e.g. Zhu et al. 2011; Pinilla et al. 2012).

While the timescale for disc dissipation has been constrained to about 5 Myr from the fading of infrared (IR)-excess (for dust) and of accretion tracers or molecular lines (for gas) in young clusters (Haisch et al. 2001; Fedele et al. 2010), very few systems are known in the transition phase between protoplanetary and debris discs. A new category of discs has emerged in recent years, which simultaneously display the typical characteristics of debris discs (with in particular a weak fractional excess of IR emission fd =LIRL < 10−2 originating from the optically thin dust component of second-generation grains), and the presence of gas tracers (mostly CO lines) that reveal a non-negligible residual gaseous component (Zuckerman et al. 1995; Moór et al. 2011). Although only a handful of such gas-rich debris discs were known when we started the present project (namely HD 21997, 49 Ceti, β Pictoris), this population has now grown to almost twenty members with typical ages in the range of 10–40 Myr (e.g. Lieman-Sifry et al. 2016; Moór et al. 2017). The origin of the remnant gas is still a subject of controversy: the age of the system must be compared to the lifetime of the detected molecular species, which itself depends on the UV/ X-ray stellar/interstellar emission and on the (complex) shielding effects (Kral et al. 2017). Those systems harbouring supposedly primordial gas are called hybrid discs (Kóspál et al. 2013). The gas-rich debris discs revealed by CO lines simultaneously display a weak dust mm continuum and an IR excess characteristic of more evolved, optically thin debris discs. These systems are thought to witness a short-lived transient phase at the interface between proto-planetary and debris discs, when the material dissipates through photo-evaporation and accretion. In Péricaud et al. (2017), we reported a correlation between the CO flux density and the mm continuum emission for a broad variety of systems including T Tauri, Herbig Ae, transition discs, and even the proto-typical young debris disc β Pic. The position of hybrid discs in the diagram, systematically above the correlation line, suggests a faster evolution of the dust in these systems, leading to a transient and unusual increase of the gas-to-dust ratio.

In this population, the Herbig B9.5V/A0Ve star HD 141569A holds a very peculiar and interesting position. The star is located at 110.6 ± 0.5 pc (Gaia Collaboration 2018), and has two M-stars companions whose gravitational link remains unclear (Reche et al. 2009). With an age of about 5 ± 3 Myr (Weinberger et al. 2000; Merín et al. 2004), it is the youngest debris disc in this population with an IR brightness twice as large as that of β Pic. With the new Gaia distance, the stellar luminosity determined in Merín et al. (2004) turns to 27.0 ± 3.6 L. In IR colour-magnitude diagrams, HD 141569A lies at the exact interface between proto-planetary and debris discs (see, e.g. Wyatt et al. 2015, and Fig. 1 therein). Its cold molecular content was discovered through its CO emission line in the spectroscopic survey of Zuckerman et al. (1995), whereas warmer CO gas was later detected in the inner system by Goto et al. (2006). The debris disc of HD 141569A (LIRL = 8.4 × 10−3, Sylvester et al. 1996) revealed by optical/IR dust emission presents a complex multiple-ring architecture with spirals, arcs, and gaps imaged in scattered light by the HST (Augereau et al. 1999; Weinberger et al. 1999; Mouillet et al. 2001; Boccaletti et al. 2003; Clampin et al. 2003; Konishi et al. 2016). More recently, a series of concentric rings was discovered with VLT/SPHERE with a range of radii of 40–90 au (Perrot et al. 2016), partially corroborating the emission detected around 40 au in the L′ band by Currie et al. (2016) with Keck/NIRC2.

While the small-grain component has been well studied in the near-infrared (NIR), a thorough analysis of the larger grains observed at mm wavelengths and a comparative physical and morphological analysis of the CO gas distribution is crucial to understanding its evolution status as well as the nature of the observed dust and gas components. Our group obtained observations of HD 141569A in 1999–2000 with the Plateau de Bure Interferometer (PdBI) of Institut de Radio-Astronomie Millimétrique (IRAM), where the gas disc was already partly resolved (briefly reported in Dutrey et al. 2004). These early observations have motivated the present study, but their quality and sensitivity is outmatched by the current dataset and will therefore not be presented here. Observations of this system by the Submilliter Array (SMA), Combined Array for Research in Millimeter-wave Astronomy (CARMA), and the Atacama Large Millimeter/submillimeter Array (ALMA) have been reported, revealing an extended gas disc (from the 12CO J = 3 → 2 transition) out to the 250 au main dust ring highlighted in early HST images, and a dust disc extending out to 56 AU from the 870 μm continuum observations (White et al. 2016; Flaherty et al. 2016). The total mass of gas was revised to values much larger than initially derived thanks to the optically thin line 13CO J = 2 → 1 (Miley et al. 2018). From the mm continuum emission, a double component with spectral indices consistent with the range of values observed in both debris and primordial discs was reported by White & Boley (2018), and an unresolved dust ring peaking at 220 au was detected from 1.3 mm ALMA observations by Miley et al. (2018). In this article, we present complementary, multi-frequency observations with ALMA and the NOrthern Extended Millimeter Array (NOEMA)1 (Sect. 2). We characterise the dust mm continuum and the gas emission, taking advantage of the spectral information to derive constraints on the grain growth and emission lines opacities (Sect. 3). We finally model the 12CO and 13CO lines with the radiative transfer code Diskfit (Piétu et al. 2007) to characterise the temperature and density laws in this disc (Sect. 4). In Sect. 5, we finally compare the derived morphology at mm wavelengths with the dust disc seen in optical/IR domain and use the new constraints on the gas and dust mass to discuss the origin of the gas and the evolution status of this peculiar system.

2 Observations

2.1 PdBI/NOEMA observations

HD 141569 was observed with NOEMA during three runs from December 2014 to April 2016. The 12CO J = 2 → 1 transition at 230.538 GHz was mapped with the 7D and 7C configurations (7 antennas) of the IRAM array, with an angular resolution of 2.4′′ × 1.2′′ (baseline lengths from 15 to 192 m; all subsequent angular resolution values correspond to natural weighting). The spectral setup used a narrow band (20 MHz bandwidth, resolution 0.20 km s−1). Excellent conditions are noteworthy for one track, with precipitable water vapour inferior to 1.5 mm, and a 15° mean phase rms.

In winter 2014/2015, we also obtained 13CO J =2→1 (220.399 GHz) observations. The 6C and 7D configuration antennas were used with a total on-source time of 12.7 h. We reached a 2.3′′ × 1.2′′ angular resolution (baseline lengths from 15 to 176 m) and used the same spectral resolution as for 12CO J = 2 → 1. For both transitions, 1546+027 and 1508-055 were used for amplitude and phase calibration and we bootstrapped their fluxes using MWC 349 observations (with a flux model of 1.86 Jy at 219 GHz and 1.91 Jy at 230.5 GHz).

The WIDEX correlator provided 4 GHz bandwidth during the observations of the 12CO and 13CO. Combining the two observations at 1.3 mm, we obtained a 2.1′′ × 1.0′′ beam for the continuum map.

2.2 ALMA observations

We obtained ALMA observations of HD 141569 in Band 3 (84− 116 GHz). The data were recorded in August 2015 for the 13CO J = 1 → 0 (110.201 GHz) and C 18 O J = 1 → 0 (115.271 GHz) transitions during two runs with 39 and 44 antennas, respectively (project 2013.1.00883, PI. Péricaud). The spatial resolution was 0.76′′ × 0.56′′ (baselines up to 1.55 km), and the spectral resolution 0.2 km s−1. The data were collected with relatively poor weather conditions (PWV ~ 4 mm). We used Ceres and 1550+054 for the flux and amplitude calibration, 1550+0527 and 1517-2422 for bandpass calibration and 1550+0527 for the phase calibration.

In addition, we collected archival raw data from the project 2012.1.00698.S (PI. Boley) to complement our analysis in Band 7 (275−373 GHz), which includes the 12CO J = 3 → 2 line at 345.796 GHz. A preliminary analysis of this data set was reported in White et al. (2016). Only a subset of the observationscould be properly calibrated. The observations carried out with 32 antennas provide a resolution of 0.43′′ × 0.35′′ (baselines up to 650 m), and a spectral resolution of 0.8 km s−1. The flux and amplitude calibrator is Titan, the phase and bandpass calibrator is 1550+0527. We used the CASA package to calibrate the ALMA data, and we carried out the imaging and analysis of all NOEMA and ALMA observationswith the GILDAS software. The imaging was performed with natural weighting for both ALMA and NOEMA.

