Discovery of an au-scale excess in millimeter emission from the protoplanetary disk around TW Hya

TL;DR

This paper reports the discovery of a small, bright dust clump in the protoplanetary disk around TW Hya, likely caused by a vortex or forming planet, using high-resolution ALMA observations at 1.3 mm.

Contribution

First detection of a small-scale dust excess in TW Hya's disk, revealing potential planet formation activity with high-resolution ALMA data.

Findings

Detected a 1.5 times brighter dust excess at 52 au from the star.

Estimated dust mass of the excess is approximately 0.03 Earth masses.

The feature is unlikely to be a background source, indicating a local disk structure.

Abstract

We report the detection of an excess in dust continuum emission at 233~GHz (1.3~mm in wavelength) in the protoplanetary disk around TW~Hya revealed through high-sensitivity observations at $\sim$3~au resolution with the Atacama Large Millimeter/submillimeter Array (ALMA). The sensitivity of the 233~GHz image has been improved by a factor of 3 with regard to that of our previous cycle 3 observations. The overall structure is mostly axisymmetric, and there are apparent gaps at 25 and 41 au as previously reported. The most remarkable new finding is a few au-scale excess emission in the south-west part of the protoplanetary disk. The excess emission is located at 52 au from the disk center and is 1.5 times brighter than the surrounding protoplanetary disk at a significance of 12$\sigma$. We performed a visibility fitting to the extracted emission after subtracting the axisymmetric protoplanetary disk emission and found that the inferred size and the total flux density of the excess emission are 4.4$\times$1.0~au and 250~$\mu$Jy, respectively. The dust mass of the excess emission corresponds to 0.03~$M_\oplus$ if a dust temperature of 18~K is assumed. Since the excess emission can also be marginally identified in the Band 7 image at almost the same position, the feature is unlikely to be a background source. The excess emission can be explained by a dust clump accumulated in a small elongated vortex or a massive circumplanetary disk around a Neptune mass forming-planet.

Equations

$$M_{\mathrm{dust}}=\frac{F_{\nu }{d}^2}{{\kappa }_{\nu }B(T_{\mathrm{dust}})},$$

Full text

protoplanetary disks — stars: individual(TW Hya)

††journal: ApJ††facilities: ALMA††software: astropy, CASA,

1 Introduction

It is widely accepted that protoplanetary disks (PPDs) are the birthplace of planets. Obtaining observational evidence of a forming planet in PPDs is crucial for the understanding of the formation of and the diversity of (exo-)planets (Ida & Lin, 2004). The detection of a substructure related to a forming planet is a promising way to investigate the planet formation process. Recent high-resolution observations using ALMA have revealed complex disk substructures, such as multiple, axisymmetric gaps, spiral arms, and lopsided emissions (e.g., DSHARP; Andrews et al., 2018), which are likely related to the planet formation process.

On the other hand, it is theoretically predicted that the forming planet accretes material from a circumplanetary disk (CPD). Such CPDs hold the promise of direct detection with sensitive detectors as a localized small-scale substructure in the PPD, and some theoretical models suggest that CPDs should be detectable at millimeter wavelengths (Zhu et al., 2016). Large-scale non-axisymmetric substructure in PPDs has been found at millimeter wavelengths for Oph IRS 48 ($\sim$60 au in azimuthal extent; van der Marel et al., 2013), HD 142527 ($\sim$150 au; Fukagawa et al., 2013; Casassus et al., 2013), MWC 758 ($\sim$30 au; Dong et al., 2018), HD143006 ($\sim$86 au; Pérez et al., 2018), HD 163296 ($\sim$17 au; Isella et al., 2018), and HD 135344B ($\sim$210 au; Cazzoletti et al., 2018). Those lopsided substructures are interpreted as the accumulation of dust particles due to a large-scale gas vortex in the PPD. On the other hand, au-scale substructures or non-axisymmetric components, which are expected to be the signatures of CPDs, have not yet been discovered at these wavelengths (Isella et al., 2014).

TW Hya is the nearest T Tauri star with a distance of 59.5 pc (Gaia Collaboration et al., 2016). The stellar mass is 0.8 $M_{\sun}$ and the stellar age is 10 Myr (Andrews et al., 2012). The disk orientation is almost face-on with an inclination of 7$\arcdeg$ (Qi et al., 2004). High-resolution observations with ALMA have resolved multiple axisymmetric gaps; in particular, deep gaps at $\sim$25 and 41 au (Andrews et al., 2016; Tsukagoshi et al., 2016; Huang et al., 2018). Evidence for non-axisymmetry or small-scale substructure has not yet been reported for the PPD around TW Hya.

