New constraints on the disk characteristics and companion candidates around T Cha with VLT/SPHERE
A. Pohl, E. Sissa, M. Langlois, A. Müller, C. Ginski, R. G. van Holstein, A. Vigan, D. Mesa, A.-L. Maire, Th. Henning, R. Gratton, J. Olofsson, R. van Boekel, M. Benisty, B. Biller, A. Boccaletti, G. Chauvin, S. Daemgen, J. de Boer, S. Desidera, C. Dominik, A. Garufi, M. Janson, Q. Kral, F. Ménard, C. Pinte, T. Stolker, J. Szulágyi, A. Zurlo, M. Bonnefoy, A. Cheetham, M. Cudel, M. Feldt, M. Kasper, A.-M. Lagrange, C. Perrot, F. Wildi
AAstronomy & Astrophysics manuscript no. pohl_ms_final c (cid:13)
ESO 2017May 25, 2017
New constraints on the disk characteristics and companioncandidates around T Cha with VLT/SPHERE (cid:63)
A. Pohl , , E. Sissa , , M. Langlois , , A. Müller , , C. Ginski , , R. G. van Holstein , A. Vigan , D. Mesa ,A.-L. Maire , Th. Henning , R. Gratton , J. Olofsson , , , R. van Boekel , M. Benisty , B. Biller , ,A. Boccaletti , G. Chauvin , S. Daemgen , J. de Boer , S. Desidera , C. Dominik , A. Garufi , M. Janson , ,Q. Kral , F. Ménard , C. Pinte , T. Stolker , J. Szulágyi , A. Zurlo , , , M. Bonnefoy , A. Cheetham ,M. Cudel , M. Feldt , M. Kasper , A.-M. Lagrange , C. Perrot , and F. Wildi (A ffi liations can be found after the references) May 25, 2017
ABSTRACT
Context.
The transition disk around the T Tauri star T Cha possesses a large gap, making it a prime target for high-resolution imagingin the context of planet formation.
Aims.
We aim to find signs of disk evolutionary processes by studying the disk geometry and the dust grain properties at its surface,and to search for companion candidates.
Methods.
We analyze a set of VLT / SPHERE data at near-infrared and optical wavelengths. We performed polarimetric imaging ofT Cha with IRDIS (1.6 µ m) and ZIMPOL (0.5–0.9 µ m), and obtained intensity images from IRDIS dual-band imaging with simulta-neous spectro-imaging with IFS (0.9–1.3 µ m). Results.
The disk around T Cha is detected in all observing modes and its outer disk is resolved in scattered light with unprecedentedangular resolution and signal-to-noise. The images reveal a highly inclined disk with a noticeable east-west brightness asymmetry. Thesignificant amount of non-azimuthal polarization signal in the U φ images, with a U φ / Q φ peak-to-peak value of 14%, is in accordancewith theoretical studies on multiple scattering in an inclined disk. Our optimal axisymmetric radiative transfer model considers twocoplanar inner and outer disks, separated by a gap of 0 (cid:48)(cid:48) .
28 ( ∼
30 au) in size, which is larger than previously thought. We derive adisk inclination of ∼
69 deg and PA of ∼
114 deg. In order to self-consistently reproduce the intensity and polarimetric images, thedust grains, responsible for the scattered light, need to be dominated by sizes of around ten microns. A point source is detected at anangular distance of 3.5 (cid:48)(cid:48) from the central star. It is, however, found not to be co-moving.
Conclusions.
We confirm that the dominant source of emission is forward scattered light from the near edge of the outer disk. Ourpoint source analysis rules out the presence of a companion with mass larger than ∼ M jup between 0 (cid:48)(cid:48) . (cid:48)(cid:48) .
3. The detection limitdecreases to ∼ M jup for 0 (cid:48)(cid:48) . (cid:48)(cid:48) . Key words.
Stars: individual: T Cha – Protoplanetary disks – Techniques: polarimetric – Radiative transfer – Scattering – Circum-stellar matter
1. Introduction
Recently developed high-resolution and high-contrast imaginginstruments provide the excellent capability to directly obtainimages of protoplanetary disks in scattered light and thermalemission. Protoplanetary disks are optically thick in the opticaland near-infrared (NIR), so that scattered light imaging probes(sub-)micron-sized dust grains in the disk surface layer. Contrar-ily, (sub-)mm observations trace larger, mm-sized grains locatedin the disk midplane. The detection of (non-)axisymmetric diskfeatures is fundamental in improving our current understandingof disk evolution and planet formation. Transition disks withlarge gas and dust gaps (see e.g. mid-infrared surveys by Brownet al. 2007; Merín et al. 2010; van der Marel et al. 2016 andmm observations by Isella et al. 2010a, 2010b; Andrews et al.2011; van der Marel et al. 2015) are particularly interestingtargets, since they host possible planet-forming hotspots andmay show signposts of planet-disk interaction processes. In (cid:63)
Based on observations made with European Southern Observatory(ESO) telescopes at the Paranal Observatory in Chile, under programIDs 095.C-0298(B), 096.C-0248(B) and 096.C-0248(C). recent observational studies, giant gaps and cavities have beendirectly imaged in scattered light observations of transition disksystems (e.g., Thalmann et al. 2010, 2015; Hashimoto et al.2012; Avenhaus et al. 2014; Follette et al. 2015; Ohta et al.2016; Stolker et al. 2016; Benisty et al. 2017).T Chamaeleontis (T Cha) is a ∼ ± Article number, page 1 of 17 a r X i v : . [ a s t r o - ph . E P ] M a y & A proofs: manuscript no. pohl_ms_final at the Very Large Telescope Interferometer (VLTI) to study theinner disk’s structure. The inner disk is found to be extremelynarrow and located close to the star with an extension from 0.13to 0.17 au. Olofsson et al. (2013) presented a radiative transfermodel accounting for several further interferometric and photo-metric observations, including VLTI / PIONIER, VLTI / MIDI andNACO / Sparse Aperture Masking (SAM) data, which furtherconstrains the inner disk to extend from 0.07 to 0.11 au. FurtherSED modeling of T Cha by Cieza et al. (2011) suggests thatthere is a high degeneracy especially for the outer disk geometry,since a very compact outer disk provides an equally good fitto the Herschel data as a much larger, but tenuous disk witha very steep surface density profile. High-resolution AtacamaLarge Millimeter / sub-millimeter Array (ALMA) observationsof the 850 µ m dust continuum as well as of several emissionlines presented by Huélamo et al. (2015) spatially resolved theouter disk around T Cha and helped to break the degeneracy ofprevious outer disk models. They report a compact dusty disk,where the continuum intensity profile displays two emissionbumps separated by 40 au, indicating an inner gap size of 20 auand an outer disk radius of ∼
80 au. In contrast, the gaseousdisk is larger by almost a factor of three, giving a radius of ∼
230 au based on the detection of CO(3–2) molecular emission.Huélamo et al. (2015) derived a disk inclination (incl) of 67 ◦ ± ◦ and a position angle (PA) of 113 ◦ ± ◦ by fitting a Gaussian tothe CO(3–2) integrated emission map.All previous observations clearly confirm that there must bea significant gap in the disk dust density distribution, while itsorigin is still debated. In general, radial gap structures can becreated by a number of processes, including grain growth (e.g.,Dullemond & Dominik 2005), e ff ects of the magneto-rotationalinstability (MRI) at the outer edge of a dead-zone (e.g., Flocket al. 2015), a close (sub-)stellar companion or the dynamicalinteraction of a planet formed within the disk (e.g., Rice et al.2003). For the latter, the disk density modification results fromthe torques exerted on the disk by the planet and by the diskitself. The planet pushes away the surrounding material, theouter part of the disk outward and the inner part inward, therebyopening a gap (e.g., Lin & Papaloizou 1979; Crida et al. 2006).Studies of observational signatures of planet-disk interactionprocesses based on numerical simulations suggest that gapsdetected in scattered light may be opened by planets (e.g.,Pinilla et al. 2015; Dong et al. 2015, 2016; Juhász et al. 2015;Pohl et al. 2015). Using NACO / SAM Huélamo et al. (2011)detected a companion candidate at a projected distance of 6.7 aufrom the primary, which is well within the previously describeddisk gap. However, an analysis of several L’ and K s data setscovering a period of three years ruled out this companionhypothesis (Olofsson et al. 2013; Cheetham et al. 2015). Theabsence of relative motion for the companion candidate favors astationary structure consistent with scattered light from a