2.3 A deep search for extra molecular lines with IRAM-30 m/EMIR

We also used the EMIR detector on the IRAM 30 m antenna to investigate the molecular content of this hybrid disc (Project 049-15). The observations were performed in Summer 2015 with a total time of 20 h on source. Two frequency bands were used in this project: 87.720–93.500 GHz and 140.520–148.300 GHz. We used Wobbler switching, resulting in very flat baselines over the observed bandwidths. The frequency coverage included lines of HCN J = 1−0, HCO+ J = 1−0, and CS J = 3−2 that are usually bright in discs around classical T Tauri (of order 0.5–1.0 Jy km s−1, see Dutrey et al. 1997) but slightly fainter in Herbig Ae stars (e.g. MWC 480 and AB Aur, Piétu et al. 2007; Schreyer et al. 2008). No emission was detected. Assuming a similar line width to that of the CO lines (7 km s−1 FWHM), the 3σ detection threshold on the integrated line flux is 45 mJy km.s−1 around 90 GHz, and 60 mJy km.s−1 around 145 GHz.

3 Results

3.1 Continuum emission

The continuum emission is resolved at the three frequencies, although with some differences due to the different instrumental resolutions. The images presented in Fig.1 were produced using the CLEAN deconvolution algorithm with naturalweighting. The emission appears compact, and centrally peaked, with a possible secondary component which looks more extended in the NOEMA map only (a low level plateau in the 1.3 mm image). The coarser resolution of the NOEMA map (231 × 110 au) is not the onlyreason: the visibility profiles presented in Fig. 1 clearly show a steep increase at the shortest baselines. This visibility rise at u, v distance shorter than 50–100 m (present at the three frequencies) is characteristic of extended emission (at few arcsec scales), while the more compact, inner disc component is resolved at intermediate baselines (e.g. B ~ 150−400 m at 0.87 mm). The spatial extent and flux values of the continuum emission were determined by fitting Gaussian functions in the u, v plane. We tested a single and a two-component model. In Table 1 we present the best-fit parameters for the two models, and the measured maximum flux density corresponding to the shortest baseline points (so called zero-baseline flux). The single-component model is dominated by the compact disc component (0.3–0.5′′) at 0.87 and 2.8 mm, while at 1.3 mm only the broad component (2–3′′) is constrained. Therefore, the flux values for this model cannot be used together to estimate the dust spectral index. Our maximum flux values from the shortest baselines are consistent with those reported by single-dish observations or compact arrays (see Table 2). At 1.3 mm, we measure F230 GHz(u, v ≃ 0) = 4.00 ± 0.25 mJy, in agreement with the JCMT/SCUBA measurement (5.4 ± 1.0 mJy, Sylvester et al. 2001), but our value is about three times larger than the ALMA measurement of 1.7 mJy reported in Miley et al. (2018). Typical calibration uncertainties would amount to ~ 10% for both instruments, but cannot justify this discrepancy. At 2.8 mm (and 0.87 mm resp.) our measurements are consistent with the CARMA (resp. SMA) values F110 GHz = 0.78 ± 0.30 mJy (resp. F230 GHz = 8.2 ± 2.4 mJy, Flaherty et al. 2016). Because the inner component is very compact, the disc orientation and aspect ratio are only roughly reproduced by the elliptical Gaussian fit in the single component model (see Table 1). In the more realistic double component model below, we therefore first deprojected the visibilities assuming the geometric parameters derived at optical wavelength (PA = 356°, i = 53°, Augereau et al. 1999), and then fit in the u, v plane with a circular Gaussian function.

The second model suggests that the emission is about equally shared between a compact component, which is concentrated within about 50 au (FWHMint ~ 30−35 au, a value comparable to the interferometric beams at 0.87 and 2.8 mm), and a much broader component that extends out to about 350 au. The u, v plane coverage does not allow precise determination of the outer radius, and most of the flux of this extended component is filtered by the interferometer at 0.87 and 2.8 mm. It is marginally detected as a low-level plateau in the 1.3 mm map thanks to the larger beam (although its shape and radial extent cannot be fully determined in our data). From higher-resolution ALMA data (0.65″ beam) at 1.3 mm, Miley et al. (2018) revealed a low-level ring-like emission peaking at 220 ± 10 au that cannot be seen in our NOEMA map. This unresolved ring is different from the extended component that we report from our analysis of the u, v plane. In our higher-resolution maps at 0.87 and 2.8 mm, no such ring-like emission is detected. Assuming a spectral index αmm = 2 (see below), and extrapolating the 220 au ring peak brightness and our brightness sensitivity at 0.87 mm (0.4″ beam) and 2.8 mm (0.66″ beam), it would be marginally detectable at the ~3 and 5 σ level, respectively. Such a two-component model was also proposed by White & Boley (2018) from u, v plane fitting, but our results disagree with their total flux values (e.g. F345 GHz = 16.8 ± 1.7 mJy), which are systematically larger than both the zero-baseline flux values, and the earlier SMA, APEX, CARMA or JCMT observations(Flaherty et al. 2016; Nilsson et al. 2010; Sylvester et al. 2001). The observed differences between the flux values summarized in Table 2 can be explained firstly by the assumed shape of the model components (Gaussian, uniform discs, rings, power-laws), and the number of emission components taken into account. Secondly, the very-short-baseline information is not always taken into account in the data analysis (e.g. White et al. 2016; Miley et al. 2018). Thirdly, the u, v coverage and the field of view can vary significantly between instruments, and the signal-to-noise ratio as a function of (u, v) radius can be very different from one dataset to another (e.g. at 1.3 mm).

We estimate the dust spectral index β from the slope of the spectral flux density, and we first add quadratically a typical calibration uncertainty of 10% to the flux density values. The spectral index is defined as β = α − 2, with , Fν. A fit to the three frequencies yields the slope α. We derive the spectral index β in two cases: (i) for the global disc using the maximum flux values measured at the shortest baselines (u, v) ~ 0, (ii) for the double component model, which provides a more realistic description of the continuum emission than the single-component model. The values reported in Table 1 suggest that the spectral index is significantly smaller than the canonical ISM value with β ≤ 0.7 at the 3σ level. Whether this is indicative of grain growth or optically thick emission can only be determined with more resolved observations. In the disc modelling by Thi et al. (2014), the dust component is optically thin from optical up to mm wavelength both in the radial and vertical directions, and we therefore expect the constraints on the spectral index to primarily reflect a change in grain properties. Radial variations of this spectral index may also be expected, as already demonstrated in a sample of TTS and HAe stars (Guilloteau et al. 2011; Pérez et al. 2012). The double-component model suggests a larger value of β for the more extended disc component, although the difference with the inner 35 au region is not statistically significant. A similar trend has been claimed by White & Boley (2018), although their spectral index value for the outer component β = 0.28 ± 0.18 may be biased by an overly large flux value at 345 GHz, as explained above. A dense u, v coverage at both short and very long baselines would be required to firmly establish a change of dust properties between the inner compact and the extended dust disc components.

Table 1

Continuum results: best-fit parameters foro: (i) a single component model (elliptical Gaussian in the u, v plane), (ii) and a double-component model (two circular Gaussian functions after deprojection in the u, v plane according to the orientation and inclination determined at optical wavelength).

thumbnail Fig. 1

Left column: continuum emission maps, all contours are set with 3σ spacing, panel a: ALMA data at 0.87 mm, σ = 0.051 mJy beam−1, contours; (c) NOEMA data at 1.3 mm, σ = 0.078 mJy beam−1; (e) ALMA data at 2.8 mm, σ = 0.008 mJy beam−1. We note that the scale of panel c is changed to include the more extended emission. Right column: real part of the visibility (with spatial binning) as a function of the u, v radius in the Fourier plane, panels b, d, and f: correspond to data in panels a, c, and e, respectively.

Open with DEXTER
Table 2

Continuum flux comparison.

