The protoplanetary disk of FT Tauri: multi-wavelength data analysis and modeling
Antonio Garufi, Linda Podio, Inga Kamp, François Ménard, Sean Brittain, Carlos Eiroa, Benjamin Montesinos, Míguel Alonso-Martinez, Wing-Fai Thi, Peter Woitke
AAstronomy & Astrophysics manuscript no. aa c (cid:13)
ESO 2018September 20, 2018
The protoplanetary disk of FT Tauri:multi-wavelength data analysis and modeling (cid:63)
A. Garufi , L. Podio , , I. Kamp , F. M´enard , , S. Brittain , C. Eiroa , B. Montesinos ,M. Alonso-Mart´ınez , W.F. Thi , and P. Woitke , , Institute for Astronomy, ETH Z¨urich, Wolfgang-Pauli-Strasse 27, CH-8093 Zurich, Switzerlande-mail: [email protected] CNRS / UJF Grenoble 1, UMR 5274, Institut de Plan´etologie et d’Astrophysique de Grenoble (IPAG), France INAF - Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125, Florence, Italy Kapteyn Astronomical Institute, Postbus 800, 9700 AV Groningen, The Netherlands UMI-FCA, CNRS / INSU, France (UMI 3386), and Dept. de Astronom´ıa, Universidad de Chile, Santiago, Chile Department of Physics & Astronomy, 118 Kinard Laboratory, Clemson University, Clemson, SC 29634, USA Dpt. F´ısica Te´orica, Facultad de Ciencias, Universidad Aut´onoma de Madrid, Cantoblanco, 28049 Madrid, Spain Dpt. de Astrof´ısica, Centro de Astrobiolog´ıa, ESAC Campus, P.O. Box 78, E-28691 Villanueva de la Ca˜nada, Madrid, Spain University of Vienna, Dept. of Astronomy, Turkenschanzstr. 17, A-1180 Vienna, Austria UK Astronomy Technology Centre, Royal Observatory, Edinburgh, Blackford Hill, Edinburgh EH9 3HJ, UK SUPA, School of Physics & Astronomy, University of St. Andrews, North Haugh, St. Andrews KY16 9SS, UKReceived ....; accepted ....
ABSTRACT
Context.
Investigating the evolution of protoplanetary disks is crucial for our understanding of star and planet formation. Severaltheoretical and observational studies have been performed in the last decades to advance this knowledge. The launch of satellitesoperating at infrared wavelengths, such as the
Spitzer
Space Telescope and the
Herschel
Space Observatory, has provided importanttools for the investigation of the properties of circumstellar disks.
Aims.
FT Tauri is a young star in the Taurus star forming region that was included in a number of spectroscopic and photometricsurveys. We investigate the properties of the star, the circumstellar disk, and the accretion / ejection processes and propose a consistentgas and dust model also as a reference for future observational studies. Methods.
We performed a multi-wavelength data analysis to derive the basic stellar and disk properties, as well as mass accre-tion / outflow rate from TNG / DOLoRes, WHT / LIRIS, NOT / NOTCam, Keck / NIRSpec, and
Herschel / PACS spectra. From the liter-ature, we compiled a complete Spectral Energy Distribution. We then performed detailed disk modeling using the
MCFOST and
ProDiMo codes. Multi-wavelength spectroscopic and photometric measurements were compared with the reddened predictions of thecodes in order to constrain the disk properties.
Results.
We determine the stellar mass ( ∼ . (cid:12) ), luminosity ( ∼ .
35 L (cid:12) ) and age ( ∼ . ∼ · − M (cid:12) / yr) to be within the range of accreting objects in Taurus.The evolutionary state and the geometric properties of the disk are also constrained. The radial extent (0.05 to 200 AU), flaring angle(power-law with exponent = . .
02 M (cid:12) ) of the circumstellar disk are typical of a young primordial disk. This objectcan serve as a benchmark for primordial disks with significant mass accretion rate, high gas content and typical size.
Key words. stars: pre-main sequence – planetary systems: protoplanetary disks – accretion, accretion disks – ISM: individual object:FT Tau
1. Introduction
Protoplanetary disks are the birthplaces of planets and the studyof their physical and chemical structure can help us understandplanet formation. By studying a large sample of young proto-planetary disks (class II) in detail, we may be able to assess thevariety in disk structure and to match that to the ever growingdiversity in exoplanetary systems architecture.The disk evolution can be observationally constrained bystudying the Spectral Energy Distribution (SED) of YoungStellar Objects (YSOs) in di ff erent evolutionary stages. In par-ticular, the mass accretion rate can be estimated from the excessin the UV and optical spectra and from emission lines that are (cid:63) Based on
Herschel data.
Herschel is an ESA space observatory withscience instruments provided by European-led Principal Investigatorconsortia and with important participation from NASA. thought to form in the magnetospheric accretion process (Basri& Bertout 1989, Edwards et al. 1994, Hartmann, Hewett, &Calvet 1994).Also geometrical properties of circumstellar disks changewith time. In the absence of spatially resolved images, the ge-ometry has to be constrained from infrared (IR) and millimetricphotometry. A flaring geometry is a natural explanation for thestrong far-infrared (FIR) flux shown by most sources (Kenyon& Hartmann 1987). Grain-grain collisions result in dust graingrowth and, on timescales of 10 − years, grains are thoughtto settle to the mid-plane leaving the gas at the disk surface ex-posed to direct stellar radiation (Bouwman et al. 2008). This evo-lution is reflected in decreasing mid-infrared (MIR) fluxes and ina widening and flattening of the 10 µ m and 18 µ m silicate fea-tures (e.g. Furlan et al. 2006, Fang et al. 2009). a r X i v : . [ a s t r o - ph . S R ] M a y arufi et al.: The protoplanetary disk of FT Tauri The diagnostic of gas emission lines from the disk is anotherpivotal tool for the study of the disk structure. CO and OH linesare commonly detected in the near-infrared (NIR) spectra of pro-toplanetary disks. In particular, the CO fundamental ( ν = − i ] 63 µ m line (Dentet al. 2013) is believed to mostly originate in the colder outerregions, between 30 and 100 AU (Kamp et al. 2010). Its flux canbe used in combination with other lines as an indicator of diskgas mass.An ever growing number of datasets is becoming availablefor circumstellar disks. Nevertheless, many studies still focus onthe interpretation of a single dataset even in the framework ofdetailed disk modeling. However, full characterization of stellarand circumstellar properties of single objects is often reachedonly by employing dust and gas diagnostic measurements (pho-tometry and line emission) that cover the entire extent of a pro-toplanetary disk. In this paper, we aim to investigate the geo-metrical and chemical properties of the disk around the T TauriStar (TTS) FT Tauri, by means of a multi-wavelength datasetand consistent dust and gas modeling.Even though the Taurus star forming region and its membersare well studied (e.g. Kenyon et al. 1994, Gullbring et al. 1998,Luhman et al. 2010, Rebull et al. 2010), a comprehensive char-acterization has been restricted so far to either extremely bright(large) disks (e.g. DM Tau, Guilloteau & Dutrey 1994) or ex-ceptional objects (e.g. LkCa15, van Zadelho ff et al. 2001). FTTau is located in the South of the Barnard 215 dark cloud and issurrounded by extended emission (see Sloan Digital Sky Surveyoptical image from Finkbeiner et al. 2004). This source is farfrom most of the known Taurus members (see extinction map ofTaurus from Dobashi et al. 2005). No X-ray emission was de-tected at the optical position of the star (Neuh¨auser et al. 1995).FT Tau was included in a number of photometric and spectro-scopic surveys at di ff erent wavelengths. The main stellar anddisk properties from the literature are shown in Table 1. Withmore sensitive interferometers such as SMA and IRAM / PdBI,FT Tau can serve as an excellent target for more detailed astro-chemical studies.The multi-wavelength data used in this work and its data re-duction is presented in Sect. 2 and a detailed analysis thereofin Sect. 3. The result of this analysis are then further used inSect. 4 to build consistent dust and gas models (MCFOST and
ProDiMo ) for FT Tau. Sect. 5 then discusses the results and thesource variability.