highlyinclined disk. Sallum et al. (2015) checked if the closure phasesignal from their VLT / NACO and Magellan / MagAO / Clio2data shows any variation in time, which is not expected forthe disk scattering model. While NACO L’ data from 2011and 2013 support the hypothesis of constant scattered lightfrom the disk, the best fits for two other NACO data sets areinconsistent, requiring temporal variability in the amount ofscattered light. Apart from this variability argument, Sallumet al. (2015) showed by means of Monte Carlo simulations thatnoise fluctuations could also cause the changing structure in theNACO and MagAO reconstructed images. In this work we present the first scattered light observa-tions of T Cha obtained with the SPHERE instrument (Spectro-Polarimeter High contrast Exoplanet REsearch, Beuzit et al.2008) at the Very Large Telescope (VLT). The target is now oneof the few T Tauri stars to have been spatially resolved in highdetail in scattered light. We used Polarimetric Di ff erential Imag-ing (PDI) complemented with total intensity images obtainedwith the angular di ff erential imaging (ADI) technique. The ob-servations provide the first spatially resolved high-contrast im-ages of T Cha in the optical and NIR. Our focus is set on an-alyzing the scattered light properties of the disk. We performphysical modeling of the disk via radiative transfer calculations,which helps us to further constrain the disk’s geometry and grainsize distribution. Images in the full Stokes vector are calculatedin order to consistently reproduce the observed total and polar-ized intensity. Furthermore, the total intensity images are usedfor a detailed search for substellar companion candidates and, incase of non-detection, to place constraints on the mass of puta-tive companions using the detection limits. This paper is laid outas follows. In Sect. 2 we describe our observations and the datareduction procedures; their results are shown in Sect. 3. Section4 presents our results of the radiative transfer disk modeling andthe search for substellar companions. A detailed discussion inSect. 5 follows. In Sect. 6 we summarize the main conclusionsof this work.
2. SPHERE observations and data reduction
Observations of T Cha were performed during the nights of30 May 2015, 19 February 2016, and 31 March 2016 with sev-eral sub-systems of the high-contrast imager SPHERE equippedwith an extreme adaptive optics system (SAXO, Fusco et al.2006, 2014) and mounted on the VLT at Cerro Paranal, Chile.All observations were part of the SPHERE consortium guar-anteed time program under IDs 095.C-0298(B) and 096.C-0248(B / C). The Infra-Red Dual-beam Imager and Spectrograph(IRDIS, Dohlen et al. 2008) and the Zurich IMaging POLarime-ter (ZIMPOL, Thalmann et al. 2008; Schmid et al. 2012) wereused in Dual-band Polarimetric Imaging (DPI) mode (Langloiset al. 2014) and in field stabilized (P2) mode, respectively. Inaddition, data were taken simultaneously with IRDIS in dual-band imaging (DBI; Vigan et al. 2010) mode and the IntegralField Spectrograph (IFS, Claudi et al. 2008). In this IRDIFSmode, IRDIS is operated in the filter pair
H2H3 (1 . µ m and1 . µ m) and IFS in YJ (0 .
95 – 1 . µ m) mode. Table 1 sum-marizes the observations and instrumental setups for each instru-ment. The Strehl ratio estimation (provided by SPARTA files) isbased on an extrapolation of the phase variance deduced fromthe reconstruction of SAXO open-loop data using a deformablemirror, tip-tilt voltages, and wavefront sensor closed-loop data(Fusco et al. 2004). The observing conditions and the di ff erentdata reduction methods for each data set taken by the varioussub-systems are described in detail in Sects. 2.1 – 2.3. The IRDIS-DPI observations of T Cha were carried out on19 February 2016 with the
BB_H filter ( λ c = . µ m) usingan apodized pupil Lyot coronagraph with a mask diameter of ∼
185 mas (Soummer 2005; Boccaletti et al. 2008). Dark and flatfield calibration were obtained during the following day. Thirtypolarimetric cycles were taken, consisting of one data cube foreach of the four half wave plate (HWP) positions (0 ◦ , 45 ◦ , 22.5 ◦ and 67.5 ◦ ). Dedicated coronagraphic images were taken at the Article number, page 2 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Table 1.
Overview of observational data sets
Date Instrument Mode Filter DIT [s] × NDIT PC t tot [min] Seeing [ (cid:48)(cid:48) ] H -band Strehl [%]2015 May 30 IRDIS IRDIFS DB_H2H3 ×
96 – 102 0.5–0.85 27 ± YJ ×
96 – 102 0.5–0.85 27 ± H × ± VBB × ± Notes.
Both filters of ZIMPOL were set to the Very Broad Band (
VBB ) filter covering a wide wavelength regime from R - to I -band. The followingcoronagraphs were used: N_ALC_YJH_SDIT for IRDIS / IFS and V_CLC_S_WF for ZIMPOL. DIT stands for the detector integration time andNDIT corresponds to the number of frames in the sequence. PC indicates the number of polarimetric cycles. The Strehl is calculated for the H -band. This leads to a Strehl of ∼
43% for the ZIMPOL data at 0.65 µ m. beginning and at the end of the science sequence to determineaccurately the position of the star behind the coronagraph. Forthis calibration a periodic amplitude is applied to the deformablemirror, which produces four equidistant, crosswise satellite spotsof the stellar PSF outside of the coronagraph. The data were re-duced following the prescriptions of Avenhaus et al. (2014) andGinski et al. (2016), who consider the radial Stokes formalism.The first step consists of standard calibration routines, includ-ing dark-frame subtraction, flat-fielding and bad-pixel correc-tion. These images are split into two individual frames represent-ing the left and right sides (parallel and perpendicular polarizedbeams, respectively), and the precise position of the central staris measured using the star center calibration frames on both im-age sides separately. Then, the right side of the image is shiftedand subtracted from the left side. To obtain clean Stokes Q and U images, that is, to correct for instrumental polarization down-stream of the HWP’s position in the optical path, Q + and Q − (0 ◦ and 45 ◦ ), and U + and U − (22.5 ◦ and 67.5 ◦ ) are subtracted,respectively. However, there might be still an instrumental polar-ization left upstream of the HWP in the final Q and U images,which is assumed to be proportional to the total intensity imageas shown in Canovas et al. (2011). To obtain this residual instru-mental signal, the azimuthal Stokes components are computedfrom (cf. Schmid et al. 2006) Q φ = + Q cos 2 φ + U sin 2 φ , (1) U φ = − Q sin 2 φ + U cos 2 φ . (2)The azimuth φ is defined with respect to the stellar position(x ,y ) as φ = arctan x − x y − y . (3)As shown by Canovas et al. (2015), the signal in the U φ frame should be small for a centrally illuminated symmetricaldisk. We thus determined the scaling factor for our secondinstrumental polarization correction such that the (absolute)signal in an annulus around the central star in the U φ frame isminimized. We then subtract the scaled Stokes I frame from the Q and U frame and use these final corrected frames to create the Q φ and U φ images displayed in Fig. 1.To cross-check the IRDIS-DPI results, and especially to testthe reliability of the U φ minimization technique in the context ofan inclined disk, we additionally perform an alternative reduc-tion procedure. This includes a proper polarimetric calibrationusing a Mueller matrix model, whose details will be presentedin van Holstein et al. (in prep.) and de Boer et al. (in prep.). Ashort explanation of the method and the corresponding resultsfor T Cha can be found in Sect. 3.1 and Fig. 2. T Cha was observed during the night of 31 March 2016 withthe SlowPolarimetry detector mode of ZIMPOL using thevery broad band filter (
VBB ). The
VBB filter covers a widewavelength range from R - to I -band (0.55–0.87 µ m). Theseobservations were also obtained with an apodized Lyot corona-graph (mask diameter of ∼