3.2 CO gas emission

The 12CO and 13CO J =   2→1 transitions are detected with NOEMA at a high signal-to-noise ratio (SN ≳ 150, and SN ≳ 20 respectively), whereas the 13CO J = 1 → 0 detection with ALMA is more marginally significant (SN ~ 5). Integrated intensities for all CO transitions are presented in Table 3 with their associated spatial resolution. The intensity we derive for the ALMA 12CO J = 3 → 2 observations 18.4± 0.3 Jy km s−1 is slightly larger than (but consistent within 2σ with) the value reported by White et al. (2016) and by Flaherty et al. (2016) with the SMA. The intensity maps (zeroth moment) are displayed in Fig. 2 for the three transitions (natural weighting). The 12CO emission appears slightly more extended in the NOEMA map because of the convolution by a larger beam. The 13CO J = 2 → 1 emission appears to be more compact than the 12CO J = 2 → 1 transition which extends beyond 2, likely as a result of a difference in opacity since the sensitivity and resolution are equivalent for these two lines. We also display the ALMA 12CO J = 3 → 2 emission for comparison, using uniform weighting in order to show evidence for the tiny inner gas cavity which is readily resolved in this high-resolution image. Given its very low level, the continuum emission was not subtracted in Fig. 2, meaning that this CO cavity cannot artificially result from over-subtraction of the continuum. The separation of the central emission peaks is about 0.8′′, which translates into a maximum cavity radius of about 45 au. This value is refined in the subsequent modelling presented in the following sections. The channel maps in Figs. A.1 and A.2 display the typical velocity pattern of rotating discs. In Fig. 3, the CO integrated area maps have been superimposed to the scattered light emission from the HST observations. These figures confirm that most of the molecular emission is apparently less extended than the optical dust disc. The 13CO J = 1 → 0 transition has a rather poor sensitivity due to bad weather conditions and a large phase noise (Fig. A.3), we only use its integrated intensity in the subsequent analysis. The two transitions C 18 O J = 2 → 1 (219.560 GHz) and C 18 O J = 1 → 0 (109.782 GHz) are not detected in our NOEMA and ALMA projects respectively, but constraining upper limits are given in Table 3.

Table 3

Summary of CO lines detections and setups.

thumbnail Fig. 2

Maps of integrated CO emission for the four detected transitions (the continuum was not subtracteddue to its very low level of emission). Top: NOEMA observations, from left to right: 12CO J = 2 → 1 (first contour at 5σ, second contour at 15σ, then 15σ spacing withσ = 31 mJy beam−1 km s−1), and 13CO J = 2 → 1 with 3σ spacing contours, σ = 23 mJy beam−1 km s−1. Dotted lines show negative contours in steps of 3σ. Bottom:ALMA observations, from left to right: 12CO J = 3 → 2 with uniform weighting to highlight the central gas cavity (3σ spacing, σ = 40 mJy beam−1 km s−1), and 13CO J = 1 → 0 1σ spacing contours starting at 2σ (for positive and negative contours), σ = 6.8 mJy beam−1 km s−1.

Open with DEXTER
thumbnail Fig. 3

Integrated intensity of the CO line superimposed on the HST scattered emission (Clampin et al. 2003). The cross indicates the position angle and aspect ratio as determined from gas modelling (Sect. 3). Left: 12CO J = 2 → 1 emission, the contour spacing is 6σ, i.e. 5.5 × 10−1 Jy beam−1 km s−1. Beam size: 2.48 × 1.45′′. Right: 13CO J = 2 → 1 emission, with a 3σ contour spacing, i.e. 7.8 × 10−2 Jy beam−1 km s−1. Beam size: 1.76 × 1.32′′.

Open with DEXTER

3.3 Opacity of CO lines

Using the line ratio of the 12CO and 13CO isotopologues for the same transition level, the line opacity can be estimated with the following formula: (1)

where τ12 and τ13 are the opacity of the 12CO and 13CO lines, respectively. If we further assume that the opacity ratio of 12CO over 13CO is equal to the elemental abundance ratio (Wilson & Rood 1994), then τ12 /τ13 = 77 and we can solve for Eq. (1). For the J = 2 → 1 transition, we measure from NOEMA observations a line ratio , and derive τ12 = 14 ± 4 and τ13 = 0.18 ± 0.05. It should be noted that here we derived the CO opacity from the 12CO /13CO flux ratio in velocity channels of 1 km s−1 to minimise the smearing effect due to rotation (Guilloteau et al. 2006) and averaged the results. The same analysis for the J = 1 → 0 transition, based on our new ALMA-Band 3 measurement and the integrated intensity Jy km s−1 reported by Flaherty et al. (2016) with CARMA yields τ12 = 8 ± 2 and τ13 = 0.11 ± 0.02, in good agreement with the J = 2 → 1 transition estimate. Even if these values could be biased by a factor of a few due to the uncertainties on the 12 C∕13C abundance ratio, as discussed in Kóspál et al. (2013), the 12CO lines appears to be optically thick and the 13CO optically thin. The 13CO transition is therefore a more robust gas mass tracer than 12CO, and this consideration leads us to reconsider the mass of gas derived by Flaherty et al. (2016) and White et al. (2016) who assumed optically thin emission for the 12CO lines. Within the uncertainties, our upper limit on C18O at 115.271 GHz and 219.560 GHz is consistent with ISM CO isotopologue abundance ratios.

4 Modelling

We independently model the 12CO J = 2 → 1, 13CO J = 2 → 1, and 12CO J = 3 → 2 lines with the Diskfit code (Piétu et al. 2007). We assume power laws for the CO surface density (), the temperature , and the gas velocity . In this study, we fixed the velocity exponent so that av = 0.5 (Keplerian rotation). The disc is assumed to have a sharp outer edge at Rout. The vertical density is assumed to have a Gaussian profile (see Eq. (1), Piétu et al. 2007), with the scale height being a (free) power law of radius . The line emission is computed assuming a (total) local line width dV independent of the radius and local thermodynamic equilibrium (LTE; i.e. T0 represents the rotation temperature of the level population distribution). In addition to the above intrinsic parameters, the model also includes geometric parameters: the source position (x0; y0), the inclination i, and the position angle (PA) of the rotation axis; and the source systemic velocity Vsys relative to the local standard of rest (LSR) frame.

The modelling is done in the u, v plane to avoid non-linear effects due to deconvolution. The minimization is completed by a Bayesian analysis using a Markov chain Monte Carlo (MCMC) code, emcee (Foreman-Mackey et al. 2013). Our approach differs from the analysis of White et al. (2016), which did not include the physical disc parameters in the modelling and focussed on the geometrical and kinematic description of the 12CO J = 3 → 2 emission. Our results are presented in Table 4. Comparison with observations are shown for the integrated spectra in Fig. 4, and for the channel maps in Fig. 5.

4.1 Spatial distribution and kinematics of the CO gas

The spatial extent of the gas disc is constrained to consistent values around Rout = 275 ± 3 au by the two transitions of 12CO. This value is slightly larger (within 3σ) than the 224 ± 20 au reported byFlaherty et al. (2016) from a simultaneous fit to 12CO J = 3 → 2 and 12CO J = 1 → 0, because these latter authors fixed the density exponent to a smaller value: p = 1.5. The 13CO line emission is slightly more compact with Rout = 240 ± 20 au. Although the discrepancy does not appear to be statistically significant, it may result from the combination of opacity and sensitivity effects. More interestingly, an inner cavity is confirmed with a radius Rin = 17 ± 3 au from 12CO J = 3 → 2, as first suggested in lower-resolution maps by Flaherty et al. (2016). The spatial resolution and sensitivity of the NOEMA observations are not sufficient to provide a robust constraint, but point towards consistent values of the inner disc edge. The 13CO J = 2 → 1 data suggest a larger value Rin = 35 ± 5 au. This discrepancy should be taken with care because the beam size of these NOEMA data is more than four to five times larger than the ALMA beam for 12CO J = 3 → 2. If confirmed, this result could be explained by the difference in opacity between the two isotopologues: since the 12CO transition is 70 times more optically thick than the 13CO transition, a larger inner radius for 13CO could be indicative of a shallow increase of the density at the inner disc edge, while a sharp edge is expected to produce similar inner radii for both lines. A similar effect has been reported in the 15–70 au gas cavity of transitional disc J160421.7-213028 (Dong et al. 2017), although the nature of the two discs is different, because HD 141569A does not show any sign of having a dust cavity (a compact continuum emission was reported at all frequencies in Sect.3.1). At the reported resolution, the presence of unresolved gaps in the gas distribution in the central 50 au may also result in a similar discrepancy between the CO isotopologues. The inner rim value (17 au) for 12CO is coherent with the spectroscopic detection of warm CO through IR ro-vibrational lines by van der Plas et al. (2015), with a modelled distribution of the gas between 14 and 45 au (after correcting for the revised stellar distance).