2. Observations and data reduction
The data analysed in this paper consist of spectroscopy at opti-cal, NIR, and FIR wavelengths from the Telescopio NazionaleGalileo (TNG), the William Herschel Telescope (WHT), theNordic Optical Telescope (NOT), the Keck Observatory, and the
Herschel
Space Observatory. The instrumental settings for thespectroscopic observations are presented in detail in the follow-ing sections and summarized in Table 2. Additional data wereretrieved from the literature and consist of MIR spectroscopyfrom
Spitzer
Space Telescope and of photometry from optical toradio wavelengths (see Table 3 and Fig. 1).
Table 1.
Properties of FT Tauri estimated in previous works.
Coordinates (J2000):Right ascension α h m s .19 a Declination δ + ◦ (cid:48) (cid:48)(cid:48) .11 a Proper motion:Right ascension µ α + ± − Declination µ δ -15.3 ± − Stellar properties:Visual extinction A V c Spectral type M3e c Luminosity L ∗ (cid:12) c Disk properties:Gas mass M d . ± .
03 M (cid:12) d Outer radius R out + − AU da Cutri et al. 2003; b Luhman et al. 2009; c Rebull et al. 2010; d Andrews & Williams 2007.
We present spectroscopic data of FT Tau obtained using theDOLORES spectrograph, the Device Optimized for the LOwRESolution (Oliva 2004) mounted on the TNG (La PalmaObservatory). The observations were performed in November2009 with the VHR-V grism ( λ/ ∆ λ = (cid:39) A J-band (1.18 - 1.40 µ m) spectrum was taken in December2009 with the LIRIS, Long-slit Intermediate Resolution InfraredSpectrograph, at the WHT (La Palma Observatory). The spec-trum was acquired with a 0.75 (cid:48)(cid:48) slit width and the LIRIS hrj grismproviding a spectral resolution of λ/ ∆ λ = A K-band spectrum of FT Tau was taken with the NOTCAM,the Nordic Optical Telescope CAMera (Aspin 1999) of the NOT(La Palma Observatory). These observations were acquired inDecember 2009 by using the K grism ( λ/ ∆ λ = µ m), andnodding the slit between two positions. Table 2.
Instrumental settings for the spectroscopic observations of FT Tau.
Observation date Instrument Slit width (”) Spectral range ( µ m) Spectral resolution (km / s) Integration time (s)2009-11-26 TNG / DOLORES 1 0.475 - 0.670 196 1400, 600, 2002009-12-04 NOT / NOTCAM 1 1.95 - 2.37 200 482009-12-09 WHT / LIRIS 0.75 1.17 - 1.35 135 1202008-12-10 Keck / NIRSPEC 0.43 4.42 - 5.53 12 960, 12002010-03-26
Herschel / PACS Integral field 62.95 - 63.40 88 1152
The observation was obtained in seeing-limited conditions(FWHM (cid:39) S photometry from 2MASS (see Table 3 andFig. 1c). An M-band high-resolution spectrum of FT Tau was taken withthe NIRSPEC, the NIR SPECtrograph (McLean et al. 1998)on the W.M. Keck Observatory. The spectra were obtained inDecember 2008 using the M-Wide filter and 0.43” slit providinga resolution of λ/ ∆ λ = , µ m. FT Tau was observed for 16 and 20 minutes in successiveexposures.Because of the thermal background in the M-band, the datawere observed in an ABBA sequence where the telescope wasnodded 12” between the A and B positions. Each frame was flatfielded and scrubbed for hot pixels and cosmic ray hits. The ob-servations were combined as (A-B-B + A) / ∼ / s). Flux calibration was performed us-ing the available Spitzer / IRAC 4.5 and 5.8 µ m photometry (seeTable 3 and Fig. 1d). In order to obtain the gas velocity with re-spect to the star, we corrected for the observed radial velocity ofFT Tau, as estimated by Guilloteau et al. (2013) (v LSR = − / s). FIR spectroscopic observations of FT Tau were obtained withthe integral-field spectrometer PACS (Poglitsch et al. 2010), onboard of the
Herschel
Space Telescope (Pilbratt et al. 2010)as part of the
Herschel
Open Time Key Project GASPS (GASin Protoplanetary Systems, PI: W. Dent, see Dent et al. 2013).The observations were carried out in the chop- nod mode to re-move the background emission and with a single pointing onthe source. They cover simultaneously a selected wavelengthrange in the blue and in the red arms. In particular, the obser- vation acquired in line mode (OBSID: 1342192790) covers theranges 63.0–63.4 µ m and 180.7–190.3 µ m, with a resolution of88 km s − and 200 km s − . The observation acquired in rangemode (OBSID: 1342243501) covers the ranges 71.8–73.3 µ mand 143.5–146.6 µ m, with a spectral resolution of 162 km s − and 258 km s − .Data were reduced using HIPE 10. Removal of saturatedand bad pixels, chop subtraction, flat-field correction, and meanof two nods were performed by means of the available PACSpipeline.Photometric FIR observations of FT Tau were also takenwith Herschel / PACS within the GASPS project. The obtainedphotometric measurements at 70, 100, and 160 µ m are presentedin Howard et al. (2013) and shown in Table 3. We collected from the literature 44 photometric measurementsof FT Tau, from 0.36 µ m to 7 mm (see Table 3 for the fluxesand references). These data have been acquired with 14 di ff erentinstruments and over more than two decades (see Sect. 5.1 andAppendix A for a further discussion on variability).A MIR spectrum of FT Tau (spectral range 5.13 µ m - 39.90 µ m) was also retrieved from the literature (Furlan et al. 2006).