185 mas).The ZIMPOL data were reduced following mostly the samestrategy as described for the IRDIS data in the previous section.The main di ff erence between the two data sets is the di ff erentstructure of the ZIMPOL data. In ZIMPOL the two perpendic-ular polarization directions for each HWP position are recordedquasi-simultaneously on the same detector pixels. For a moredetailed description of the instrument and the specialized datareduction steps involved we refer to Thalmann et al. (2008) andSchmid et al. (2012). We process both ZIMPOL detector imagesindependently and only combine the images after the final datareduction to increase the S / N. We first bias subtract and flat fieldthe individual frames. We then extract the two perpendicular po-larization directions from the interlaced rows in each frame, re-sulting in two 1024 ×
512 pixel images per original frame. Wethen correct for charge-shifting artifacts by always combiningtwo consecutive frames of the observation sequence. To createquadratic images we then bin each image by a factor of twoalong the x-axis. Finally we subtract the two perpendicular po-larization directions from each other to get Q + , Q − , U + and U − frames (depending on the HWP position). These are then com-bined to create the final Q and U frames identical to the IRDISreduction. In a last step we again calculate the azimuthal Stokescomponents Q φ and U φ and employ the instrumental polariza-tion correction from Canovas et al. (2011). The resulting images(after combination of both ZIMPOL images) are also displayedin Fig.1. T Cha IRDIFS observations where obtained during the nightof 30 May 2015 as part of the SpHere INfrared survey forExoplanets (SHINE; Chauvin et al., in prep.) using the SHINEstandard setup: pupil-stabilized images with IFS operating in YJ mode (39 channels between 0.95 and 1.35 µ m) and IRDIS work-ing in DBI mode using the H2H3 filter pair ( λ H = . µ m; λ H = . µ m). This observing strategy allows for performingADI (Marois et al. 2006) in order to reach high contrast. Thespectral resolution of IFS YJ data amounts to R ∼
50. Theobservations lasted about 6100 seconds with a field rotation of ∼ ◦ . Since the target is located far to the south, obtaining agood rotation is challenging. The unstable weather conditions Article number, page 3 of 17 & A proofs: manuscript no. pohl_ms_final (Di ff erential Image Motion Monitor (DIMM) seeing variedfrom 0 (cid:48)(cid:48) . (cid:48)(cid:48) .
85 and clouds passing by) caused flux variationsof up to one order of magnitude during the sequence.The basic steps of the first IRDIS data reduction consistof flat-field and bad-pixel correction, cosmic ray detectionand correction, and sky subtraction. Because of the variableatmospheric conditions during the observations, a very strictframe selection is applied at the end of the basic reductionand eventually only 42 out of 96 frames were used. Thiscorresponds to selecting frames with Strehl ratio larger than ∼ ff ect of selfsubtraction for an extended source is minimized and the S / Nmaximized.To cross-check the IRDIS-DBI results, we perform a parallelreduction using the SPHERE Data Reduction and Handlingpipeline (DRH, Pavlov et al. 2008) implemented at the SPHEREData Center. This includes dark and sky subtraction, bad-pixelremoval, flat-field correction, anamorphism correction, andwavelength calibration. After these first steps the data weresorted according to their quality. Because of the di ffi cult observ-ing conditions, we use stringent frame selection (using 77 framesout of 96). This roughly corresponds to selecting frames withStrehl ratio greater than 15% and leads to an average H- bandStrehl of ∼ ffl e patternon the deformable mirror. Then, to remove the stellar halo and toachieve high-contrast, the data were processed with the SpeCalpipeline developed for the SHINE survey (Galicher et al., inprep.); this implements a variety of ADI-based algorithms:Classical Angular Di ff erential Imaging (cADI, Marois et al.2006), Template Locally Optimized Combination of Images(TLOCI, Marois et al. 2006) and PCA (Soummer et al. 2012;Amara & Quanz 2012). In the following we discuss the resultsbased on the TLOCI, and PCA images for the morphology andphotometric analyses. Separate reductions were performed forthe extraction of the disk. In particular, for the SpeCal PCAreduction a small number of PCA modes is used in order toenable optimal retrieval of the disk. The reduced numbers ofmodes, between two and four, are determined by maximizingthe SNR inside a region delimited by the disk location. For thecontrast curves, TLOCI images were considered because theyprovide the best compromise of contrast, stellar rejection, andthroughput correction for the point source detection.The data reduction for the IFS is performed using tools avail-able at the SPHERE Data Center at IPAG following the proce-dure described in Mesa et al. (2015) and in Zurlo et al. (2014). Using the SPHERE DRH software we apply the appropriate cal-ibrations (dark, flat, spectral positions, wavelength calibrationand instrument flat) to create a calibrated datacube composed of39 images of di ff erent wavelengths for each frame obtained dur-ing the observations. Similar to the procedure used for IRDIS, inorder to take into account the very variable weather conditions,we apply a frame selection resulting in 76 frames out of the orig-inal 96. A frame is considered as ‘bad’ if the adaptive optics loopopens or the star exits the coronagraph region, causing an excessof light in the central part of the image. For each frame two cen-tral areas with 20 and 160 pixels per side are defined, for whichthe flux ratio is determined. Frames are rejected by an automatedsorting if this ratio exceeds 130% of the median value. The posi-tion of the star behind the coronagraph is estimated from imageswith four satellite spots, symmetric with respect to the centralstar taken just before and after the standard coronagraphic obser-vations. Exploiting these images we are then able to define there-scaling factor for images at di ff erent wavelengths to maintainthe speckle pattern as stable as possible. Moreover, we are ableto combine those images using the PCA algorithm from Soum-mer et al. (2012) to implement both ADI and spectral di ff erentialimaging (SDI, Racine et al. 1999) in order to remove the specklenoise.
3. Results
The disk of T Cha is detected in all data sets presented in thisstudy. The analysis of the disk geometry primarily focuses onthe IRDIS-DPI and IRDIS-ADI images (see Sect. 4.1). Further-more, the ADI images are used to search for point-source sig-nals focusing on non-polarized companions, because this dataset reaches a higher contrast (Sects. 4.2 and 4.3).
Figure 1 shows the reduced Q φ and U φ images of the IRDIS-DPI H -band and ZIMPOL VBB observations of T Cha describedin Sects. 2.1 and 2.2. The dark central region correspondsto the area masked by the coronagraph. The disk is clearlydetected in the IRDIS Q φ image, which gives by far the bestquality view of the outer disk structure and its rim in scatteredlight for T Cha. Our SPHERE observations support a highdisk inclination with respect to the line of sight, in agreementwith the model by Huélamo et al. (2015). Scattered light isdetected out to a projected distance of ∼ (cid:48)(cid:48) .
39 ( ∼
42 au) fromthe central star concentrated in a bright arc with, however, asignificant di ff erence in brightness between the east and westsides (factor of ∼ U φ image there is some residual signalleft, which has usually been interpreted as instrumental e ff ectsor as imperfect centering of the images. However, since the U φ / Q φ peak-to-peak value amounts to 9% and owing to the highinclination of the disk around T Cha (69 deg is determined fromthe total intensity image modeling, see Sect. 4.1.2), multiplescattering (e.g., Bastien & Menard 1990; Fischer et al. 1996;Ageorges et al. 1996), that is, scattering of already polarizedlight, in the inner disk might be the prime contributor to thissignal. This is consistent with a theoretical study by Canovaset al. (2015), who found that even for moderate disk inclinationsmultiple scattering alone can produce significant non-azimuthalpolarization above the noise level in the U φ images. Theyshowed that the U φ / Q φ peak-to-peak value can even go up to50% for a disk inclination of 70 deg depending on the mass andgrain size distribution of the disk model. We note that the exact Article number, page 4 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. 1.