The inclination and position angle are found to be in the range i = 56−58° and PA = 86 ± 1° for 12CO lines, and i = 53 ± 2° and PA = 90 ± 2° for 13CO, which is in good agreement with the optical values of i = 52.5 ± 4.5° and PA = 85.4 ± 1° respectively (Augereau et al. 1999). The disc is flared, although the flaring exponent h is not well constrained, a value consistent with h = 1.25 is found for both 12CO transitions, while it is naturally not constrained for the optically thin transition of 13CO. For 12CO, a scale height H0 = 17 au is found at 100 au. This is twice the value of the hydrostatic scale height (Hhydro = 8 au) for a disc with T0 = 30 K, as measured from the optically thick transitions of 12CO. This is consistent with direct measurements of the 12CO layer location in HD 163296 (de Gregorio-Monsalvo et al. 2013) and the Flying Saucer (Guilloteau et al. 2016), where CO is located at 2−3 scale heights (Dutrey et al. 2017).

In our preliminary fits, the velocity exponent proved to be very close to the Keplerian value av = 0.5, and we adopted this fixed value for the MCMC exploration reported in Table 4. At 100 au, the rotation velocity is better constrained for 12CO lines, with a value V0 = 4.44 ± 0.01 km s−1 that can be converted into a central stellar mass of 2.22 ± 0.01 M. This value is slightly smaller than the mass of 2.39 ± 0.05 M reported by White et al. (2016), but is very consistent with the spectral analysis of Merín et al. (2004) that yielded . The 13CO analysis leads to a slightly larger mass M = 2.35 ± 0.05 M.

thumbnail Fig. 4

Observed line profiles and superposed best models obtained with Diskfit (colour dashed lines) for, from left to right: 12CO J = 2 → 1, 13CO J = 2 → 1 NOEMA data, and 12CO J = 3 → 2 ALMA data.

Open with DEXTER
Table 4

Results of Diskfit modelling for the 12CO and 13CO with 1σ error bars.

4.2 Temperature and density laws

The modelling of the two optically thick transitions 12CO J = 2 → 1 and 12CO J = 3 → 2 provides very consistent estimates for the excitation temperature, T0 ~ 30 K with a relatively flat radial profile q < 0.4 (q = 0.28 ± 0.01 for 12CO J = 3 → 2 and q = 0.10 ± 0.09 for 12CO J = 2 → 1). The temperature is ill constrained for the optically thin 13CO transition. A value of 30 K is consistent with the analysis of the 12CO J = 3 → 2 detection with SMA by Flaherty et al. (2016, who used a fixed temperature exponent q = 0.5 and found K). It is somewhat smaller than the gas temperature of about 50 K derived from 12CO line for a similar, albeit younger and more optically thick disc around the Herbig Ae star MWC 480 (Piétu et al. 2007). Given the star luminosity of 30L, blackbody grains in an optically thin disc are predicted to have a temperature of 60 K at a distance of 100 au, and would be much hotter than the detected gas. The derived gas temperature is also significantly smaller than the range of kinetic temperatures computed by Thi et al. (2014) with their Prodimo modelling based on unresolved detections of [OI], C[II] and 12CO J = 3 → 2 lines around HD 141569. In their modelling, the gas is heated by Polycyclic Aromatic Hydrocarbons (PAH) and cooled by [OI] and CO lines, and the computed temperature reaches about 100 K at 100 au in the inner disc and drops down to 40 K at 500 au (i.e. well beyond the outer edge of the emission we detect with ALMA/NOEMA). The excitation temperature we derive from 12CO may still be smaller than the actual gas kinetic temperature. Low values of the gas temperature (although with a much more dramatic deviation) have also been reported in the hybrid disc HD 21997 by Kóspál et al. (2013), where the excitation temperature derived from CO line ratios drops to values as low as 6− 9 K. Flaherty et al. (2016) propose that this low gas temperature could be explained by enhanced cooling from CO (following Hollenbach & Tielens 1999), and/or NLTE effects (following Matrà et al. 2015).

In the absence of C18O emission, we use the more optically thin emission of the 13CO J = 2 → 1 line to constrain the CO surface density. Our modelling yields a best-fit value cm−2 at 100 au with an exponent p = 1.8 ± 0.5. We note that the two other transitions of 12CO lead to much steeper profiles with p ~ 3 and require a surface density of the order of 6 × 1016 cm−2. These values of the surface density are at least one order of magnitude smaller than those found for the younger, gas-rich disc around the Herbig Ae star MWC 480: Piétu et al. (2007) report, at a distance of 100 au, Σ0 = 8 × 1016 cm−2 for 13CO and 6 × 1018 cm−2 for 12CO. This system is taken here as a reference as it is one of the very few where the surface density of 13CO has been precisely constrained through similar disc emission modelling.

We derived a total mass of CO gas based on Σ0 and used the minimum outer disc edge value Rout = 232 au inferred from the 13CO line (hence a minimum mass value). Because most of the mass lies in the inner regions with a power law distribution, the computed value is sensitive to the inner disc edge, and we take into account the inner cavity detected with the 12CO emission in the ALMA high-resolution maps (Rin = 17 au). Assuming a standard isotopic ratio 12CO/13CO = 77, these parameters yield M (= 4 × 1023 kg), where we take into account the uncertainties on Rin, Rout, Σ0 and p. This value is 50 times larger than the CO mass reported by White et al. (2016) based on the 12CO emission.

The conversion to a total mass of gas is very uncertain. A low CO content compared to the dust emission was already noted in the early detection of CO isotopologues in DM Tau (Dutrey et al. 1997). Several recent studies (Bergin et al. 2013; Reboussin et al. 2015) favour much lower values for the [CO/H2] ratio than the canonical ISM value of 10−4 (see also thediscussion on carbon depletion in e.g. Miotello et al. 2017). We discuss this point again in Sect. 5.2. Here, in order to compare with other works, we nevertheless convert this minimum mass of CO into M or MJup ≃ 2.10−4 M using the standard ISM abundance. Using the grid of models developed by Miotello et al. (2016), which include the effect of isotope selective photodissociation, the luminosity of the 13CO J = 2 → 1 line and the non-detection of C 18 O J = 2 → 1 favours a total mass in the range 1−5 × 10−4 M, in good agreement with the previous estimate. This is one order of magnitude larger than the value reported by Flaherty et al. (2016) from the SMA detection of 12CO J = 3 → 2, but is in good agreement with the earlier modelling of Thi et al. (2014) and Jonkheid et al. (2006), based on spatially unresolved observations of the same line, who found a total gas mass in the range 67− 164M. Our value is a factor of three to five smaller than the total gas mass of 6− 9 × 10−4 M reported by Miley et al. (2018) based on the integrated flux of the optically thin transition 13CO J = 2 → 1, but these remain in reasonable agreement due to the different methodologies used. This total gas mass value is clearly on the lower bound of the range of disc masses found for Herbig Ae stars, and usually derived from mm dust flux converted to gas mass with an assumed gas-to-dust ratio of 100 (see, e.g. Williams & Cieza 2011). It is instead typical of discs around very low-mass TTS (for instance, TTS disc masses obtained from 13CO and C 18 O J = 3 → 2 in Lupus range between 10−5 and 10−3 M, based on models assuming isotope selective photo-dissociation and CO freeze-out on grains Miotello et al. 2017). HD 141569 also harbours one of the most massive molecular gas components among the gas-rich debris discs reported in Moór et al. (2017) based on supposedly optically thin CO isotopologues. Even if we underestimate the CO mass due to 13CO optical depth effects or the uncertain [CO/H2] abundance, such a low mass supports the process of ongoing gas dissipation in this 5 Myr-old disc.

thumbnail Fig. 5

Results of the modelling showing the channel maps for the transition (from left to right) 12CO J = 2 → 1, 13CO J = 2 → 1 (NOEMA) and 12CO J = 3 → 2 (ALMA). In each sub-panel, we present from left to right: the observations, the model (see best-fit parameters in Table 4), and the residual emission. The latter is obtained by subtraction in the u, v plane before imaging and cleaning. Contour spacing is 5σ for the 12CO transitions (σ = 16 mJy beam−1 for J = 2 → 1 and σ = 6.3 mJy beam−1 for J = 3 → 2), and 3σ for the 13CO J = 2 → 1 transition (σ = 10 mJy beam−1). Residual contour spacing is 3σ for all transitions.