3. Results from observations
The spectra obtained by reducing the observations presented inSect. 2 are shown in Fig. 1. The optical TNG spectrum (Fig. 1a)shows a number of molecular absorption bands, that are typicalof late-type stars, and several emission lines, which are thoughtto originate from the accretion columns or from the outflow. TheJ-band WHT spectrum (Fig. 1b) and the K-band NOT spectrum(Fig. 1c) show prominent Pa β and Br γ emission lines, producedin the accretion process. In the Keck spectra (Fig. 1d, 1d1, 1d2)we detected the Pf β and Hu (cid:15) recombination lines, and CO ro-vibrational lines which are thought to be produced in the disk bythermal excitation or by UV fluorescence. The spectroscopic ob-servations collected with Herschel / PACS cover a number of disktracers, e.g. water lines (o-H O 7 -6 , p-H O 4 -3 ), high-Jrotational CO lines (CO J = = + J = i ] P - P and [O i ] P - P lines at 63.184and 145.525 µ m. However, only the [O i ] 63.184 µ m line wasdetected in the central spatial pixel (spaxel) of the PACS in-tegral field unit (Fig. 1f), while the other lines remained unde-tected. For those lines we report the 3 σ upper limit in Table 4.The Spitzer / IRS spectrum shows prominent silicate features at10 and 18 µ m (Fig. 1e), which are believed to originate in theoptically thin disk surface layer. The properties of all detectedlines are listed in Table 4. Wavelength (μm) x -11 x -10 x -9 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) Hβ FeII HeI NaD [OI] Hα HeIFeII (a)
Wavelength (μm) x -10 x -10 x -10 x -10 x -10 x -10 x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) Paβ(b)
Wavelength (μm) x -10 x -10 x -10 x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) Brγ (c)
Wavelength (μm) x -11 x -10 x -10 x -10 x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) PfβHuε (d)
Wavelength (μm) x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) Silicate features (e)
63 63.1 63.2 63.3 63.4
Wavelength (μm) x -11 x -11 x -11 x -11 x -11 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) [OI] 63μm (f) Wavelength (μm) x -11 x -10 x -10 x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) - - - - - - - - - - - - - - (d1) Wavelength (μm) x -11 x -10 x -10 x -10 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) - - - - - - - - - - (d2) Fig. 1.
Spectroscopic observations. (a) TNG optical spectrum; (b) WHT J-band spectrum; (c) NOT K-band spectrum; (d) Keckhigh-resolution spectra; (e) Spitzer MIR spectrum; (f)
Herschel
FIR spectrum of the only detected line ([O i ] 63 µ m); (d1) and(d2) zoom on the CO lines detected in the Keck spectra. Photometric non-simultaneous observations are plotted as color symbols.Spectra shown in (a), (b), (c) were flux-calibrated by means of the photometry. Table 3.
Photometric measurements of FT Tau.
Instrument λ Band Flux λ F λ ( µ m) (publication units) (erg · cm − · s − )SDSS 0.36 u band 16.681 ± a (6 . ± . · − USNO 0.44 B band 15.48 mag b . · − SDSS 0.48 g band 15.804 ± a (1 . ± . · − USNO 0.55 V band 14.69 mag b . · − SDSS 0.62 r band 14.254 ± a (3 . ± . · − USNO 0.68 R band 12.06 mag b . · − SDSS 0.76 i band 15.023 ± a (6 . ± . · − USNO 0.80 I band 11.36 mag b . · − SDSS 0.90 z band 12.456 ± a (1 . ± . · − ± d (3 . ± . · − ± d (4 . ± . · − S band 8.59 ± d (3 . ± . · − WISE 3.37 7.75 ± e (2 . ± . · − Spitzer / IRAC 3.6 7.64 ± f (2 . ± . · − Spitzer / IRAC 3.6 7.89 ± f (1 . ± . · − Spitzer / IRAC 4.5 7.12 ± f (1 . ± . · − Spitzer / IRAC 4.5 7.44 ± f (1 . ± . · − WISE 4.61 7.10 ± e (1 . ± . · − Spitzer / IRAC 5.8 6.81 ± f (1 . ± . · − Spitzer / IRAC 5.8 7.12 ± f (7 . ± . · − Spitzer / IRAC 8.0 5.95 ± f (1 . ± . · − Spitzer / IRAC 8.0 6.27 ± f (7 . ± . · − IRAS 12 0.46 Jy ± g (1 . ± . · − WISE 12.08 5.09 ± e (7 . ± . · − WISE 22.19 3.08 ± e (6 . ± . · − Spitzer / MIPS 24 3.15 ± f (4 . ± . · − IRAS 25 0.65 Jy ± g (7 . ± . · − IRAS 60 0.86 Jy ± g (4 . ± . · − Spitzer / MIPS 70 0.28 ± h (2 . ± . · − Herschel / PACS 70 0.73 ± i (3 . ± . · − IRAS 100 1.92 Jy ± g (5 . ± . · − Herschel / PACS 100 0.95 ± i (2 . ± . · − Herschel / PACS 160 1.27 ± i (2 . ± . · − CSO 350 1106 ±
82 mJy j (9 . ± . · − JCMT 450 437 ±
56 mJy j (2 . ± . · − CSO 624 260 ±
100 mJy k (1 . ± . · − CSO 769 250 ±
50 mJy k (9 . ± . · − JCMT 850 121 ± j (4 . ± . · − SMA 880 111 ± l (3 . ± . · − CSO 1056 137 ±
40 mJy k (3 . ± . · − IRAM 1300 130 ±
14 mJy m (3 . ± . · − IRAM 2700 25 ± n (2 . ± . · − VLA 7000 1.62 ± o (6 . ± . · − References: a Finkbeiner et al. 2004; b Monet et al. (2003); d Cutri et al. (2003); e Wright et al. 2010; f Luhman et al. (2010), two observations perwavelength; g Beichman et al. (1988); h Rebull et al. (2010); i Howard et al. (2013); j Andrews & Williams (2005); k Beckwith & Sargent (1991); l Andrews & Williams (2007); m Beckwith et al. (1990); n Dutrey et al. (1996); o Rodmann et al. (2006).
In this section we describe the methods applied to derivethe stellar properties (Sect. 3.1), a few disk properties (Sect. 3.2),and the mass accretion and outflow rates (Sect. 3.3).
First, we determined the spectral type of FT Tau by comparingthe optical TNG spectrum with the spectra of the MILES stel-lar libraries (Sanchez-Blazquez et al. 2006). The observed ab-sorption features suggest that FT Tau is an M2 or M3 star, inagreement with the result by Rebull et al. (2010). The tempera-ture scale is based on Cohen & Kuhi (1979). Then, we adopted a PHOENIX model ( T e ff = g ) = .
5, [Fe / H] = . ff er from strongextinction by the dust along the line of sight, either foregroundor in the disk. We estimated the visual extinction, A V , by com-paring the colors of the adopted PHOENIX model spectrum withthe available photometry. Since the UV excess of young accret-ing stars may extend up to red optical wavelengths and the IRexcess may start at ∼ µ m, we used (J-H) colors. Using the ex-tinction law by Cardelli et al. (1989) and R V = .
1, we obtained
Table 4.
Emission lines detected in our spectra and respective vacuum wavelengths. When the line presents instrumental gaps, lowerand higher estimates due to the missing region are included in the error. Dereddened fluxes are omitted when the correction for theextinction is negligible.
Line Wavelength Instrument Observed flux Dereddened flux ( A V = .
8) Likely origin( µ m) (10 − erg · s − · cm − ) (10 − erg · s − · cm − )H β . ± . ii . ± .
03 4.9 AccretionFe ii . ± .
03 5.6 AccretionFe ii . ± .
07 7.1 AccretionHe i . ± .
31 12.1 AccretionNaD 0.589 TNG 0 . ± .
03 4.8 AccretionNaD 0.590 TNG 0 . ± .
03 2.7 Accretion[O i ] 0.630 TNG 0 . ± .
21 2.6 Disk / OutflowH α . ± . i . ± .
05 6.9 AccretionPa β . ± . γ . ± . β . ± .
21 - AccretionCO 1 − . ± .
11 - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 1 − . + . − . - DiskCO 2 − . ± .
13 - DiskCO 1 − . + . − . - DiskCO 1 − . ± .
12 - DiskCO 2 − . ± .
10 - DiskCO 1 − . ± .
13 - DiskCO 2 − . ± .