IRDIS-DPI H -band and ZIMPOL P2 VBB -filter Q φ (top row) and U φ (bottom row) images. North is up, east is left. All images arenormalized to the highest disk brightness. The dynamical range for the color scaling is the same for the two images of the top (1000) and bottomrow (20), respectively. An apodized Lyot coronagraph with a mask diameter of ∼
185 mas was used. The inner 0 (cid:48)(cid:48) .
18 are masked, represented bythe black circular area. Negative values of U φ are saturated at dark blue and dark red color, respectively. geometrical structure of the U φ signal might be influenced bythe reduction method described in Sect. 2.1 (correction for theinstrumental crosstalk by minimizing U φ ). Therefore, we are notgoing to force our model to also fit the U φ in addition to the Q φ .However, in order to prove that the U φ signal is indeed real, weevaluate our IRDIS-DPI data with a newly developed reductionmethod, independent of the one presented in Sect. 2.1. By usingthe detailed Mueller matrix model of van Holstein et al. (inprep.) and de Boer et al. (in prep.) we correct our measurementsfor instrumental polarization e ff ects. This model describes thecomplete optical path of SPHERE / IRDIS, i.e. telescope and instrument, and has been fully validated with measurementsusing SPHERE’s internal source and observations of unpolar-ized standard stars (van Holstein et al., in prep.). The imagesof Stokes Q and U incident on the telescope are computed bysetting up a system of equations describing every measurementof Q and U and solving it — for every pixel individually —using linear least-squares. The resulting Q φ image (Fig. 2,left panel) is very similar to the one from the first reduction(Fig. 1, top left panel). The new U φ image (Fig. 2, middle panel)has a higher accuracy than the one from the reduction thatminimizes U φ (Fig. 1, bottom left panel), in particular because Article number, page 5 of 17 & A proofs: manuscript no. pohl_ms_final
Fig. 2.
Mueller matrix model-corrected IRDIS-DPI Q φ (left), U φ (middle) and polarized intensity PI = (cid:112) Q + U (right) images. North is up, eastis left. Note that the Q φ and U φ images are not normalized / saturated here on purpose to emphasize their partially negative signal. The inner 0 (cid:48)(cid:48) . Fig. 3.
IRDIS-ADI
H2H3 -band images (mean across the wavelengths) based on three di ff erent reduction pipelines ( (cid:48)(cid:48) .
18 region masked by the coronagraph is represented by the black circular area. no assumptions about the angle of linear polarization of thesource are made to correct for the instrumental polarization.The U φ image is cleaner and shows more symmetry in thesense that there is also a strong signal to the south-west. On thetop right, the positive U φ signal from Fig.1, bottom left, is notvisible anymore. The right panel of Fig. 2 shows the polarizedintensity overplotted with polarization vectors representing theangle of linear polarization. This strengthens that there is a cleardeparture from azimuthal polarization. For the model-correctedimages the U φ / Q φ peak-to-peak value increases from 9% to14%, suggesting that some of the actual physical U φ signal hasbeen removed in the original reduction method due to the U φ minimization procedure. A detailed comparison between di ff er-ent reduction methods and the specific influence on the left-over U φ signal will be the topic of the two follow-up SPHERE papers.The optical images obtained with ZIMPOL (Fig. 1, rightpanel) corroborate the disk geometry, but the bad weather con-ditions and low Strehl ratio (43% at 0.65 µ m) of this observationlead to a rather blurred structure. Again, positive and negativepatterns (dark and bright color) alternate in the U φ image, wherethese negative patterns are practically at the same location as inthe IRDIS U φ image. In addition to the polarimetric images, the IRDIS-ADI
H2H3 intensity images in Fig. 3 also clearly show the inclined diskaround T Cha. It even more strongly brings out the inner rim of
Article number, page 6 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. 4.
IFS image after PCA + SDI reduction with 100 modes: medianacross the entire wavelength range YJ . North is up, east is toward theleft; the image is normalized to the highest disk brightness and the colorscale considers the same dynamical range as in Fig.3. The inner 0 (cid:48)(cid:48) . the outer disk on the far side, visible as a faint arc below the coro-nagraph. The double-arch structure is a recurrent new form offeatures we have been detecting with high-contrast imaging in-struments such as SPHERE (cf. Janson et al. 2016; Garufi et al.2016). We note that because of the ADI processing this imagemay have been biased and is not a faithful representation of thetrue intensity and geometry. The residual signal northeast of theimage center is likely due to stellar residuals. However, this sig-nal is almost aligned with the near minor axis, so the possibilitythat it is real cannot be completely excluded. In fact, it couldbe the marginal detection of some material outward of the ringwith high scattering e ffi ciency. The brightness asymmetry be-tween the west and east disk wings is as pronounced as in thepolarimetric images, especially for reduction Y - to J -band. The disk is nicely resolved, confirming the diskgeometry and surface brightness extension stated above.
4. Analysis
We build a radiative transfer model for T Cha aiming to repro-duce the basic structure of its disk. We take earlier e ff orts (Olof-sson et al. 2011, 2013; Huélamo et al. 2015) as a starting pointfor independent three-dimensional (3D) radiative transfer calcu-lations using the Monte Carlo code RADMC-3D (Dullemondet al. 2012). We aim to complement the current understanding ofthe disk geometry by also taking into account our new SPHEREdata. RADMC-3D is used to calculate the thermal structure ofthe dust disk and ray-traced synthetic scattered light images inthe NIR. The polarization of RADMC-3D was investigated by The RADMC-3D source code and more details are available onlineat http: // / ∼ dullemond / software / radmc-3d / . Kataoka et al. (2015), who performed a benchmark test againstthe numerical models presented in Pinte et al. (2009).
The disk around T Cha is parametrized using constraints ob-tained from previous analyses of data sets (cf., Olofsson et al.2011, 2013; Huélamo et al. 2015) and from the new SPHEREobservations presented in this paper. We assume the disk to becomposed of two spatially separated zones with an inner (r in ) andouter radius (r out ) each, a narrow inner disk close to the centralstar that is responsible for the NIR excess and a more extendedouter disk. Inner and outer disks are assumed to be coplanar,since there is no significant evidence for a misaligned inner disk,which would cast shadows onto the outer disk (cf. the cases ofHD142527, Marino et al. 2015 and HD100453, Benisty et al.2017). The surface density structure is defined by a power-lawprofile and an exponential taper at the outer edge of the outercomponent (e.g., Hughes et al. 2008), Σ ( r ) = Σ (cid:32) rr c (cid:33) − δ exp − (cid:32) rr c (cid:33) − δ , (4)where r c corresponds to a characteristic radius and δ denotes thesurface density index. For the sake of simplicity, we assume auniform distribution along the azimuth in our model and con-centrate on the radial disk structure. The disk scale height is pa-rameterized radially as H ( r ) = H ( r / r ) β , where H is the scaleheight at a reference radius r and β is the flaring index. The ver-tical density distribution follows a Gaussian profile, so that thedust volume density is given by ρ ( R , ϕ, z ) = Σ ( R ) √ π H ( R ) exp (cid:32) − z H ( R ) (cid:33) , (5)where the spherical coordinates R and z can be converted intocylindrical coordinates with R = r sin( θ ) and z = r cos( θ ),where θ is the polar angle. We consider a power-law grain sizedistribution with an index p = − .