Open with DEXTER
thumbnail Fig. 6

Global scheme of dust and gas distribution in HD 141569 summarizing the most recent resolved observations.

Open with DEXTER

5 Discussion

5.1 Comparison with optical/NIR emission

The distribution of matter is strikingly different at optical/NIR and mm wavelengths. A global sketch summarizing and comparing the respective radial extent of the main disc components detected at optical/NIR and mm wavelengths is presented in Fig. 6. The continuum emission at mm wavelength is clearly much more compact than the 250–400 au broad rings of the debris disc revealed at optical/NIR wavelengths, and is also more compact and centrally peaked than the gas emission characterised in the following section. Furthermore, it is coincident with the innermost component recently reported by various authors in the NIR domain. Indeed, in the inner 100 au recent observations with VLT/SPHERE have revealed the presence of a series of concentric rings and arcs with semi-major axes in the range 45− 90 au (Perrot et al. 2016). The link between the μm grains responsible for the IR emission and the mm grains detected with ALMA/NOEMA cannot be clarified because of the limited resolution of our data. VLT/VISIR observations in the mid-IR also suggest the presence of an even more internal component (Thi et al. 2014), which was recently detected at L′ band with the Keck vortex coronagraph within 70 au (Mawet et al. 2017) and may extend up to 100 au. According to the extended scattered-light emission reported with HST/STIS by Konishi et al. (2016), this inner warm component may also be spatially associated with the inner rim of the gas disc that our modelling locates around 15−20 au from 12CO J = 3 → 2 maps (similarly to the warm CO detected in spectroscopy in the 14− 60 au region by Goto et al. 2006; van der Plas et al. 2015). The disc is not classified as a transitional disc because of the lack of a mm dust cavity. A small cavity of no more than 15 au in radius is observed in the CO maps, consistent with the earlier spectroscopic detections of a warm CO ring, but the continuum emission is strikingly centrally peaked. Inside this gas cavity, there is probably not much material: VLTI-PIONIER observations in the H-band reported a completely unresolved structure (Lazareff et al. 2017), which might be dominated by the stellar emission itself (accounting for 86% of the total flux), without suggestion of an inner ring near the sublimation distance of small grains.

The scattered light images show many asymmetries at various spatial scales, with spirals, arcs, and multiple ring structures. The outer limit of the gas disc is precisely coincident with the so-called “inner ring” near 250 au reported previously. Among the possible ring formation mechanisms, Takeuchi & Artymowicz (2001) proposed a complex mechanism based on the outward migration of small grains at the outer edge of a gas disc. HD 141569 is the only resolved disc so far that displays this combination of distant dust rings and a gas-rich inner disc.

Large-scale asymmetries are also suggested in both wavelength domains. In Fig. 7, we superimpose the SPHERE/ Ks image on the map of our best-fit residual of 12CO J = 3 →2 ALMA emission. This residual emission, which was already reported by White et al. (2016), is significant (up to 6σ) at all velocity channels and peaks between 6 and 9 km s−1 (see Fig. 5). We present here the integrated emission, which extends on the western side with an elliptical shape. It displays two brightness peaks centred at 60–70 au of deprojected distance at azimuth angles of about 220° and 290°, in the region of the three concentric rings detected with SPHERE. Since the 12CO emission is optically thick and the feature is not detected in the isotopologues due to insufficient sensitivity and angular resolution, we cannot disentangle between a local temperature offset and a density feature. However, the significance of this 12CO brightness increase, the coincidence with the location of the inner rings, and the similar brightness asymmetry raise questions about the interplay between gas and dust in this region. Lyra & Kuchner (2013) proposed a mechanism based on the development of a clumping instability to create eccentric rings from the coupling of gas and grains, but this appears to be efficient in a regime of low gas-to-dust ratio, which does not seem to be favoured in the inner region (see discussion in Sect. 5.2). The 250 au dust ring (and possibly also the outermost 400 au ring) at the edge of the detectable gaseous component (see Fig. 3) may instead be an ideal location to reach the necessary very low value of gas-to-dust ratio (GD < 1) for such an instability to develop. So far, no planetary companion candidate has been reported in the inner 100 au region with a sensitivity limit of the order of 1−3 MJup (Perrot et al. 2016). The destruction of planetesimals (by collision or evaporation) could produce both large quantities of small grains and release substantial amounts of gas, as has been suggested by the detection of an asymmetrical emission of CO in β Pictoris (Dent et al. 2014). In HD 141569, such a collisional event or enhanced planetesimal evaporation might be more difficult to demonstrate because of the large residual amount of primordial gas in the region of the NIR rings.

thumbnail Fig. 7

Montage of the 12CO J = 3 → 2 residual emission with the series of concentric rings detected in NIR scattered light with SPHERE by Perrot et al. (2016).

Open with DEXTER

5.2 Gas-to-dust ratio

The global gas-to-dust mass ratio in the disc is expected to trace the evolution of the system, but getting an accurate value is challenging because of the numerous assumptions that must be made especially in the computation of the total gas mass, as discussed in the previous section. The mass of dust is no less uncertain as it relies on the unknown value of the dust mass opacity κν, and we will follow the prescription: cm2 g−1, which is similar to that adopted by Beckwith et al. (1990) (with β =1). Mdust can be estimated on the basis of the mm dust flux (Fν), which is constrained in our observations at three different frequencies. From the modelling of the SED, earlier models by Jonkheid et al. (2006) and Thi et al. (2014) showed that the dust emission is optically thin from opticalto mm wavelengths, meaning that the mm flux can be used to provide a robust estimate of the dust mass. Using the mm dust fluxes reported in Table 1, the extreme values of β in the range [0,0.7], and typical values of the temperature for dust grains located at 50–100 au from this 30 L star (Td = 65−90 K), Mdust is computed with the standard equation: (2)

We obtain values for Mdust in the range 0.03−0.52M. This value is significantly smaller than for other Herbig Ae systems. In order to compare with dust masses derived for other gas-rich debris discs reported in Moór et al. (2017), we alternatively used the same dust opacity (κν=230 GHz = 2.3 cm2 g−1) and their assumed temperature (Td = 50 K), which leads to Md = 0.46M, which makes it the most massive gas-rich debris disc with HD 131488 in this study. With our gas mass estimate, 70 M, and our first choice of dust opacity, yielding Mdust ~ 0.03−0.52M, we can obtain a gas-to-dust mass ratio GD in the range 135−2370. This is most likely a severe lower limit, since we assumed a high CO/H2 abundance ratio to derive the gas mass. In T Tauri discs, it is in general necessary to invoke an overall depletion of CO by a factor of ten or so to bring the derived G/D ratio to the ISM value of 100 (see, Piétu et al. 2007; Williams & Best 2014), although the warmer discs around HAeBe stars have higher CO/dust ratios, and can sometimes be directly interpreted by GD = 100 and a normal CO/H2 ratio (e.g. AB Aur Piétu et al. 2007, or also HD 163296 in Williams & Best 2014). The estimated gas-to-dust mass ratio appears at least one order of magnitude larger than the values reported in Miotello et al. (2017) for a sample of sources in Lupus, where gas mass estimates were also derived from optically thin CO isotopologues (although with more complex models including isotope selective processes, as mentioned earlier).