07 - DiskCO 1 − . + . − . - DiskCO 1 − . ± .
19 - DiskCO 2 − . ± .
14 - DiskCO 1 − . ± .
26 - DiskCO 1 − . ± .
21 - DiskCO 2 − . ± .
10 - DiskCO 1 − . ± .
08 - DiskCO 2 − . ± .
07 - DiskCO 1 − . ± .
13 - DiskCO 1 − . ± .
14 - DiskCO 1 − . ± .
18 - DiskCO 1 − . ± .
09 - Disk[O i ] 63.185 Herschel . ± . / Outflowo − H O 71.947
Herschel < .
41 - -CH + − Herschel < .
37 - -CO 36 −
35 72.843
Herschel < .
97 - -p − H O 144.518
Herschel < .
32 - -CO 18 −
17 144.784
Herschel < .
33 - -[O i ] 145.525 Herschel < .
31 - - an extinction of A V = .
8. For a discussion on the uncertaintiesa ff ecting the determination of the extinction, see Appendix A.To determine the stellar luminosity and radius we imposedthe reddened PHOENIX flux at 1.25 µ m to be equal to the ob-served one, as available from 2MASS (see Table 3), assumingthat the latter is only due to photospheric emission.The estimated radius is R ∗ = . (cid:12) and the luminosity L ∗ = .
35 L (cid:12) . Uncertainties on those estimates are due to the distance(140 ±
10 pc, Kenyon et al. 1994) and e ff ective temperature (3%in the M2-M4 range, Kenyon & Hartmann 1995). Finally, using pre-main sequence (PMS) evolutionary tracksby Siess et al. (2000) and assuming [Fe / H] = M ∗ = . (cid:12) and age 1 . · yr).The derived stellar properties (see Table 5) are typical for TTSs(Beckwith et al. 1990, Kenyon & Hartmann 1995, Hartigan et al.1995). Wavelength (μm) x -15 x -14 x -13 x -12 x -11 x -10 x -9 λ · F ( λ ) ( e r g · s ⁻ ¹ · c m ⁻ ²) PhotometryTNGWHTNOTKeckSpitzer/IRSHerschel/PACSReddened photospheric model
Fig. 2.
SED of FT Tau. All the available photometric and spectroscopic non-simultaneous observations are plotted. The reddened
Phoenix model is reproducing the stellar photospheric emission (grey line). The SED clearly shows UV / optical excess emission atwavelengths shorter than ∼ µ m and infrared excess beyond ∼ µ m. IRAS photometry is excluded and the higher IRAC 5.8 and8 µ m photometry is omitted giving more weight to the IRS spectrum. Error bars are not visible at this scale. Table 5.
Estimated stellar properties.
Visual extinction A V ± ± ff ective temperature T e ff ±
200 KLuminosity L ∗ . ± .
09 L (cid:12)
Radius R ∗ . ± . (cid:12) Mass M ∗ . ± . (cid:12) Age 1 . ± . FT Tau shows a very prominent IR excess (see Fig. 2). By inte-grating the excess emission over the photosphere, we obtainedan IR excess luminosity L IR = . ± .
01 L (cid:12) .We also measured the IR excess at di ff erent wavelengths (seeTable 6) by subtracting the photospheric model from the ob-served photometry. This provides qualitative information on thegeometry of the disk. As the dust grows and settles toward themid-plane, the vertical scale height of the disk decreases, caus-ing less reprocessing of the stellar radiation and, thus, a smallerMIR excess. According to the evolutionary scheme by Fang etal. (2009), the excess shown by FT Tau is typical of objects withdisks that are evolving from a mildly flaring to a flat geometry.The MIR spectrum of FT Tau (see Fig. 1e) shows prominent,narrow and smooth silicate features. These are thought to origi-nate in the warm, optically thin disk surface and provide infor-mation on the silicate dust in this layer. As shown by Bouwmanet al. (2001) for Herbig Ae / Be stars, silicate features peaking at ∼ µ m, as in the case of FT Tau, are indicative of a dust popu-lation dominated by grains as small as 0.1 µ m. The flattening ofthese features (see Furlan et al. 2006 for a large sample of TTSs)can be a tracer of the evolution of the dust population at the disksurface. Some processes such as stellar winds and radiation pres-sure can deplete sub- µ m size grains (Olofsson et al. 2009). Thenarrow and prominent nature of the silicate features shown byFT Tau suggests that these processes are not yet e ffi cient in thisdisk. The 10 µ m feature can also provide insight into the crys-tallinity of the silicate (e.g. Sargent et al. 2006). The absence ofsubstructure in the MIR spectrum of FT Tau indicates that thesilicates are mostly amorphous.Finally, the MIR spectrum does not show PolycyclicAromatic Hydrocarbon (PAH) emission features. PAH emissionis indeed hardly detected in TTSs (Furlan et al. 2006), while it iscommon in more massive Herbig Ae / Be stars (see e.g. Meeus etal. 2001). This can be explained either in terms of di ff erent graincomposition or the weaker UV radiation field of low-mass stars. Most of the CO ro-vibrational lines detected withKeck / NIRSPEC are from transitions from the first vibra-tional level ( ν = −
0) and their fluxes are typically a factor of afew higher than those from ν = − ν = − =
12) lines are strongly contam-inated by atmospheric absorption / emission lines which does notallow to recover the full line profile. The uncertainty on the linefluxes is obtained by assuming a lower / upper flux limit equal tothe line intensity at the edges of the region a ff ected by the telluriclines. On the contrary, ν = − =
30 to 40) linesdo not su ff er from telluric contamination, and their profiles are Table 6.
Infrared excess with respect to the photospheric modelmeasured at di ff erent wavelengths. Reported errors are due to theinstrumental errors and to di ff erent measurements from di ff erentobservations, where available. Wavelength ( µ m) Excess (mag)3.6 0.83 ± ± ± ± ± -100 -50 0 50 100 Velocity (km/s) N o r m a li z ed F l u x HWZI
Telluriccontamination
High-J lines (J=30-40) Low-J lines (J=1-13)
Fig. 3.
Average profiles of the CO ν = − obs of the lines. The red and blue crossesindicate the position of ν = − ν = − R in = GM ∗ ( ∆ V obs / sin i ) (cid:39) . · (sin i ) AU (1)where M ∗ is the stellar mass, ∆ V obs =
65 km / s is the Half Widthat Zero Intensity (HWZI) of the CO profile, and i is the diskinclination (see Sect. 4.1.1). As shown by e.g. Pringle (1981), the accretion luminosity, L acc ,released in accretion disks or boundary layers, is related to themass accretion rate, ˙ M acc , and depends on the assumed widthover which the emission occurs (see e.g. Bertout, Basri, & Table 7.
Accretion luminosity and mass accretion rate esti-mates from optical / NIR line luminosities and from optical ex-cess. References: (1) Fang et al. (2009); (2) Muzerolle et al.(1998c); (3) Eq. (3) of Hartigan et al. (1995).