5, dn(a) ∝ a p da between aminimum (a min ) and maximum grain size (a max ). During ourmodeling process we use di ff erent values for the parameters a min and a max , where two distributions are eventually find to give anequally good match for the total intensity image (cf. Sect. 4.1.2).The dust is assumed to be a mixture made of silicates (Draine2003), carbon (Zubko et al. 1996), and water ice (Warren &Brandt 2008) with fractional abundances of 7%, 21%, and42%, consistent with Ricci et al. (2010). The remaining 30% isvacuum. The opacity of the mixture is determined by means ofthe Bruggeman mixing formula. The absorption and scatteringopacities, κ scat and κ abs , as well as the scattering matrix elements Z i j are calculated for spherical, compact dust grains with Mietheory considering the BHMIE code of Bohren & Hu ff man(1983).The radiative transfer calculations start with computing thedust temperature consistently by means of a thermal MonteCarlo simulation using 10 photon packages . Hence, an equi-librium dust temperature is calculated considering the star as thesource of luminosity. The main inputs for the radiative trans-fer modeling are the dust density structure from Eq. 5 and the Each single package actually represents many photons at once as-suming that these photons follow the same path.Article number, page 7 of 17 & A proofs: manuscript no. pohl_ms_final
Table 2.
Overview of the best RADMC-3D model parameters
Parameter Inner disk Outer diskr in [au] ∗ a out [au] ∗ a c [au] – 50 b M dust [M (cid:12) ] 2 · − a · − b δ / r / a / b β { a min ,a max } [ µ m] ∗ { } { }∼ ∼
10p -3.5 -3.5incl [deg] ∗
69 69PA [deg] ∗
114 114
Notes. δ denotes the exponent of the surface density power-law and β corresponds to the disk flaring index. For the radiation source we takethe following star parameters: T e ff = K , M = . M (cid:12) , R = . R (cid:12) ,where the star is assumed to be spherical. All parameters marked withan asterisk symbol ( ∗ ) were varied during the radiative transfer model-ing. The grain size distributions as well as inclination and PA were takento be the same for the inner and outer disk. References: a Olofsson et al.(2013); b Huélamo et al. (2015). dust opacities. Full non-isotropic scattering calculations are per-formed that take multiple scattering and polarization into ac-count. To compare with the observations, synthetic Stokes I in-tensity images, and Stokes Q and U polarized intensity imagesare produced at H -band (1.6 µ m) using 10 photon packages.These theoretical images are then convolved with a GaussianPSF with a FWHM of 0 (cid:48)(cid:48) .
04 assuming the object to be at 107 pc.Moreover, the synthetic total intensity images are run throughthe MPIA-PCA and SpeCal-PCA processing described in Sect.2.3 to have a proper comparison. The polarimetric Stokes Q and U images are eventually converted into their azimuthal coun-terparts Q φ and U φ . All synthetic images are normalized to thehighest disk surface brightness and displayed using the same dy-namical range as for the observational data. The coronagraphused in our IRDIS observations is mimicked by masking the in-ner 0 (cid:48)(cid:48) .
18 of the disk (19.3 au at 107 pc distance).
The inner disk geometry parameters are adopted from Olofssonet al. (2013) and are kept fixed in the modeling process.By adjusting the parameters from Huélamo et al. (2015) wegenerate the outer disk and run a grid of models exploring apre-defined parameter space for r in , r out , a min , a max , incl andPA. The fiducial model is defined by the set of parameters thatcauses a minimization in the residuals between observations andmodel within the paramater ranges set. For this determinationthe images of both, data and models, are normalized to thehighest flux value outside of the coronagraph. There is, however,no automatic fitting routine since computing tens of thousandsof 3D models for the T Cha system is computationally far tooexpensive. The best parameters are summarized in Table 2. Themodeling approach taking into account the new high-contrastSPHERE images allow us to better constrain the position of theinner rim of the outer disk, which we find to be at a significantlylarger radius of ∼
30 au ( ∼ (cid:48)(cid:48) .
28) compared to earlier work. Hence, the cavity size between the inner and outer disk iscorrespondingly larger than previously thought (Olofsson et al.2013). Furthermore, the polarimetric measurements provide usbetter estimates of the grain sizes.
Synthetic total intensity images
Figure 5, left panel, shows the synthetic Stokes I image at H -band from the first of our two radiative transfer models. It isproduced at a disk position angle of PA = ◦ and an inclina-tion angle of i = ◦ , which is similar to those values derivedin Huélamo et al. (2015). Our disk model gives a qualitativelygood match with the IRDIS total intensity images from Fig.3.The bright arc as the dominant source of scattered light is wellreproduced and corresponds to forward scattered light from thenear side of the inclined disk. The ADI images may, however,be significantly altered by the software processing, which wasalready shown by Garufi et al. (2016) for the case of HD100546.This ADI bias is especially important for T Cha, since theself-subtraction is strong due to the small field rotation and highinclination. Thus, we apply the ADI processing routines to themodel image. To do so we process the model image rotated by70 deg with the raw data considering the same PCA parameters.The middle and right panels of Fig. 5 show the resulting post-processed images depending on the PCA reduction method. TheADI procedure damps the signal of the backside of the disk andintroduces a brightness asymmetry along the disk surface. Thus,an original azimuthally symmetric feature can be seen as anasymmetric double-wing structure for a specific disk geometryand orientation. We note here that we additionally favor aphysical reason for this asymmetry, since this is also seen in thepolarimetric images (cf. Sect. 5.3). The ADI processed modelimage supports that the geometrical parameters used in ourmodel, in particular the gap size, reproduce the observationsnicely.The fraction of star light scattered o ff the disk surface layertowards the observer depends on the disk properties (e.g., massand scale height), but also on dust grain properties that determinethe phase function. Dust grains, which are large compared tothe wavelength, have strongly forward peaking scattering phasefunction, while small grains scatter photons almost isotropically.When keeping the minimum dust grain size fixed at 0.01 µ m, amaximum grain size of at least 100 µ m is requested to matchthe observations. This serves to reduce the influence of the smallgrains that are in the Rayleigh limit and absorb radiation muchmore e ffi ciently than they scatter it. Except for very turbulentdisks, one would, however, expect very large grains ( > µ m) tosettle below the scattering surface. The need for large grains inthe disk surface can be avoided by removing the smallest grains.Hence, the minimum and maximum values for the dust grain sizedistribution in our model are somehow degenerate. An equallygood image, that also achieves the desired brightness contrastof the arc with respect to the disk backside, is obtained by us-ing a narrow distribution around 10 µ m. Grains of about ten mi-crons in size are strong forward scatterers in the H -band. If evenlarger particles were primarily present, the forward scattering ef-ficiency would be too strong, and the brightness of the disk’s farsides would be too faint. The corresponding synthetic intensityimages for the second model and their appearance after the ADIpost processing with PCA can be found in Fig. 6. Article number, page 8 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. 5.
Synthetic total intensity images from our radiative transfer model a min = . µ m and a max = µ m. The left panel shows the theoretical Stokes I image convolved with a Gaussian PSF with FWHM of 0 (cid:48)(cid:48) .
04 (at107 pc distance). The middle and right panels show the theoretical model image at 70 ◦ processed together with the raw DBI data by the di ff erentPCA methods as described in Sect. 2.3. The central 0 (cid:48)(cid:48) .
18 of the image are masked to mimic the e ff ect of the coronagraph on the observations. Theunits are arbitrary, but the dynamical range of the color bar is taken the same as in Fig.3. Fig. 6.
Synthetic total intensity images from our radiative transfer model ∼ µ m.The layout and color scale is identical to Fig. 5. Fig. 7.