In this paper, we report the results of our high-sensitivity observations at Band 6 using ALMA and the first finding of an au-scale substructure at a radius of 52 au in the PPD around TW Hya. We describe our observations and data reduction in §2. In §3, we show the resulting images of the Band 6 continuum emission, and the finding of the au-scale substructure in the TW Hya disk is presented. In §4, we discuss the detailed structures and the expected origin of the substructure.

2 Observations

2.1 Sensitive ALMA Observation at Band 6

Our 233 GHz (1.3 mm) continuum observations at Band 6 toward TW Hya were carried out using ALMA on 15 May 2017 with array configuration C40-5 (2016.1.00842.S) and in the period from 20 to 25 November 2017 with C43-8 (2017.1.00520.S). The total on-source integration times were $\sim$12 and $\sim 200$ minutes, respectively. The correlator was configured to detect dual polarizations in four spectral windows with a bandwidth of 1.875 GHz each, resulting a bandwidth of 7.5 GHz in total. The four spectral windows were tuned to detect continuum emission centered at 225, 227, 239, and 241 GHz. The phase fluctuations of the complex gain due to atmospheric noise were calibrated by observing quasars J1051-3138 or J1103-3251. Quasars J1107-4449, J1058+0133, or J1037-2934 were used for the calibration of the bandpass characteristics and flux scales were determined by observing J1058+0133, J1107-4449, or J1037-2934.

The observed visibilities were reduced and calibrated using the Common Astronomical Software Application (CASA) package (McMullin et al., 2007). The initial flagging of the visibilities and the calibrations for the bandpass characteristics, complex gain, and flux scaling were performed using the pipeline scripts provided by ALMA. After flagging the bad data, the corrected data were concatenated and imaged by CLEAN. The CLEAN map was created by adopting Briggs weighting with a robust parameter of 0.5. We also employed the multiscale CLEAN with scale parameters of [0, 42, 126] mas for better reconstruction of extended emission. After that, self-calibration was applied for the concatenated data set. The spatial resolution of the CLEANed image was 46.88$\times$41.56 mas, full width at half maximum (FWHM) with a position angle of $-78.9^{\circ}$, corresponding to 2.79$\times$2.48 au. The 1$\sigma$ root-mean-square (rms) noise level achieved was 9.1 $\mu$Jy beam*-1*.

2.2 Data Reduction of the Band 7 Archive Data

To validate the new results revealed by our Band 6 observation, we have also analysed eight sets of ALMA archival data at Band 7 from Cycle 3. The highest resolution data obtained by Andrews et al. (2016) is included, which is taken at the beginning of December 2015, two years before our observations. Data reduction has been done in the same manner as for our Band 6 observations described above. Line free channels were used for making the map of continuum emission at Band 7. The center frequency of the concatenated data was 325.4 GHz. The calibrated data were concatenated and imaged by CLEAN with an iterative self-calibration. For the imaging, we adopt Briggs weighting with a robust parameter of -1.0 so that the synthesized beam can be described well by a single Guassian function. Multiscale CLEAN with scale parameters of [0, 30, 90] mas was employed. The final CLEANed image had a synthesized beam size of 36.4$\times$28.9 mas (2.2$\times$1.7 au) with a position angle of 69.9∘. To reduce phase noise, we have smoothed the final CLEANed image to 50 mas resolution using the imsmooth task on CASA. The rms noise level of the resultant image is 27.7 $\mu$Jy beam*-1*.

3 Results

Figure 1a shows the global distribution of the 233 GHz continuum emission. The total flux density from the PPD and the known axisymmetric features, gaps at 25 and 41 au, are consistent with previous works (Andrews et al., 2016; Tsukagoshi et al., 2016; Huang et al., 2018).

In the south-west part of the PPD, the resolved excess of emission is discovered at a radius of 52 au, as shown in Fig. 1b. The peak intensity at this feature is measured to be 308.4 $\mu$Jy beam*-1*, corresponding to signal-to-noise (SN) ratio of 34. The excess of emission is significant in comparison to the surrounding background emission of the PPD. The excess is clearly detected both in the radial and in the azimuthal profiles as shown in Figure 2. The average of the surrounding emission is measured to be 197.1 $\mu$Jy beam*-1* from position angles of 226–230∘ and 244–248∘; hence the excess is calculated to be 111.5$\pm$9.1 $\mu$Jy beam*-1* ($\sim 12\sigma$), i.e., the emission feature is 1.5 times brighter than the surrounding PPD. There is no clear excess of emission above 3$\sigma$ in this radial region of the PPD except for the emission feature. We also note that no clear gap-like structure is found at 52 au from the central star.