Method L acc ˙ M acc Ref.(L (cid:12) ) (10 − M (cid:12) / yr)H β luminosity 0 . + . − . . + . − . (1)He i luminosity 0 . + . − . . + − . (1)H α luminosity 0.17 + . − . . + . − . (1)Pa β luminosity 0.11 + . − . . + . − . (2)Br γ luminosity 0.19 + . − . . + − . (2)Visual excess 0.12 + . − . . + . − . (3) Bouvier 1989). In this work, we consider the accretion luminos-ity released in the impact of the accretion flow, as: L acc (cid:39) (cid:32) − R ∗ R in (cid:33) GM ∗ R ∗ ˙ M acc (2)(Gullbring et al. 1998) where R ∗ and M ∗ are the stellar radiusand mass, and R in is the disk inner radius. We were not able todirectly measure the accretion luminosity since the UV region isonly partially covered by the available observations. Therefore,we estimated the mass accretion rate by employing observed em-pirical correlations between L acc and the luminosity of opticaland NIR emission lines which are thought to be excited in the ac-cretion columns (Fang et al. 2009, Muzerolle et al. 1998c) suchas H α , H β , He i , Br γ , and Pa β (see Table 4). Furthermore, weestimated the accretion luminosity from the visual excess withrespect to the stellar photosphere as in Hartigan et al. (1995).The derived estimates are in agreement within a factor 2,with an average value of L acc = .
15 L (cid:12) , corresponding to˙ M acc = . · − M (cid:12) / yr (see Table 7). The largest uncertain-ties on the estimated L acc and ˙ M acc are due to the scattering ofthe empirical correlations. On the contrary, the errors on L acc ob-tained from the visual excess are due to the continuum determi-nation and hence to the uncertainty on the estimated A V . Furtheruncertainty may be due to the employed bolometric corrections. Optical and IR forbidden lines (e.g. atomic oxygen lines) are typ-ical jet tracers. Following the correlation found by Hollenbach(1985), the [O i ] 63 µ m line is commonly used to constrainthe mass outflow rate (see e.g. Ceccarelli et al. 1997, Podio etal. 2012). A similar correlation has been found for the optical[O i ] 6300 Å (see e.g. Hartigan et al. 1995).The [O i ] 63 µ m from FT Tau was detected only in the centralspaxel (see Sect. 3). Thus, the line originates in a region aroundthe source smaller than ∼ i ] 63 µ m and the continuum flux at 63 µ m forTaurus sources showing no evidence of outflow. Jet sources in-stead show a line flux exceeding the value predicted by the cor-relation by up to two orders of magnitude, indicating that a sig-nificant fraction of the emission is produced in the jet / outflow.The correlation by Howard et al. (2013) indicates that forFT Tau, up to 85% of the observed [O i ] 63 µ m line flux could Table 8.
Mass outflow rate estimate. References: (1) Eq. (A11)of Hartigan et al. (1995); (2) Eq. (A13) of Hartigan et al. (1995)
Method ˙ M W Ref.(10 − M (cid:12) / yr)[O i ] 6300 Å luminosity < . i ] 63 µ m luminosity < . originate in the disk. Similarly, also a fraction of the observed[O i ] 6300 Å flux could be produced in the disk. Thus, we usedthe [O i ] 6300 Å and the [O i ] 63 µ m line luminosity and thecorrelation by Hollenbach (1985) and Hartigan et al. (1995) toderive an upper limit on the mass outflow rate ˙ M W (see Table8). We found that ˙ M W < . · − M (cid:12) / yr. Another uncertaintyof the [O i ] 63 µ m flux can be source variability, since our op-tical spectrum has been flux-calibrated through photometry (seeSect. 5.1).
4. Modeling the disk of FT Tau
To interpret the SED and the available line emission from thedisk, we use the Monte Carlo radiative transfer code MCFOST(Pinte et al. 2006) and the thermo-chemical disk modeling code
ProDiMo (Woitke et al. 2009, Kamp et al. 2010) sequentially.We fix the stellar properties as estimated from the data analy-sis (see Sect. 3.1) and run MCFOST to determine a set of dustproperties that reproduces the observed SED. We perform a χ minimization of the SED and obtain a base set of parameters.As a second step, we run a grid of ProDiMo models to study thebehavior of predicted line fluxes by comparing the results withthe observations. The results of the dust and gas modeling arediscussed in Sect. 4.1 and Sect. 4.2 respectively.
MCFOST calculates thermal and chemical properties of the dustdisk by treating grains as spherical and homogeneous particles.It is based on the Monte Carlo method, allowing monochromaticphoton packets to propagate through the circumstellar environ-ment. Both photospheric emission and dust thermal emission areconsidered as radiation sources.As shown in Sect. 3.3.1 and Table 7, observed emission lineluminosities suggest L acc values between 0.09 and 0.19 L (cid:12) . Tostudy the impact of the UV luminosity, we assumed in the mod-els an average value of L acc = .
15 L (cid:12) . Then, to explore possiblevariability we also assumed a value three times lower. To trans-late from accretion luminosity to f UV = L UV (90 −
250 nm) / L ∗ ,we assumed a blackbody spectrum at 10,000 K for the accretionshock. This yields f UV = .
07 and 0.025 (in the following de-noted as high and low f UV ). In the models, the UV spectrum inthe narrow range between 90 and 250 nm is approximated by apower-law F λ ∝ λ p UV with p UV = . χ minimization leaving only the disk dust mass andoptical extinction as free parameters. The fixed / explored param-eters and the results from the fitting are listed in Table 9 (see Fig.4). The χ fit of the SED results in an optical extinction A V = . Fig. 4.
SED predicted by MCFOST using the input parame-ters listed in Table 9 (black solid line) overplotted on the ob-served photometric points (red dots). The blue spectrum is theSpitzer / IRS spectrum.in agreement with the result inferred by photometric colors. Weobtain a disk dust mass M d = · − M (cid:12) . The outer radius ispoorly constrained and we explore in the following the impactof three di ff erent outer radii, 50, 100 and 200 AU (see Sect. 5.3for a discussion). Below, we discuss some modeling aspects anddegeneracies in more detail. Guilloteau et al. (2011) derived an inclination of 23 ± ◦ fromCO sub-millimeter line kinematics. However, the performedSED modeling indicates that for a disk inclination i < ◦ ,the predicted SED has a NIR bump considerably lower than ob-served. On the other hand, if the inclination is i > ◦ , the photo-spheric emission is shielded too much by the disk. We found thatthe SED is well reproduced for i = ◦ . However, a fine-tuningof this parameter is hardly feasible and the degeneracy betweeninclination and inner radius (arising from the analysis of the COlines, see Sect. 3.2.2) remains unsolved (see Appendix A).Assuming an inclination of 60 ◦ , the disk inner radius inferredfrom the profile of the CO ro-vibrational lines is R in = .
05 AU(see Eq. 1). The estimated disk inner radius is at (cid:39) R ∗ , inagreement with estimates of the magnetic truncation radius fortypical CTTSs (Shu et al. 1994, Donati et al. 2008, Long et al.2011).The outer radius of the disk is poorly constrained by theSED. Models with R out =
50, 100, and 200 AU (all other pa-rameters kept constant) result in the same SED within the pho-tometric error bars (see Sect. 5.3 for a more detailed discussion).
The 10 µ m silicate feature suggests that sub- µ m size dust grainsare still present in the surface layer of the disk (see Sect. 3.2.1).The mineralogy used in our models is amorphous MgFeSiO olivine (Dorschner et al. 1995). As the most simple working hy-pothesis, we assume that the dust is homogeneous in composi-tion throughout the disk. Furthermore, we have constraints onthe gas inner radius R in from the analysis of CO ro-vibrational Table 9.