Phase functions Z , − Z and degree of polarization − Z / Z of the dust grains dependent on the scattering angle θ and calculated at λ = . µ m. Model Synthetic polarimetric images
Our results so far demonstrate that we find a quite good model tomatch the disk geometry of T Cha. The goal is, however, to alsoanalyze the grain properties compatible with the polarimetricdata. For scattering in the Rayleigh and Mie regime, that is, for grains with sizes smaller than or approximately equal thewavelength (2 π a (cid:46) λ ), maximum polarization is expectedalong a scattering angle of 90 deg. The phase function Z , thescattering matrix element − Z and the degree of polarization − Z / Z of the dust grains used in our radiative transfer models Article number, page 9 of 17 & A proofs: manuscript no. pohl_ms_final are shown as a function of the scattering angle θ in Fig. 7.Comparing those quantities for both models allows us to ruleout the first model covering a wide range of grain sizes from0.01 to 1000 µ m. This, rather, produces maxima in polarizedintensity along the semi-major axis (see Appendix A), whichis clearly not observed in the SPHERE PDI data from Fig. 1.Although one can recognize an extreme forward peak in − Z ,the resulting peak in polarized intensity is hidden behind thecoronagraph. For our second model with grains of ∼ µ m thephase function is also dominated by small-angle scattering asseen in the Z plot, but the − Z curve has a strong peak atsmall angles of ∼ ◦ . This leads to the spatial shift of brightnessmaxima away from the semi-major axis (i.e., scattering at90 deg), meaning that the maximum polarized intensity occursat the forward scattering position. This is in good agreementwith our polarimetric SPHERE observations.Figure 8 shows the synthetic Q φ and U φ images at H -bandfor the second model, with a disk position angle of PA = ◦ and an inclination of i = ◦ ; both determined from the fit to thetotal intensity image. The Q φ image is dominated by large posi-tive signal, which is consistent with forward scattering from theclose edge of the disk. The small-scale brightness blobs could bedue to self-scattering of thermal emission or the result of mul-tiple scattering treatment in the radiative transfer calculations.Monte Carlo noise can be ruled out as the source of these fea-tures since the best models were also run with a higher numberof photon packages (10 ) for testing, confirming that our calcu-lations are converged. Similar to the observed U φ image, the U φ model image shows an alternation of positive (white) and nega-tive (dark blue) signal, although the exact geometry appears dif-ferent. The extension of the south-east lobe with negative signalis comparable to that in the observational image in Figs. 1 (bot-tom left panel) and 2. The positive signal is a bit less pronouncedin our calculated model. Since the U φ signal in the observationalimage can be substantially influenced by noise, instrumental ef-fects, and the data reduction procedure, which is not included inour modeling, such a deviation was to be expected. The U φ / Q φ peak-to-peak value for the best model is about 15%, which is stillin very good agreement with the observations (9% and 14%), butlower than calculated in the study by Canovas et al. (2015) onnon-azimuthal linear polarization. For their models and in ourRADMC-3D calculations we consider a full treatment of polar-ized scattering o ff randomly oriented particles. Due to the ab-sence of any instrumental influence on the polarization, the U φ signal visible in the model images should be primarily connectedto multiple scattering events happening in the disk. However,the contribution of multiple scattering strongly depends on thedisk inclination, the grain population and the mass of the disk.In Canovas et al. (2015) the signal in U φ reaches up to 50% ofthe Q φ , but only for an inclination of 70 deg, a grain size dis-tribution with a min , max = (5 , µ m, and a disk significantlymore massive than assumed for our T Cha model. A higher diskmass produces more scattering events as simply more scatter-ing particles are available. Furthermore, the higher scattering ef-ficiency of the grains relative to their absorption e ffi ciency re-sults in stronger multiple scattering signature in the models ofCanovas et al. (2015). These e ff ects can explain the discrepancyto our U φ / Q φ peak-to-peak value of only 15%. East-west brightness asymmetry
The clear asymmetry in brightness along the inner edge of theouter disk from the observations is naturally not produced with
Table 3.
Astrometry and photometry relative to the star of the compan-ion candidate in the T Cha system
IRDIS companion candidateFilter
H2 H3 λ [ µ m] 1.593 1.667Contrast [mag] 11.65 ± ± ± ± ± ± ± ± ± ± Table 4.
Relative astrometry of the companion candidate for di ff erentinstrumental data NACO SPHEREDate 5 March 2004 30 May 2015JD 2453070 2457173Separation [mas] 3868.9 ± ± ± ± Notes.
The NACO data was published in Chauvin et al. (2010). our symmetric disk model with spatially invariant dust proper-ties. A slightly o ff set disk is one possibility for explaining theorigin of the asymmetry and we explore this scenario in the fol-lowing. We take our best axisymmetric model and slightly dis-place the star along the semi-major axis with respect to its origi-nal central position, while keeping the general disk structure un-changed. This is directly implemented into the radiative trans-fer code and not performed in a post-processing manner. A gridof additional models is computed, where the magnitude of thephysical o ff set between the center of the T Cha disk and the po-sition of its host star is changed between 0.5 and 2.5 au. We areonly interested whether such a scenario is principally reliable,so we abstain from a fitting procedure. A value of x = . x is measured along the disk’s semi-major axis, representsa reasonable match. The o ff set we apply is equivalent to a diskeccentricity of e ≈ .
07. This way a brightness contrast betweenthe east and west sides of 2 (Stokes I ) and 3 (Stokes Q φ ) canbe reached (see Fig. 9), consistent with the observational con-straints. Other possible scenarios for the brightness asymmetryare discussed in Sect. 5.3. One candidate companion (CC) is detected in the IRDIS field ofview (Fig. 10), whereas no point-like sources are found in theIFS image. The speckle pattern is reduced in each frame of thesequence by subtracting an optimized reference image calcu-lated by the TLOCI algorithm (Marois et al. 2010) implementedin SpeCal. We estimate the astrometry and photometry of thiscompanion candidate using the calibrated unsaturated PSF(Galicher & Marois 2011) to remove biases. First, we roughlyestimate the flux and position of the source in the TLOCI image.The SpeCal pipeline then creates a data cube of frames that onlycontain the unsaturated PSF at the candidate position on thedetector, accounting for the field-of-view rotation in each frame.The TLOCI coe ffi cients used to generate the TLOCI imagewhere the candidate is detected are applied on the candidate data Article number, page 10 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. 8.
Synthetic Q φ (left) and U φ (right) images at H -band. They are convolved with a Gaussian PSF with FWHM of 0 (cid:48)(cid:48) .
04 (at 107 pc distance).The color scale is arbitrary, the dynamical range is similar as in Fig.1. Negative values of U φ are saturated at dark blue color. Fig. 9.
Synthetic Stokes I (left) and Q φ (right) images at H -band of model ff set from its original central positionalong the semi-major axis. The images are convolved with a Gaussian PSF with FWHM of 0 (cid:48)(cid:48) .
04 (at 107 pc distance). The color scales are identicalto Figs. 5 and 8, respectively. cube. The resulting frames are rotated to align north up. Themedian of these frames provides the estimation of the candidateimage in the TLOCI image. We then adjust the estimated imagesubpixel position and its flux to minimize the integrated fluxof the di ff erence between the real and estimated candidateimages. We use a 3 × FWHM diameter disk for the minimization.The 1 σ error bars are the required excursions in position or in flux to increase the minimum residual flux by a factor of √ .
15 (cf. Galicher et al. 2016, Galicher et al., in prep.). Weempirically determine this factor running tests on sequences,in which we inject known fake planets. Using the calibratedunsaturated PSF, we also estimate the TLOCI throughput in allTLOCI sections following a procedure similar to the one usedfor the candidate position and flux estimation. The images were
Article number, page 11 of 17 & A proofs: manuscript no. pohl_ms_final
Fig. 10.
Signal-to-noise ratio map of the PCA reduction (
H2H3 data. The point source considered as a companion can-didate (CC) is marked with a circle.
Fig. 11.