This analysis must furthermore be refined because the morphology of the dust and gas components is quite different. The GD ratio can obviously display ample radial variations due to migration of large grains towards the inner disc regions, or local enhancements of the gas and dust density. The resolved images of the HD 141569 disc suggest that the bulk of the mm grains reside in a small disc confined within Rout,dust ~ 100 au, while the gas reservoir extends between about 20 and 270 au. There is possible evidence for a second, low-emission component of dust grains in the lower-resolution NOEMA maps (and concurrent hints at very short ALMA baselines), which could extend out to the limit of the gas disc and may contribute up to 30% of the dust emission. We checked however that the spatial variations of GD do not affect the overall conclusion. Similar large values of the GD ratio have also been reported in the 30 Myr-old hybrid disc HD 21997 by Kóspál et al. (2013): the authors quote a CO-to-dust ratio of 0.4−0.9 within Rout = 150 au that converts to 280−630 for H2-to-dust. However, the morphology of the two systems is very different, with a dust cavity within 55 au for HD 21997, but no evidence for a gas cavity. We also reported in Péricaud et al. (2017) the tendency of gas-rich debris discs to display a larger ratio of CO to mm continuum flux compared to other Herbig/TTauri systems, which could suggest that there is a larger gas-to-dust ratio in this transient phase. Beyond 270 au, the abrupt rarefaction of the gas naturally decreases GD, although neither mm grains nor gas emission are detected in this cold region. The outermost rings can be regarded as a nascent debris disc, where the grains detected in scattered light emission are produced by a collisional cascade from a reservoir of unseen parent bodies.

5.3 Origin of the gas

From the optically thin 13CO emission, we derived a 12CO gas mass in the range 8–15 ×10−2 M, a factor of 50 larger than the estimate by White et al. (2016) based on the optically thick 12CO transition, and a factor of 2 larger than the gas mass derived by Miley et al. (2018) based on optically thin 13CO observations. The gas disc appears also twice more massive than the C18O mass estimate in the 30 Myr-old hybrid disc HD 21997 (Kóspál et al. 2013). Based on the typical lifetime for self-shielded CO molecules and using estimates of CO content and gas production in solar system comets, these latter authors computed that more than 6000 Hale-Bopp-like comets would have to be destroyed each year to reproduce the observed level of CO in HD 21997 if the gas were of secondary origin. In the hypothesis of an exo-cometary origin, CO molecules can only be protected through self-shielding, which is however a more efficient mechanism for 12CO than for its isotopologues. We estimate the lifetime of self-shielded 13CO molecules from Visser et al. (2009) using the gas column density derived from our previous analysis and adopting Tex = 50 K, and we find 860 yr as a maximum value under the ambient interstellar UV field (the actual lifetime might even be shorter considering the additional contribution of the stellar UV field). Using the mass measured for the optically thin molecule 13CO and assuming a standard isotopic ratio, the 12CO production rate required to sustain the observed amount of 13CO is 4.7 × 1020 kg yr−1, and the 12CO production rate is 3 × 1021 kg yr−1, a value more than 2000 times larger than the gas production rate inferred for the β Pic disc by Dent et al. (2014) and Matrà et al. (2017). In Péricaud et al. (2017), we alternatively calculated the maximum amount of 12CO that can be locked in comets and planetesimals for a massive Herbig Ae disc (Md = 0.1M), assuming that 10% of the total amount of available oxygen in the disc can ultimately be evaporated in the form of CO (an upper limit derived from solar system comets). We found an upper limit for the initial 12CO mass in comets and planetesimals of about 4 × 1025 kg, hence a 13CO mass of 5.2 × 1023 kg, which is 100 times the CO mass value estimated for HD 141569. With a 13CO lifetime of about 860 yr, the observed mass of 13CO would imply a replenishment rate of 6 × 1018 kg yr−1, and hence a maximum lifetime of the initial 13CO reservoir of less than 105 yr. This value is likely an upper limit because we overestimated the CO content in comets relying on CO-rich comets onlyand assumed a very massive disc. It is however much shorter than the system age of 5 ± 3 Myr (Merín et al. 2004).

Moreover, the morphology of the CO emission, which continuously extends from 15–20 au out to 275 au favours a primordial origin for the gas, in contrast to the localised emission detected around β Pic where the gas and dust are likely produced in a planetesimal belt (Matrà et al. 2017). Our modelling of the CO maps shows that the gas disc is well represented by a protoplanetary-like disc with a surface density an order of magnitude lower than in the typical 1–2 Myr-old discs around Herbig Ae stars. The total amount of dust and gas in this system is similar to the older debris discs around HD 21997, HD 121617, HD 131488 and HD 131835 (Moór et al. 2017). No large inner cavity is found in HD 141569 in contrast to HD 21997 or HD 121617, indicating possible different mechanisms at play. A possible alternative scenario wasrecently proposed by Kral et al. (2017, 2019), in which even the most massive hybrid discs could be explained by secondary CO gas shielded by accumulated neutral carbon that is a photodissociation product of molecular gas released by planetesimals. This would require specific (low) values of the viscosity to accumulate a sufficient amount of C0, and it is uncertain whether this scenario could efficiently be applied to such a young system as HD 141569A.

The detection of other molecular species in the gas phase can provide a complementary diagnosis on the nature of the gas and reveal the specific chemistry in such a system. Our attempt to find HCO+, H2 CO, HCN, HC3N and CS with the IRAM-30 m radio antenna did not provide any successful detection (see Sect. 2.3). With our sensitivity limit, we would have been able to detect HCO+ J = 1 → 0 and CS J = 3 → 2 lines at >20σ if these were as bright as in the disc of AB Aur (Pacheco-Vázquez et al. 2015). However, the comparison should account for the smaller size of the HD 141569 disc, which is 250 au in radius, instead of about 500 au for that of AB Aur, and of the differences in distance. We therefore conservatively conclude that the average brightness for these transitions is at least a factor of three below that found for AB Aur. This may be the result of lower column densities or lower excitation temperatures. Both possibilities favour a low gas surface density in HD 141569. The presence of a prominent OI line (in addition to CII) revealed by Herschel also underlines the similarity with other Herbig Ae discs, because this feature is not detected in any of the more evolved, gas-rich debris discs (Meeus et al. 2012; Moór et al. 2016). The presence of prominent PAH features is another argument favouring the presence of primordial material in at least some regions of the disc, since such strong features have not been detected in IR spectra of Vega-like stars so far (e.g. Chen et al. 2006). This again reinforces its classification as an evolved disc around a Herbig Ae star rather than a debris disc. The disc is thus likely dominated by primordial gas in the phase of dissipation, and differs from its less evolved siblings due to its reduced CO surface density, and most notably its extremely low continuum which suggests a faster evolution of the dust component as already highlighted in Péricaud et al. (2017).

Acknowledgements

This work was supported by “Programme National de Physique Stellaire” (PNPS from INSU/CNRS.) This research made use of the SIMBAD database, operated at the CDS, Strasbourg, France. This paper makes use of the following ALMA data: ADS/JAO.ALMA2012.1.00698.S and ADS/JAO.ALMA.2013.1.00883.S. ALMA is a partnership of ESO (representing its member states), NSF (USA), and NINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea) in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ. This paper is also based on observations carried out with the IRAM NOEMA interferometer and the IRAM 30-m telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain). We thank the anonymous referee for his/her careful reading and useful suggestions to improve the manuscript.

Appendix A Channel maps for the 12CO and 13CO lines

thumbnail Fig. A.1

Channel map for the 12CO J = 2 → 1 transition observed with NOEMA, the spatial resolution is 2.5 × 1.4′′, the contour spacing is 7σ.

Open with DEXTER
thumbnail Fig. A.2

Channel map for the 13CO J = 2 → 1 transition observed with NOEMA, the spatial resolution is 1.9 × 1.1′′, the contour spacing is 2σ, same channel selection as for Fig. A.1.

Open with DEXTER
thumbnail Fig. A.3

Channel map for the 13CO J = 1 → 0 transition observed with ALMA, the spatial resolution is 0.7 × 0.5′′, the contour spacing is 3σ, same channel selection as for Fig. A.1.