Parameters of the FT Tau model. The di ff erence between an ‘explored’ and ‘free’ parameter is that the former is set afteran exploratory parameter study while the latter is derived using χ fitting of the SED. Parameter Symbol Comments ValueStellar luminosity L ∗ Derived from observations 0.35 L (cid:12)
Stellar mass M ∗ ” 0.3 M (cid:12) Stellar radius R ∗ ” 1.7 R (cid:12) E ff ective temperature T e ff ” 3400 KDistance d ” 140 pcSlope of UV excess distribution p UV Fixed in MCFOST 2.0Slope of grain size distribution a pow ” 3.5Dust mass density ρ d ” 3.5 g cm − Slope of surface mass density (cid:15) ” − R Explored with MCFOST 100 AUFlaring reference height H ” 12, 14 AU (high / low f UV )Flaring exponent β ” 1.15Disk inner radius R in ” 0.09, 0.05 AU (high / low f UV )Minimum dust grain size a min ” 0.05, 0.1 µ m (high / low f UV )Maximum dust grain size a max ” 1 cmStratification exponent s set ” 0.2, 0.3 (high / low f UV )Stratification grain dimension a set ” 0.05, 0.1 µ m (high / low f UV )Inclination i ” 60 ◦ Disk outer radius R out ” 50, 100, 200 AUOptical extinction A V Free parameter in MCFOST 1.6Disk dust mass M d ” 9 · − M (cid:12) Cosmic Ray Ionization rate ζ Fixed in
ProDiMo . · − s − UV excess f UV Explored with
ProDiMo f PAH ” 10 − , − , − Disk gas mass M g ” (9, 4.5, 1.8) · − M (cid:12) line profiles ( R in = .
05 AU, see Sect. 3.2.2 and above). If weassume that the dust and gas inner radii are coincident, the dusttemperature at that radius has to be below the sublimation tem-perature. Thus, grains with a < . µ m cannot survive at thatradius and we use a min = . µ m. Another possibility is that afraction of the CO ro-vibrational line emission originates fromgas inside the dust sublimation radius. a max is not well con-strained and degenerate with the slope of the grain size distri-bution. Thus, we use a max = f UV model requires a slightly larger R in of 0.09 AU and has a slightly smaller minimum grain size, a min = . µ m. The surface mass density of the disk is parametrized as Σ ∝ r − (cid:15) (3)and the disk scale height as H ( r ) = H · (cid:32) rR (cid:33) β (4)with r being the distance from the star and H the disk height atthe reference radius R . According to the qualitative analysis ofthe IR excess in Sect. 3.2.1, the disk is mildly flared and we thusfixed the flaring angle to be β = .
15. The best fit resulted in ascale height of H =
12 AU and H =
14 AU for the high andlow f UV respectively at R =
100 AU. Smaller scale heights leadto an underprediction of the FIR fluxes.Dust settling is parametrized assuming that the scale heightchanges with grain size for grains larger than a set H ( r , a ) = H ( r ) · ( a / a set ) − s set / (5) The best match of the observed silicate features is found by in-cluding dust settling with all particles involved, i.e. a set = µ m, and an exponent s set = ProDiMo calculates the chemistry and heating / cooling of the gasself-consistently using a large chemical network of 111 speciesand 1462 reactions. An extensive list of all heating and cool-ing processes can be found in Woitke et al. (2009, 2012). In thiswork, we do not feed the gas temperatures back into the verti-cal hydrostatic equilibrium, but instead keep the vertical flaringstructure given by the MCFOST parametrization found for thebest fitting SED model. Using the results from the MCFOSTmodels described in the previous section, we ran a small grid of ProDiMo models with di ff erent values of UV excess, gas massand PAH abundance.Two UV excess cases were considered (high state, f UV = .
07, and low state f UV = . M dust as suggested by MCFOST and explored dust-to-gas mass ratiosof 0.01, 0.02, and 0.05, i.e. M gas = (cid:12) (hereafter denoted as hGAS, iGAS, and lGAS). Even the mostmassive model with M gas = (cid:12) is gravitational stable ac-cording to the Toomre criterion (see Eq. A.10 of Kamp et al.2011). The abundance of PAHs, f PAH , was set to 10 − , − , and10 − times the one in the ISM (hereafter denoted as hPAH, iPAH,and lPAH). The combination of these three parameters yields atotal of 18 disk models.The level populations for the line radiative transfer are cal-culated from statistical equilibrium and escape probability (seeWoitke et al. 2009, for details). Using these populations, wecarry out a detailed line radiative transfer using ray tracing andtaking into account the disk rotation and inclination (Woitke et Fig. 5.
From top to bottom, left to right: the total hydrogen number density, the dust temperature, the gas temperature, and the COabundance distribution in the reference disk model, namely the case with low UV excess, low gas mass, and low PAH abundance.The black contours indicate the total A V = . , i ] 63 µ mline and three representative CO ro-vibrational lines, ν = ν = A V >
1. The CO abundance reaches a maximum value of ∼ − already well above that line and the top CO layer resides at tem-peratures above 1000 K inside 10 AU. Since the CO fundamen-tal ν = − . The main cooling processes are CO rotationaland ro-vibrational line cooling as well as water line cooling. We use in this study the large CO model molecule compiled byThi et al. (2013) including IR and UV pumping. The model uses 7 vibrational levels of the X Σ + and A Π electronic states and60 rotational levels within each of them. ProDiMo calculatesthe level populations from statistical equilibrium and performsa detailed line radiative transfer to obtain the emerging CO linefluxes (Woitke et al. 2011). This type of thermo-chemical model-ing leaves no freedom to adjust CO densities, column densities,densities of collision partners or gas temperatures.The
ProDiMo models show that CO ro-vibrational linefluxes are very weakly a ff ected by the PAH abundance whilethey substantially correlate with the gas mass (Table 10). Modelswith high UV excess generally over-predict the observed fluxes(up to a factor 15). However, those with low gas mass well re-produce or slightly over-predict (by a factor 2) all ν = − ν = − ν = − N = · cm − ). Theonly collision partner is atomic hydrogen; this is taken as a rep- Fig. 6.