Relative astrometry of the companion candidate labeled as ‘CC’in Fig. 10 measured in SPHERE, NACO and HST data. The black solidline displays the motion of the companion if co-moving and the blacksquares are the positions expected at the time of HST and NACO obser-vations. thus flux calibrated. The systematic errors for the astrometryof the detected companion candidate include the uncertaintieson the pixel scale, North angle, frame centering using thesatellite spots, accuracy of the IRDIS dithering procedure,anamorphic correction and SPHERE pupil o ff set angle inpupil-tracking mode (Vigan et al. 2016; Maire et al. 2016). Thecalibration uses pixel scales of (12 . ± . / pix and(12 . ± . / pix for the H2 and H3 filters, respectively,and a true North o ff set of ( − . ± . ◦ is considered (Maireet al. 2016). −4 −2 0 2H2−H3201510 M H M0−M5M6−M9L0−L5L6−L9T0−T5T6−T8HR8799bHR8799cHR8799dHR8799e
PZ Tel B1RXS1609b2M1207bCD−352722BHN Peg B UScoCTIO 108BGJ 758b >T8 cc Fig. 12.
Color-magnitude diagram displaying our candidate compan-ion, which is marked in red and labeled with CC, compared to knownsubstellar field (colored symbols) and young objects. Note that this plotassumes that CC is at the same distance as the star. Since CC is eventu-ally classified as a background object based on a common proper motiontest (cf. Fig. 11), it is likely located much further.
CC is located at a separation of (3 . ± . (cid:48)(cid:48) with contrast( ∆ m H = . ± .
04) mag (see Table 3). This same com-panion was already detected by Chauvin et al. (2010) with m K = (11 . ± .
1) mag and is also present in HST data takenin coronagraphic mode with STIS in March 2000. Combiningthe new position measured from SPHERE with the old data werule out this object as being gravitationally bound to T Cha,because the motion observed over these 15 years is too largeto be explained by a Keplerian orbit around this star; it istherefore a contaminant object. For completeness, given thatthe T Cha proper motion is µ α = ( − . ± .
2) mas / yr and µ δ = ( − . ± .
19) mas / yr (Gaia Collaboration et al. 2016),we also notice that CC has a high relative proper motion withrespect to a background object (Fig. 11). For completeness, weshow the CMD in Fig. 12. We note that this plot assumes thatCC is at the same distance as T Cha, since its actual distanceis unknown. This is rather unlikely based on our previousconclusion that it is not physically associated with T Cha. CC islikely located much further, and thus, likely intrinsically muchbrighter than an object at the L-T transition. Given the H2-H3color ∼
0, we conclude that this object could be either a floatingbrown dwarf or a low mass star of the galactic thick disk or halo.
The IRDIS detection limits for point sources are determinedusing the TLOCI data reduction. We estimate the 5 σ noise level,where σ is the azimuthal robust deviation of the residual flux inannuli of λ / D width rejecting pixels with no flux. Finally, the5 σ noise levels are divided by the stellar flux estimated fromthe unsaturated images. The maximum contrast reached withIFS is obtained by applying the PCA technique. The contrast Article number, page 12 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. 13.
Contrast curves and companion mass limits derived for IFS(black) after applying PCA, and for IRDIS H2 and H3 bands (red andblue, respectively) from the TLOCI reduction ( (cid:48)(cid:48) .
12. The detection limits from NACO K s band data aregiven for comparison (gray dashed line, inner cut at 0 (cid:48)(cid:48) . ff ected by the presence of the disk in scattered light. limits are estimated by an azimuthal standard deviation, that is,between pixels at the same separation from the star, for eachangular separation, corrected by the star flux (obtained fromthe o ff -axis PSF images taken immediately before and afterthe coronagraphic observations) and the algorithm throughput(using synthetic companions injected into the data before thedata processing as described above).In Fig. 13 the contrast curves obtained for the di ff erent datasets are shown. The IRDIS data give a 5 σ contrast for a sep-aration larger than 1.0 (cid:48)(cid:48) of greater than 12.5 mag and 12.6 magin the H2 and H3 bands, respectively. Compared with NACO K s band results (cf. Chauvin et al. 2010), these observationsare deeper by more than three magnitudes at a separation of0 (cid:48)(cid:48) .
7, that is, at the outer edge of the NACO coronagraph, whilethe contrast values at wider separations are comparable. IFSis deeper in contrast for separations closer than 0 (cid:48)(cid:48) . YJ band at a separation of ∼ (cid:48)(cid:48) .
7, assuming a gray contrast between the two objects.Using the theoretical atmospheric models AMES-COND(Allard et al. 2003) we convert the contrast limits into upper lim-its on the mass of possible objects orbiting around T Cha. Thesemodels are valid for T e ff < ∼ M jup in the innermost regions ( ∼ (cid:48)(cid:48) . (cid:48)(cid:48) . ∼ M jup for a separation between 0 (cid:48)(cid:48) . (cid:48)(cid:48) . Ournew SPHERE observations, therefore, improve the NACO masslimits especially up to ∼ (cid:48)(cid:48) . Both, the contrast and mass curveare cut at 0 (cid:48)(cid:48) .
12. The whole coronagraph system (apodizer, mask,stop) produces a radial transmission profile, which has not beenaccounted for in the derivation of the detection limits. The e ff ectis visible at the region near the edge of the mask plus λ/ D , thus,we exclude the inner 0 (cid:48)(cid:48) .
12. Furthermore, we note that the de-tection limits represent an average value around the star, whichmight be a ff ected by the disk signal at the location of the disk.However, we expect this e ff ect to be small given the rather com-pact nature of the disk around T Cha.
5. Discussion
Our analysis and modeling of the SPHERE data set confirmsthat the disk around T Cha consists of an inner disk part and anouter disk part, separated by a cavity. Compared to the previ-ous study by Olofsson et al. (2013) we find the small dust cav-ity size to be larger by a factor of ∼
2. Besides, it is even largerthan the mm dust cavity of 20 au estimated in Huélamo et al.(2015). This is rather unexpected, as the dust trapping scenariofor transition disks is supposed to work such that bigger dustis trapped at a ring located outside of the small dust / gas cavityedge (see e.g., Pinilla et al. 2012; van der Marel et al. 2015). Thispossible contradiction could be, however, due to uncertainties inthe model fitting of data with low resolution by Huélamo et al.(2015). An inclination angle of ∼ ◦ and PA of ∼ ◦ best matchour SPHERE observations, which is in agreement with Huélamoet al. (2015). Our new optical and near-infrared data do not, how-ever, help us to constrain the outer disk radius. In our radiativetransfer model we considered a tapered density profile for thedust density description of the outer disk, meaning that the sur-face density falls o ff gradually and hence, there is a smooth de-crease of the dust mass per radius bin. However, simultaneouslyreproducing the gas and dust components of the disk remainschallenging, and including this in our modeling e ff ort is beyondthe scope of this paper. To simultaneously match the total intensity and polarimetric im-ages obtained during our SPHERE observations, intermediatesized grains of ∼ µ m must be present in the disk. This pro-vides a better match with the observed properties of the diskthan dust distributions covering several orders of magnitudesin size or a narrow distribution peaking at (sub-)micron size.This is in accordance with current grain growth models pro-ducing systematically larger grains, although we cannot guar-antee that ∼ µ m grains are located at the upper surface layer.By means of scattered light observations in the NIR we onlytrace the disk surface where the micron-sized grains are locatedfor sure. Compact grains of a few tens of microns are expected Article number, page 13 of 17 & A proofs: manuscript no. pohl_ms_final to start settling down toward the disk midplane. The e ffi ciencyand timescale of vertical mixing depends, however, on the levelof turbulence in disks which is still uncertain. With strong tur-bulence (high α -viscosity, Shakura & Sunyaev 1973) all grainsizes are better mixed. Thus, even larger grains can be present inthe disk surface where they can contribute to the scattering. Fur-thermore, the amount of porosity of dust grains is unknown andstill debated (e.g., Ossenkopf 1993; Dominik & Tielens 1997;Kataoka et al. 2014). For fractal aggregates with high porositythe phase function is supposed to di ff er (Tazaki et al. 2016),which might alter our grain picture