Since the emission from the PPD is almost wholly axisymmetric, the emission feature is extracted by subtracting the axisymmetric background emission in the visibility (i.e., $u$-$v$) domain. The size, position and total flux density of the emission feature are measured by model fitting in the $u$-$v$ domain using the residual visibilities.

The subtraction process works well as shown in Figure 3a and b. To deduce the structure of the emission feature, we performed a least squares fitting to the subtracted visibilities using the uvmodelfit task installed on CASA. A single component described by a Gaussian function was assumed for the model. The modeled visibilities determined by the fitting were converted to an image using CLEAN with the same parameters as for the original image, and the resulting image is shown in Fig. 3c. From the fitting, the peak position of the emission feature is measured to be ($\Delta\alpha$, $\Delta\delta$)$=$($-726\pm 1$, $-471\pm 1$) mas from the central star, corresponding to ($-43.2\pm 0.1$, $-28.0\pm 0.1$) au or a radius of 51.5$\pm$0.1 au. The FWHMs of major and minor axes of the fitted Gaussian are measured to be (74.7$\pm$3.3)$\times$(16.2$\pm$3.2) mas with a position angle of $-38.3^{\circ}\pm 2.1^{\circ}$, corresponding to (4.4$\pm$0.2)$\times$(1.0$\pm$0.2) au. The total intensity of the fitted Gaussian is 250$\pm$5 $\mu$Jy.

The brightness temperature of the emission feature can be converted from the total flux density. Assuming that the millimeter emission of the PPD behind the resolved feature is uniform, the total flux density of the background inside the feature area is estimated to be 122 $\mu$Jy from the average of the surrounding emission, and thus the summation of the total flux density inside this area is 372 $\mu$Jy. Therefore, the brightness temperature at the position of the emission feature is converted using the Planck function to be 11.6 K. The brightness temperature is less than the temperature of the dust disk at 52 au, where $\sim$14–18 K is expected from the midplane temperature profile (Andrews et al., 2016; Tsukagoshi et al., 2016; Huang et al., 2018), indicating that the emission feature may be partially optically thin.

the mass of the emission feature is estimated from the equation

[TABLE]

where $F_{\nu}$ is the integrated flux density, $d$ is the distance to the source, and $B(T_{\mathrm{dust}})$ is the Planck function for the dust temperature $T_{\mathrm{dust}}$. We employ a dust mass opacity coefficient $\kappa_{\nu}$ of 2.3 cm2 g*-1* at a frequency of 233 GHz, which is determined from $\kappa_{\nu}$ of 2.8 cm2 g*-1* at a wavelength of 870 $\mu$m (Andrews et al., 2012) and $\beta$, the power-law index of $\kappa_{\nu}$, of $\sim 0.5$. Assuming $T_{\mathrm{dust}}=18$ K (Huang et al., 2018), the total dust mass of the emission feature is (2.83$\pm$0.06)$\times 10^{-2}$ $M_{\oplus}$, where $M_{\oplus}$ indicates the Earth mass.

The emission feature can also be seen in a high-resolution 325 GHz ALMA image within close proximity to the location of our 233 GHz emission feature. The 325 GHz continuum image we reconstructed from the archive data is shown in Figure 4a and b, and the subtracted image made by using the same procedure as done for our 233 GHz data is shown in Figure 4c and d. It is clear that there is a local emission peak near the position of the emission feature that we found in the 233 GHz map, while the emission seems to be azimuthally elongated. The positional offset between the emission feature at Band 6 and the residual emission at Band 7 is much less than the proper motion of the TW Hya system. If the excess emission is a background source, the positional offset must be 136 mas for 2 yrs in the R.A. direction according to the proper motion of the TW Hya system (($-68.225$, $-13.934$) mas yr*-1*; Gaia Collaboration et al., 2016). The positional deviation of the residual emission at Band 7 is within only $\sim$50 mas with respect to the emission feature at Band 6, as shown in Figure 4. Therefore, we conclude that the emission feature we found is situated in the PPD and is likely orbiting the central star. Due to the significant phase noise of the Band 7 data, we conservatively conclude that it is unclear whether there is a positional offset due to the Keplerian motion for 2 years. (see Fig. 4d).