Cumulative flux distribution from vertical escape proba-bility in a
ProDiMo model with low UV excess, low gas mass,and low PAH abundance. The top panel of each box shows thevertical optical depth for the line and continuum. The middlepanel of each box shows the cumulative flux from the
ProDiMo model (black) and best fit slab model (red). The bottom plot ofeach box shows the CO density distribution in the disk. Outlinedwith black contours is the region in which radially and verticallybetween 15 and 85% of the flux originates. The top box is forthe ν = − ν = − ν = − χ minimization, we find a hydrogen volumedensity profile n H ( r ) = . · ( r / R in , slab ) − cm − and a gastemperature profile T ( r ) = r / R in , slab ) − . K. The inner ra-dius R in , slab is only loosely constrained to 0 . + . − . AU, becausethe line wings have a rather low S / N. The outer radius is largelyunconstrained due to the degeneracy between the surface densityof the gas N and the outer radius of the emitting area; R out , slab hasto be larger than 0.9 AU. The turbulent line broadening b is foundto be 2 km / s, although N and b are degenerate. The gas tempera-ture found at the inner radius is T = + − K. Integrated linefluxes have been measured from the spectrum generated withthe slab model in the same way as for the observed spectra. Themodel predicts correctly the ν = − ν = − T ex ∼ ν = − ff use (lower volume density), thatthe line flux declines steeper with distance from the star (steeperdensity power-law), or that some additional non-LTE e ff ects arestill missing in the slab model.Fig. 6 shows the cumulative flux distribution from simplevertical escape probability in the lUV / lPAH / lGAS ProDiMo model for the three representative lines, ν = ν = ν = / lPAH / lGAS model, the three lines studied here as repre-sentative lines are a factor 2-3 lower than their respective LTEvalues. The agreement between ProDiMo models and the moresimple slab models on how the flux is building up as a functionof radius is very good (see comparison in Fig. 6). The
ProDiMo models also show a similar temperature of ∼ ff ect the CO ro-vibrational lines much stronger in the former case. Another pos-sibility could be that part of the CO ro-vibrational emission orig-inates from gas inside the dust sublimation radius of our models.At this stage, a further analysis of these CO ro-vibrational linesis largely limited by the observations, which have limited spec-tral resolution and su ff er from a low signal-to-noise and telluriccontamination (see Sect. 2.4). i ] 63 µ m line As shown in Table 10, the [O i ] 63 µ m line flux is a ff ected by theUV excess, the gas mass, and the PAH abundance. Di ff erencesof a factor two are found between the high and low UV excessmodels. The dependence on the gas mass stems from the factthat the gas temperature changes with disk mass and, in turn,a ff ects the line flux. The fact that the flux increases with PAHabundance is explained by the increasing photoelectric heatingof the gas in the upper disk layer (Jonkheid et al. 2004). Thus,gas mass, UV excess and PAH abundance are to some degreedegenerate in the prediction of the [O i ] line flux. J -71-70-69-68-67-66 l n ( F / v A ( J + )) High UV/High Gas ModelHigh UV/Intermediate Gas ModelHigh UV/Low Gas ModelLow UV/High Gas ModelLow UV/Intermediate Gas ModelLow UV/ Low Gas ModelSlab ModelObserved 1-0Observed 2-1 J -74-73-72-71-70-69-68 Fig. 7.
Rotational diagram of CO ro-vibrational lines observed (left: ν = −
0, right: ν = −
1) versus model predicted: slab modeldescribed in the text and
ProDiMo runs with intermediate PAH abundance, with high, intermediate, and low gas mass, and withhigh and low UV excess. The vertical dotted lines indicate discontinuity in the x-axis.
Table 10.
Fluxes of the [O i ] 63 µ m line and of three representa-tive CO ro-vibrational lines as predicted by the grid of ProDiMo models using detailed line radiative transfer, the slab model de-scribed in the text, and as observed.
Model Flux (10 − W / m )[O i ] CO CO COUV / PAH / GAS 63 µ m 1-0 P4 2-1 P4 1-0 P36h / h / h 30.4 6.51 5.07 10.4h / i / h 21.7 5.10 4.79 10.3h / l / h 20.9 4.95 4.77 10.3h / h / i 24.7 4.35 2.99 7.05h / i / i 18.3 3.55 2.86 6.95h / l / i 17.7 3.41 2.84 6.59h / h / l 17.5 2.11 1.31 3.09h / i / l 13.4 1.81 1.26 3.02h / l / l 12.9 1.77 1.25 3.02l / h / h 15.6 2.97 2.03 5.31l / i / h 11.3 2.17 1.80 5.05l / l / h 10.8 2.08 1.78 5.03l / h / i 13.1 1.94 1.15 3.04l / i / i 9.92 1.44 1.06 2.91l / l / i 9.59 1.38 1.05 2.90l / h / l 10.1 1.03 0.55 1.26l / i / l 7.99 0.79 0.52 1.21l / l / l 7.77 0.77 0.52 1.20Slab model - 1.25 0.40 1.31Observed 1.6 ± . ± .
38 0.28 ± .
10 0.83 ± .
5. Discussion
In this section, we discuss the source variability and the resultsof our detailed disk modeling in the context of the available ob-servational data.
Young circumstellar systems are often highly variable objects(see e.g. Bouvier et al. 1993). Flux variability up to a factor ∼ ∼ ff ective temperature of M3 stars. The B bandshows an amplitude slightly smaller than expected for those hotspots but this can be explained by the presence of a hot contin-uum in addition to the pure photospheric emission.Given this, the brightness at the minimum of the light curvesprovides an upper limit to the photospheric brightness of the star.As we see from Fig. 8, the reddened photospheric emission inthe V band assumed in Sect. 3.1 is lower than measurements,from either ASAS or Calar Alto surveys. This indicates that theextinction cannot be much lower than estimated in Sect. 3.1, be-cause this would increase the reddened photospheric emission ofthe model to values higher than the observed one. The USNO Vband photometry used to flux-calibrate the TNG spectrum (and,thus, to estimate the mass accretion rate, see Sect. 2.1 and 3.3.1)turns out to be an average value of all measurements (see Fig. 8).In order to address the origin of the observed variability, atime-dependent study of optical / NIR emission lines is necessary.Any relation between these lines and contemporary observationsof optical photometric variations can clarify if and how muchof the observed variability is due to the accretion process. Inaddition, we must be careful in the interpretation of line emissionfrom the disk surface especially if these lines result from UVpumping by stellar radiation.
The CO ro-vibrational lines are very sensitive to the extent ofthe hot gas surface layer. The observations clearly indicate thatthe lines are typically very wide (HWZI (cid:39)
65 km / s). In the mod-els the CO ro-vibrational lines predominantly arise from this hotsurface layer (Sect. 4.2). The UV radiation field a ff ects the extent of this hot surface and it can change due to the particular choiceof the dust opacities, the scale height of the disk and the flaring.The quality of the available Keck CO ro-vibrational line profilesis not good enough to derive the extent of the hot surface layerdirectly from their shape. In case of exquisite data quality, thiscan be done as shown by Goto et al. (2012) for the example ofHD100546, a Herbig Ae star. So, as new data will become avail-able, these parameters should be refined keeping the constraintson the SED. i ] 63 µ m line All models presented here over-predict the [O i ] 63 µ m line.Since the line is optically thick, its flux is mostly a ff ected bythe gas temperature in the emitting region and the total emittingsurface area. The models indicate that the [O i ] 63 µ m line typ-ically originates between ∼
10 and 200 AU. Roughly 15% ofthe total line flux builds up between 100 and 200 AU. In orderto understand the dependence of the predicted [O i ] line flux onthe adopted disk size, we calculated models with di ff erent outerradii (50, 100, and 200 AU, see Table 9). We find that the lineflux decreases by only a factor 4 for the smallest disk size. Thisis due to the fact that the smaller emitting area is partially com-pensated by the higher gas temperature of the emitting region. Atthe same time, the CO ro-vibrational lines do not change withinthe modeling uncertainties. Hence, the observed emission linesdo not allow us to put any stronger constraint on the size of thegaseous disk.Guilloteau et al. (2013) derive from IRAM 30-mCN N = − R out =
57 AU(Guilloteau et al. 2011). However, the disk is barely resolved andbetter interferometric images at shorter wavelength are requiredto measure R out for the dust; at the same time, interferometricline data e.g. for CO isotopologues are required to obtain areliable outer gas radius. Previous work shows that gas and dustouter radii at submm wavelength can actually di ff er (e.g. Isellaet al. 2007, Andrews et al. 2012). Given the existing uncertaintyon the estimate of R out , models with di ff erent gas and dust outerradii have to await better observational data.In the region where the [O i ] 63 µ m line emits, photoelectricheating is one of the dominant heating processes. Since the PAHfeatures are not observed in the Spitzer spectra, their abundancecan be arbitrarily low. Supressing the PAH abundance even be-low the lowest value in the grid, f PAH = − , does not a ff ectthe [O i ] line flux anymore. A lower disk gas mass shifts the lineforming region to lower depth in the disk, thus making the lineflux weaker.