Open with DEXTER

References

  1. Andrews, S. M., Wilner, D. J., Espaillat, C., et al. 2011, ApJ, 732, 42 [Google Scholar]
  2. Andrews, S. M., Huang, J., Pérez, L. M., et al. 2018, ApJ, 869, L41 [NASA ADS] [CrossRef] [Google Scholar]
  3. Augereau, J. C., Lagrange, A. M., Mouillet, D., & Ménard, F. 1999, A&A, 350, L51 [NASA ADS] [Google Scholar]
  4. Avenhaus, H., Quanz, S. P., Garufi, A., et al. 2018, ApJ, 863, 44 [NASA ADS] [CrossRef] [Google Scholar]
  5. Beckwith, S. V. W., Sargent, A. I., Chini, R. S., & Guesten, R. 1990, AJ, 99, 924 [NASA ADS] [CrossRef] [Google Scholar]
  6. Bergin, E. A., Cleeves, L. I., Gorti, U., et al. 2013, Nature, 493, 644 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
  7. Boccaletti, A., Augereau, J.-C., Marchis, F., & Hahn, J. 2003, ApJ, 585, 494 [NASA ADS] [CrossRef] [Google Scholar]
  8. Chen, C. H., Sargent, B. A., Bohac, C., et al. 2006, ApJS, 166, 351 [NASA ADS] [CrossRef] [Google Scholar]
  9. Clampin, M., Krist, J. E., Ardila, D. R., et al. 2003, AJ, 126, 385 [NASA ADS] [CrossRef] [Google Scholar]
  10. Currie, T., Grady, C. A., Cloutier, R., et al. 2016, ApJ, 819, L26 [NASA ADS] [CrossRef] [Google Scholar]
  11. de Gregorio-Monsalvo, I., Ménard, F., Dent, W., et al. 2013, A&A, 557, A133 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  12. Dent, W. R. F., Wyatt, M. C., Roberge, A., et al. 2014, Science, 343, 1490 [NASA ADS] [CrossRef] [Google Scholar]
  13. Dong, R., van der Marel, N., Hashimoto, J., et al. 2017, ApJ, 836, 201 [NASA ADS] [CrossRef] [Google Scholar]
  14. Dutrey, A., Guilloteau, S., & Guelin, M. 1997, A&A, 317, L55 [NASA ADS] [Google Scholar]
  15. Dutrey, A., Lecavelier Des Etangs, A., & Augereau, J. C. 2004, Comets II, eds. M. C. Festou, H. U. Keller, & H. A. Weaver (Tucson, AZ: University of Arizona Press), 81 [Google Scholar]
  16. Dutrey, A., Guilloteau, S., Piétu, V., et al. 2017, A&A, 607, A30 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  17. Espaillat, C., Muzerolle, J., Najita, J., et al. 2014, Protostars and Planets VI (Tucson, AZ: University of Arizona Press), 497 [Google Scholar]
  18. Fedele, D., van den Ancker, M. E., Henning, T., Jayawardhana, R., & Oliveira, J. M. 2010, A&A, 510, A72 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  19. Flaherty, K. M., Hughes, A. M., Andrews, S. M., et al. 2016, ApJ, 818, 97 [NASA ADS] [CrossRef] [Google Scholar]
  20. Foreman-Mackey, D., Hogg, D. W., Lang, D., & Goodman, J. 2013, PASP, 125, 306 [NASA ADS] [CrossRef] [Google Scholar]
  21. Gaia Collaboration (Brown, A. G. A., et al.) 2018, A&A, 616, A1 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  22. Garufi, A., Benisty, M., Pinilla, P., et al. 2018, A&A, 620, A94 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  23. Goto, M., Usuda, T., Dullemond, C. P., et al. 2006, ApJ, 652, 758 [NASA ADS] [CrossRef] [Google Scholar]
  24. Guilloteau, S., Piétu, V., Dutrey, A., & Guélin, M. 2006, A&A, 448, L5 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  25. Guilloteau, S., Dutrey, A., Piétu, V., & Boehler, Y. 2011, A&A, 529, A105 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  26. Guilloteau, S., Piétu, V., Chapillon, E., et al. 2016, A&A, 586, L1 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  27. Haisch, K. E. J., Lada, E. A., & Lada, C. J. 2001, ApJ, 553, L153 [NASA ADS] [CrossRef] [Google Scholar]
  28. Hollenbach, D. J., & Tielens, A. G. G. M. 1999, Rev. Mod. Phys., 71, 173 [NASA ADS] [CrossRef] [Google Scholar]
  29. Hughes, A. M., Duchêne, G., & Matthews, B. C. 2018, ARA&A, 56, 541 [NASA ADS] [CrossRef] [Google Scholar]
  30. Jonkheid, B., Kamp, I., Augereau, J.-C., & van Dishoeck, E. F. 2006, A&A, 453, 163 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  31. Konishi, M., Grady, C. A., Schneider, G., et al. 2016, ApJ, 818, L23 [NASA ADS] [CrossRef] [Google Scholar]
  32. Kóspál, Á., Moór, A., Juhász, A., et al. 2013, ApJ, 776, 77 [NASA ADS] [CrossRef] [Google Scholar]
  33. Kral, Q., Matrà, L., Wyatt, M. C., & Kennedy, G. M. 2017, MNRAS, 469, 521 [NASA ADS] [CrossRef] [Google Scholar]
  34. Kral, Q., Marino, S., Wyatt, M. C., Kama, M., & Matra, L. 2019, Turk. J. Phys., 43, 126 [CrossRef] [Google Scholar]
  35. Lazareff, B., Berger, J.-P., Kluska, J., et al. 2017, A&A, 599, A85 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  36. Lieman-Sifry, J., Hughes, A. M., Carpenter, J. M., et al. 2016, ApJ, 828, 25 [NASA ADS] [CrossRef] [Google Scholar]
  37. Lyra, W., & Kuchner, M. 2013, Nature, 499, 184 [NASA ADS] [CrossRef] [Google Scholar]
  38. Matrà, L., Panić, O., Wyatt, M. C., & Dent, W. R. F. 2015, MNRAS, 447, 3936 [NASA ADS] [CrossRef] [Google Scholar]
  39. Matrà, L., Dent, W. R. F., Wyatt, M. C., et al. 2017, MNRAS, 464, 1415 [NASA ADS] [CrossRef] [Google Scholar]
  40. Mawet, D., Choquet, É., Absil, O., et al. 2017, AJ, 153, 44 [NASA ADS] [CrossRef] [Google Scholar]
  41. Meeus, G., Montesinos, B., Mendigutía, I., et al. 2012, A&A, 544, A78 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  42. Merín, B., Montesinos, B., Eiroa, C., et al. 2004, A&A, 419, 301 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  43. Miley, J. M., Panić, O., Wyatt, M., & Kennedy, G. M. 2018, A&A, 615, L10 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  44. Miotello, A., van Dishoeck, E. F., Kama, M., & Bruderer, S. 2016, A&A, 594, A85 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  45. Miotello, A., van Dishoeck, E. F., Williams, J. P., et al. 2017, A&A, 599, A113 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  46. Moór, A., Ábrahám, P., Juhász, A., et al. 2011, ApJ, 740, L7 [NASA ADS] [CrossRef] [Google Scholar]
  47. Moór, A., Kóspál, Á., Ábrahám, P., et al. 2016, ApJ, 826, 123 [NASA ADS] [CrossRef] [Google Scholar]
  48. Moór, A., Curé, M., Kóspál, Á., et al. 2017, ApJ, 849, 123 [NASA ADS] [CrossRef] [Google Scholar]
  49. Mouillet, D., Lagrange, A. M., Augereau, J. C., & Ménard, F. 2001, A&A, 372, L61 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  50. Nilsson, R., Liseau, R., Brandeker, A., et al. 2010, A&A, 518, A40 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  51. Pacheco-Vázquez, S., Fuente, A., Agúndez, M., et al. 2015, A&A, 578, A81 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  52. Pérez, L. M., Carpenter, J. M., Chandler, C. J., et al. 2012, ApJ, 760, L17 [NASA ADS] [CrossRef] [Google Scholar]
  53. Péricaud, J., Di Folco, E., Dutrey, A., Guilloteau, S., & Piétu, V. 2017, A&A, 600, A62 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  54. Perrot, C., Boccaletti, A., Pantin, E., et al. 2016, A&A, 590, L7 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  55. Piétu, V., Dutrey, A., & Guilloteau, S. 2007, A&A, 467, 163 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  56. Pinilla, P., Birnstiel, T., Ricci, L., et al. 2012, A&A, 538, A114 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  57. Reboussin, L., Guilloteau, S., Simon, M., et al. 2015, A&A, 578, A31 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  58. Reche, R., Beust, H., & Augereau, J.-C. 2009, A&A, 493, 661 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  59. Schreyer, K., Guilloteau, S., Semenov, D., et al. 2008, A&A, 491, 821 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  60. Sylvester, R. J., Skinner, C. J., Barlow, M. J., & Mannings, V. 1996, MNRAS, 279, 915 [NASA ADS] [CrossRef] [Google Scholar]
  61. Sylvester, R. J., Dunkin, S. K., & Barlow, M. J. 2001, MNRAS, 327, 133 [NASA ADS] [CrossRef] [Google Scholar]
  62. Takeuchi, T., & Artymowicz, P. 2001, ApJ, 557, 990 [NASA ADS] [CrossRef] [Google Scholar]
  63. Thi, W.-F., Pinte, C., Pantin, E., et al. 2014, A&A, 561, A50 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  64. van der Marel, N., van Dishoeck, E. F., Bruderer, S., et al. 2016, A&A, 585, A58 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  65. van der Marel, N., Williams, J. P., Ansdell, M., et al. 2018, ApJ, 854, 177 [Google Scholar]
  66. van der Plas, G., van den Ancker, M. E., Waters, L. B. F. M., & Dominik, C. 2015, A&A, 574, A75 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  67. Visser, R., van Dishoeck, E. F., & Black, J. H. 2009, A&A, 503, 323 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  68. Weinberger, A. J., Becklin, E. E., Schneider, G., et al. 1999, ApJ, 525, L53 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
  69. Weinberger, A. J., Rich, R. M., Becklin, E. E., Zuckerman, B., & Matthews, K. 2000, ApJ, 544, 937 [NASA ADS] [CrossRef] [Google Scholar]
  70. White, J. A., & Boley, A. C. 2018, ApJ, 859, 103 [NASA ADS] [CrossRef] [Google Scholar]
  71. White, J. A., Boley, A. C., Hughes, A. M., et al. 2016, ApJ, 829, 6 [NASA ADS] [CrossRef] [Google Scholar]
  72. White, J. A., Boley, A. C., MacGregor, M. A., Hughes, A. M., & Wilner, D. J. 2018, MNRAS, 474, 4500 [NASA ADS] [CrossRef] [Google Scholar]
  73. Williams, J. P., & Best, W. M. J. 2014, ApJ, 788, 59 [NASA ADS] [CrossRef] [Google Scholar]
  74. Williams, J. P., & Cieza, L. A. 2011, ARA&A, 49, 67 [NASA ADS] [CrossRef] [Google Scholar]
  75. Wilson, T. L.,& Rood, R. 1994, ARA&A, 32, 191 [NASA ADS] [CrossRef] [Google Scholar]
  76. Wyatt, M. C. 2008, ARA&A, 46, 339 [NASA ADS] [CrossRef] [Google Scholar]
  77. Wyatt, M. C., Panić, O., Kennedy, G. M., & Matrà, L. 2015, Ap&SS, 357, 103 [NASA ADS] [CrossRef] [Google Scholar]
  78. Zhu, Z., Nelson, R. P., Hartmann, L., Espaillat, C., & Calvet, N. 2011, ApJ, 729, 47 [NASA ADS] [CrossRef] [Google Scholar]
  79. Zuckerman, B., Forveille, T., & Kastner, J. H. 1995, Nature, 373, 494 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]