Lastly, we also note that the millimeter flux density fluctuates azimuthally (see the profile at 45 au in Fig. 2b). The fluctuation might be related to a moving surface brightness asymmetry which is probably due to disk shadow (Debes et al., 2017), but we do not focus on this structure in this paper.

4 Discussion

We have found a few au-scale, elongated emission feature in the PPD of TW Hya. Similar asymmetric structures have been found in several other PPDs as listed in §1, but the feature in the TW Hya disk is the smallest ever discovered.

In many PPDs, the asymmetric features are interpreted as the dust particles trapped in a gas vortex (e.g., Raettig et al., 2015). The morphology of the emission feature we found may also be interpreted along this line. The width of the gas vortex is limited to several times the gas scale height. The radial half width of the emission feature is only $\sim$0.5 au and is much smaller than the expected gas scale height at this radius which is $\sim$4–5 au (Andrews et al., 2012). The emission feature is azimuthally elongated and the ratio of the azimuthal to the radial widths, or aspect ratio, is $\sim$4. The shape of gas vortices has been investigated by several authors (Lesur & Papaloizou, 2009; Richard et al., 2013), and the aspect ratio of stable vortices is of the order of $\sim$several. The dust particles trapped in a gas vortex is more concentrated at the vortex center, but the aspect ratio of the dust distribution is similar to that of gas (Lyra & Lin, 2013).

The mean surface density of the emission feature is expected to be $\sim$2 times higher than that of the surrounding PPD under the assumption that the temperature at 52 au is constant at 18 K (Huang et al., 2018). If the overall dust-to-gas mass ratio of the TW Hya disk is $\sim$100, the excess emission region may have a dust-to-gas mass ratio of around 50. Such a small overdensity may be realized even with a weak gas vortex.

The gas disk may be full of such “weak vortices” if the PPD is moderately turbulent. A chain of vortices in the same radial location will merge into one vortex (e.g., Ono et al., 2018), but there may be more vortices at different radii. In this sense, the emission feature could be a “tip of the iceberg” of even smaller dust concentrations.

An alternative scenario that may explain why there is only one emission feature is the existence of a planet. If there is a planet accreting gas and dust from the surrounding PPD (e.g., Pollack et al., 1996; Canup & Ward, 2002), the temperature and the density may increase locally (e.g., Owen, 2014)

However, it should be noted that the observed emission feature may not be fully accounted for by emission from a CPD. We use a simple model by Zhu et al. (2016) to estimate the flux density at millimeter wavelengths from the CPD

around a Neptune-mass planet.

Figure 5 shows the flux density of the modeled CPDs for a range of viscous parameters $\alpha$ and mass accretion rates onto the planet $\dot{M}$. If the mass accretion rate is $10^{-7}~{}M_{\mathrm{Nep}}$ per year, the flux density of the CPD is no larger than $\sim 100~{}\mu$Jy.

This value may be compared with the observed emission within a circle with a diameter of 1 au,

Therefore, a part of the excess emission may come from the CPD while the entire emission may be accounted for by the overdensity due to a surrounding envelope-like structure of accreting material around the system of the planet and the CPD.

With observations at millimeter wavelengths only, it seems difficult to judge whether the hypothetical planet actually exists at the location of the emission feature. The simple CPD model described above suggests that the millimeter emission from the CPD is similar to that of dust concentration within a gas vortex, making it difficult to determine whether the planet exists or not. The most prominent feature of an accreting planet is that the material heated up to $\gtrsim 1000$ K in the close vicinity of the planet. Therefore, direct imaging observations at infrared wavelengths or detection of an accretion signature in emission lines may be critical to prove (or rule out) the existence of a planet.

This paper makes use of the following ALMA data: ADS/JAO.ALMA#2017.1.00520.S and 2016.1.00842.S. We also use the following public ALMA archive data: ADS/JAO.ALMA#2015.1.00308.S, 2015.1.00686.S, 2016.1.00229.S, 2016.1.00311.S, 2016.1.00440.S, 2016.1.00464.S, 2016.1.00629.S and 2016.1.01495.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. A part of the data analysis was carried out on the common-use data analysis computer system at the Astronomy Data Center of NAOJ. This work is partially supported by JSPS KAKENHI grant number 17K14244. TJM thanks STFC for support under grant reference ST/P000321/1. CW acknowledges financial support from STFC (grant reference ST/R000549/1) and the University of Leeds.