6. Summary
We have performed analysis and modeling of the SED andemission lines of the TTS FT Tau to fully characterize thestellar, disk, and accretion properties. We reduced and anal-ysed five spectra from optical to FIR wavelengths, taken withthe optical Telescopio Nazionale Galileo, the NIR WilliamHerschel Telescope, the NIR Nordic Optical Telescope, the high-resolution NIR Keck Observatory, and the FIR
Herschel
SpaceTelescope. Additional data were retrieved from the literature andconsist of a MIR Spitzer Space Telescope spectrum and of 44photometric measurements from optical to radio wavelengths.We studied a set of models of the source generated by means of the radiative transfer code MCFOST and the thermo-chemicaldisk modeling code
ProDiMo .We found that FT Tau is a low-mass (0 . ± . (cid:12) ) and -luminosity (0 . ± .
09 L (cid:12) ) M3 star showing very high vari-ability (probably due to photospheric hot spots). The estimatedproperties are typical of young unevolved systems in the Taurusstar forming region with ages of roughly 1 Myr. The optical ex-tinction ( A V = .
8) is also within the range of typical values inTaurus.The inner radius of the circumstellar disk is small(0 . + . − . AU) indicating an early stage of internal disk dis-sipation. This is in agreement with the fact that the sourceis strongly accreting. In fact, the derived mass accretion rate((3 . ± . · − M (cid:12) / yr) is a high value for M-stars, since thiscorresponds to L acc ∼ . L ∗ . The ratio of mass outflow to massaccretion rate is lower than 0 .
03, in agreement with typical ob-served values for TTSs (Hartigan et al. 1995). The disk is quitemassive ( ∼ .
02 M (cid:12) ) with respect to the stellar mass (mass ratio ∼ . ffi cient at removingthis small grain population. These findings are consistent withthe apparent primordial nature of this disk (e.g. no gaps, holes).The PAH abundance is inferred to be extremely low ( ∼ − times that in the ISM).From this work it is clear that FT Tau can be considered asbenchmark for primordial disks in the Taurus molecular cloudwith high mass accretion rate, high gas content, and typical disksize. It is an interesting target for follow-up chemical studies aswell as for informing surveys of star forming regions on moreprototypical objects to expect. Acknowledgements.
We acknowledge the referee for valuable comments thatconsiderably improved the paper. We thank Matilde Fern´andez and Victor Terr´onfor observing FT Tau and reducing the data at the Calar Alto Observatory. Wereally appreciate the helpful discussion about the origin of the variability. Wealso gratefully thank Gwendolyn Meeus for her work in acquiring data with theTNG and Ilaria Pascucci and Veronica Roccatagliata for reducing the data fromSpitzer. This work is supported by the Swiss National Science Foundation. LPacknowledges the funding from the FP7 Intra-European Marie Curie Fellowship(PIEF-GA-2009-253896). IK, WFT, FM, and PW acknowledge funding from anNWO MEERVOUD grant and from the EU FP7- 2011 under Grant Agreementnr. 284405. FM acknowledges support from the Millennium Science Initiative(Chilean Ministry of Economy), through grant Nucleus P10-022-F. I. Pascucciacknowledges NASA / ADP Grant NNX10AD62G. This research has made useof the SIMBAD database, operated at CDS, Strasbourg, France.
Appendix A: Uncertainties of the analysis
In this appendix we discuss the limitations of our analysis dueto non-simultaneous observations and quantify thoroughly theuncertainties on the inferred results.The mentioned stellar variability does not have strong impacton the determination of the stellar properties because it does nota ff ect significantly the shape of the optical spectrum (used todetermine the spectral type) and the J and H band fluxes (used toestimate luminosity and radius).On the contrary, the estimate of the visual extinction is af-fected by large uncertainties. We firstly remark that the use ofthe (J-H) color as tracer of the extinction relies on the assump-tion that the observed flux at those wavelengths is entirely emit-ted by the stellar photosphere. Secondly, the determination ofthe optical extinction A V is pretty sensitive to the surface grav-ity of the assumed model. By varying the stellar radius or massby 30%, we obtain A V values between 1.2 and 2.5. This mayadd a further factor 15% uncertainty to the estimates of stellar Time (HJD-2450000) M agn i t ude Calar Alto surveyASAS surveyUSNO photometry (used in this work)Photospheric emission (assumed in this work)
Fig. 8.
ASAS, Calar Alto, and USNO photometric measurements of FT Tau in the V band. The dashed line indicates the reddenedphotospheric magnitude assumed in our analysis. The vertical line indicates a gap in the x-scale. The position in time of the USNOphotometry is arbitrary.properties. However, the fact that the accretion luminosity val-ues estimated by using di ff erent tracers from 0.45 and 2.17 µ mdoes not show a dependence with the wavelength (see Fig. A.1)is a strong sanity check for the determination of A V . The factthat we find the same A V values by using two independent meth-ods (the observed colors, Sect. 3.1, and the modeling approach,Sect. 4.1) further reinforces our result. The large di ff erence be-tween our estimate of the stellar luminosity and the result fromRebull et al. (2010) (see Table 1) is due to the determination of A V which is in turn due to the assumed surface gravity.The spectral type- T e ff relation can actually introduce an ad-ditional error. Di ff erences up to some hundreds of Kelvin arisefor M-type stars among di ff erent works (see e.g. Da Rio et al.2010). Finally, further uncertainty in the determination of thestellar properties is provided by the PMS star tracks adopted toinfer the stellar mass and age. Hartmann (2001) suggested thatthe age spread inferred for TTSs in Taurus may exclusively bedue to uncertainties towards individual members.In Sect. 3.2.2 we estimated the disk inner radius by measur-ing the width of the CO ro-vibrational lines. The largest uncer-tainty in the determination of R in is set by the adopted inclina-tion. The width of the CO lines is equally reproduced by config-urations with ( i : R in ) = (60 ◦ : 0.05 AU), (45 ◦ : 0.03 AU), and(30 ◦ : 0.02 AU).The estimates of the mass accretion and outflow rate may bea ff ected by variability, since the optical and NIR spectra usedto measure the line luminosities were flux-calibrated by usingnon-simultaneous photometry. This is particularly true for esti-mates based on optical lines (optical flux variability ∼ L acc (Sect. 3.3.1). Thelowest and the highest estimates for L acc have been found bymeans of emission lines from the same spectrum (thus taken si-multaneously, see Table 7). This is indicating that the scattering H β . H e I . H α . P a β . B r γ . -2-101 l og ( A cc r e t i on l u m i no s i t y ) ( L ⊙ ) Visual extinction = 0.0Visual extinction = 1.8Visual extinction = 4.0
Fig. A.1.
Accretion luminosity estimated from the luminosity ofemission lines at optical to NIR wavelengths (x-axis) for di ff er-ent A V values. The grey stripe indicates the range of values in-ferred from the Br γ line, which is the least a ff ected by extinction.It is clear the trend with wavelength for high and no extinction.Slight displacement between points has been put for a better vi-sualization.of the empirical relations might play the major source of uncer-tainty on the accretion luminosity. References
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