1

Based on observations carried out with the IRAM Plateau de Bure Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain).

All Tables

Table 1

Continuum results: best-fit parameters foro: (i) a single component model (elliptical Gaussian in the u, v plane), (ii) and a double-component model (two circular Gaussian functions after deprojection in the u, v plane according to the orientation and inclination determined at optical wavelength).

Table 2

Continuum flux comparison.

Table 3

Summary of CO lines detections and setups.

Table 4

Results of Diskfit modelling for the 12CO and 13CO with 1σ error bars.

All Figures

thumbnail Fig. 1

Left column: continuum emission maps, all contours are set with 3σ spacing, panel a: ALMA data at 0.87 mm, σ = 0.051 mJy beam−1, contours; (c) NOEMA data at 1.3 mm, σ = 0.078 mJy beam−1; (e) ALMA data at 2.8 mm, σ = 0.008 mJy beam−1. We note that the scale of panel c is changed to include the more extended emission. Right column: real part of the visibility (with spatial binning) as a function of the u, v radius in the Fourier plane, panels b, d, and f: correspond to data in panels a, c, and e, respectively.

Open with DEXTER
In the text
thumbnail Fig. 2

Maps of integrated CO emission for the four detected transitions (the continuum was not subtracteddue to its very low level of emission). Top: NOEMA observations, from left to right: 12CO J = 2 → 1 (first contour at 5σ, second contour at 15σ, then 15σ spacing withσ = 31 mJy beam−1 km s−1), and 13CO J = 2 → 1 with 3σ spacing contours, σ = 23 mJy beam−1 km s−1. Dotted lines show negative contours in steps of 3σ. Bottom:ALMA observations, from left to right: 12CO J = 3 → 2 with uniform weighting to highlight the central gas cavity (3σ spacing, σ = 40 mJy beam−1 km s−1), and 13CO J = 1 → 0 1σ spacing contours starting at 2σ (for positive and negative contours), σ = 6.8 mJy beam−1 km s−1.

Open with DEXTER
In the text
thumbnail Fig. 3

Integrated intensity of the CO line superimposed on the HST scattered emission (Clampin et al. 2003). The cross indicates the position angle and aspect ratio as determined from gas modelling (Sect. 3). Left: 12CO J = 2 → 1 emission, the contour spacing is 6σ, i.e. 5.5 × 10−1 Jy beam−1 km s−1. Beam size: 2.48 × 1.45′′. Right: 13CO J = 2 → 1 emission, with a 3σ contour spacing, i.e. 7.8 × 10−2 Jy beam−1 km s−1. Beam size: 1.76 × 1.32′′.

Open with DEXTER
In the text
thumbnail Fig. 4

Observed line profiles and superposed best models obtained with Diskfit (colour dashed lines) for, from left to right: 12CO J = 2 → 1, 13CO J = 2 → 1 NOEMA data, and 12CO J = 3 → 2 ALMA data.

Open with DEXTER
In the text
thumbnail Fig. 5

Results of the modelling showing the channel maps for the transition (from left to right) 12CO J = 2 → 1, 13CO J = 2 → 1 (NOEMA) and 12CO J = 3 → 2 (ALMA). In each sub-panel, we present from left to right: the observations, the model (see best-fit parameters in Table 4), and the residual emission. The latter is obtained by subtraction in the u, v plane before imaging and cleaning. Contour spacing is 5σ for the 12CO transitions (σ = 16 mJy beam−1 for J = 2 → 1 and σ = 6.3 mJy beam−1 for J = 3 → 2), and 3σ for the 13CO J = 2 → 1 transition (σ = 10 mJy beam−1). Residual contour spacing is 3σ for all transitions.

Open with DEXTER
In the text
thumbnail Fig. 6

Global scheme of dust and gas distribution in HD 141569 summarizing the most recent resolved observations.

Open with DEXTER
In the text
thumbnail Fig. 7

Montage of the 12CO J = 3 → 2 residual emission with the series of concentric rings detected in NIR scattered light with SPHERE by Perrot et al. (2016).

Open with DEXTER
In the text
thumbnail Fig. A.1

Channel map for the 12CO J = 2 → 1 transition observed with NOEMA, the spatial resolution is 2.5 × 1.4′′, the contour spacing is 7σ.

Open with DEXTER
In the text
thumbnail Fig. A.2

Channel map for the 13CO J = 2 → 1 transition observed with NOEMA, the spatial resolution is 1.9 × 1.1′′, the contour spacing is 2σ, same channel selection as for Fig. A.1.

Open with DEXTER
In the text
thumbnail Fig. A.3

Channel map for the 13CO J = 1 → 0 transition observed with ALMA, the spatial resolution is 0.7 × 0.5′′, the contour spacing is 3σ, same channel selection as for Fig. A.1.

Open with DEXTER
In the text

Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.

Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.

Initial download of the metrics may take a while.