for T Cha. A larger poros-ity for the same grain size might reduce the settling, where thesize of the monomers still determines the absorption and scat-tering opacities. As shown by Min et al. (2012), the appearanceof a disk in scattered light could be di ff erent depending on thefraction of flu ff y aggregated dust particles compared to compactgrains contained in the disk. We also note that very large grains(mm size) are indeed also expected to be present in the midplanein order to match the (sub-)mm data (cf. Huélamo et al. 2015). The intensity and polarized intensity distributions observed forT Cha are asymmetric with respect to the minor axis of thedisk. Similar brightness variation has also been detected in otherdisks, such as RY Tau (Takami et al. 2013) and AK Sco (Jansonet al. 2016). In our radiative transfer modeling we explored theorigin of the east-west asymmetry seen along the semi-majoraxis in the SPHERE observations by looking into the simplestpossibility of a slightly o ff set disk. We approximate such aneccentric disk by calculating scattered light images of an az-imuthally symmetric disk, but introducing an o ff set between thedisk center and the star. A planetary companion on an eccentricorbit could shape the outer disk into an eccentric disk, causingthe o ff set. Keeping T Cha’s stellar properties as the photonsource in the radiative transfer code, but adding a positionalo ff set, already reproduces well the observed asymmetry.However, alternative explanations for the east-west bright-ness di ff erence cannot be ruled out, and several e ff ects may in-teract. Another idea is that an asymmetry in the inner disk orat the gap edges can lead to illumination e ff ects helping to ex-plain the dips in scattered light. The circumstellar disk aroundT Cha may be actually still in an early stage of planetary for-mation. Thus, a dense dust clump formed in the inner, densestparts of the disk, or an already formed yet undetected planetaryperturber below the detection limit, could cause this asymme-try. However, this scenario also raises the question of whethersuch an anisotropy is indeed stationary or moves with the localKeplerian velocity. A third scenario deals with spatially variantdust properties leading to a di ff erent scattering e ffi ciency, whichis especially related to grain size, structure, and composition. Apossibility would be that unequal dust grain size distributions arepresent in the east and west wings of the disk, whose origin, how-ever, remains unexplained. A fourth possible scenario leading toshadows in the outer disk is an inner disk significantly tilted withrespect to the outer disk’s plane. However, we find this scenariounlikely, since this arrangement would rather lead to relativelysharp, dark lanes, which are not apparent in our T Cha images.
6. Conclusions
We have carried out VLT / SPHERE optical and NIR ob-servations in polarimetric di ff erential imaging mode with SPHERE / ZIMPOL in
VBB and SPHERE / IRDIS in H -bandof the evolved transition disk around the T Tauri star T Cha.Alongside the polarimetric observations, intensity images fromIRDIS H2H3 dual-band imaging with simultaneous spectro-imaging with IFS in YJ -band were obtained. The disk is clearlydetected in all data sets presented in this work and resolvedin scattered light with high angular resolution, allowing us toreview the current understanding of the disk morphology andsurface brightness. The basic structure of a classical transitiondisk previously reported by interferometric and (sub-)mmstudies, has been confirmed. We developed a radiative transfermodel of the disk including a truncated power-law surface den-sity profile. The conclusions of this paper are summarized below.1. Our RADMC-3D radiative transfer model with updateddisk parameters accounts well for the main geometry of thedisk, the cavity, and the outer disk with its bright inner rimlocated at 0 (cid:48)(cid:48) .
28 ( ∼
30 au). This is significantly further outthan previously estimated. A disk inclination of ∼ ◦ and aposition angle of ∼ ◦ matches the SPHERE data sets best.2. We confirm that the dominant source of emission is forwardscattered light from the near edge of the disk, given thehigh disk inclination. While small grains in the Rayleighlimit scatter photons rather isotropically and absorb verye ffi ciently, large dust grains with sizes (2 π a > λ ) havestrong forward scattering properties. This demands a certainrange of grain sizes to be present in the disk. We foundthat a power-law distribution with a min = . µ m and a max = µ m reproduces the total intensity observa-tions well, but fails to be consistent with the polarimetricimages. Thus, we propose a dominant grain size in thedisk of ∼ µ m. Such grains bring the desired amount offorward scattering and lead to a model that is in accordancewith the complete SPHERE data set presented. We notethat we restricted ourselves to the analysis of Mie theoryand spherical compact grains. However, for asphericalaggregates with high porosity the phase function is sup-posed to di ff er, which might alter our grain picture for T Cha.3. Our highly inclined disk model shows a significant U φ signalat H -band, which is in accordance with the observational U φ / Q φ peak-to-peak value of 14% and theoretical studies onmultiple scattering events. The exact geometrical U φ patternobserved with IRDIS is not reproduced, but the alternatingstructure of positive and negative lobes is well recognizable.4. The brightness asymmetry between the east and west sidescan be reproduced with a slight o ff set of the star’s position,representing a disk eccentricity of e ≈ .
07. A planetarycompanion on an eccentric orbit could force the outer diskto become eccentric, causing this o ff set. However, a locallydi ff erent grain size distribution and therefore a changeof the scattering properties, or illumination e ff ects due toasymmetric structures in the inner disk could also contributeto the brightness contrast observed.5. A previously known companion candidate is detected in theIRDIS field of view at a separation of (3 . ± . (cid:48)(cid:48) withcontrast m H2 = (11 . ± .
04) mag. We, however, rule outthe possibility that this object is bound and, thus, concludethat it is not part of the T Cha system.
Article number, page 14 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
6. Our analysis rules out the presence of a companion withmass larger than ∼ M jup between 0 (cid:48)(cid:48) . (cid:48)(cid:48) . ∼ M jup for wider separations.There could still be lower-mass planets in the outer diskregions and / or planets in the very inner disk. Acknowledgements.
We would like to thank the ESO Paranal Sta ff for their sup-port during the observations. We are very grateful to C.P. Dullemond for insight-ful discussions. A. P. is a member of the International Max Planck ResearchSchool for Astronomy and Cosmic Physics at Heidelberg University, IMPRS-HD, Germany. INAF-Osservatorio Astronomico di Padova acknowledges sup-port from the "Progetti Premiali" funding scheme of the Italian Ministry ofEducation, University, and Research. M. L., M. B., F. M. and C. P. acknowl-edge funding from ANR of France under contract number ANR-16-CE31-0013.J. O. acknowledges support from ALMA / Conicyt Project 31130027, and fromthe Millennium Nucleus RC130007 (Chilean Ministry of Economy). SPHEREis an instrument designed and built by a consortium consisting of IPAG (Greno-ble, France), MPIA (Heidelberg, Germany), LAM (Marseille, France), LESIA(Paris, France), Laboratoire Lagrange (Nice, France), INAF-Osservatorio diPadova (Italy), Observatoire de Genève (Switzerland), ETH Zurich (Switzer-land), NOVA (Netherlands), ONERA (France) and ASTRON (Netherlands), incollaboration with ESO. SPHERE was funded by ESO, with additional contri-butions from CNRS (France), MPIA (Germany), INAF (Italy), FINES (Switzer-land) and NOVA (Netherlands). SPHERE also received funding from the Eu-ropean Commission Sixth and Seventh Framework Programmes as part of theOptical Infrared Coordination Network for Astronomy (OPTICON) under grantnumber RII3-Ct-2004-001566 for FP6 (2004-2008), grant number 226604 forFP7 (2009-2012) and grant number 312430 for FP7 (2013-2016).
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Appendix A: Synthetic Q φ image of model Article number, page 16 of 17. Pohl et al.: VLT / SPHERE observations of the disk around T Cha
Fig. A.1. Q φ (left) and U φ (middle) images of model H -band. The right panel considers the same model, but the star is slightly o ff set fromits original central position along the semi-major axis. The images are convolved with a Gaussian PSF with FWHM of 0 (cid:48)(cid:48) ..