Draftversion January30,2013 PreprinttypesetusingLATEXstyleemulateapjv.5/2/11 MAGNETICALLY CONTROLLED ACCRETION ON THE CLASSICAL T TAURI STARS GQ LUPI AND TW HYDRAE Christopher M. Johns–Krull & Wei Chen DepartmentofPhysics&Astronomy,RiceUniversity,6100MainSt. MS-108,Houston, TX77005, USA Jeff A. Valenti SpaceTelescopeScienceInstitute, 3700SanMartinDr.,Baltimore,MD21210,USA 3 1 0 Sandra V. Jeffers 2 Georg-August-Universit¨at,Institutfu¨rAstrophysik,Friedrich-Hund-Platz1,D-37077G¨ottingen, Germany n a Nikolai E. Piskunov, Oleg Kochukhov, V. Makaganiuk, & H. C. Stempels J DepartmentofAstronomyandSpacePhysics,UppsalaUniversity,75120Uppsala,Sweden 9 2 Frans Snik, Christoph Keller, & M. Rodenhuis ] SterrewachtLeiden,LeidenUniversity,NielsBohrweg2,2333CA,Leiden,TheNetherlands R Draft version January 30, 2013 S ABSTRACT . h We present high spectral resolution (R ≈ 108,000) Stokes V polarimetry of the Classical T Tauri p stars (CTTSs) GQ Lup and TW Hya obtained with the polarimetric upgrade to the HARPS spec- - o trometeronthe ESO3.6mtelescope. We presentdataonbothphotosphericlines andemissionlines, r concentrating our discussion on the polarization properties of the He I emission lines at 5876 ˚A and t s 6678˚A.TheHeIlinesintheseCTTSscontainbothnarrowemissioncores,believedtocomefromnear a theaccretionshockregiononthesestars,andbroademissioncomponentswhichmaycomefromeither [ a wind or the large scale magnetospheric accretion flow. We detect strong polarizationin the narrow 2 componentofthetwoHeIemissionlinesinbothstars. We observeamaximumimpliedfieldstrength v of 6.05±0.24 kG in the 5876 ˚A line of GQ Lup, making it the star with the highest field strength 2 measured in this line for a CTTS. We find field strengths in the two He I lines that are consistent 8 with each other, in contrast to what has been reported in the literature on at least one star. We do 1 notdetectanypolarizationinthe broadcomponentofthe HeIlines onthese stars,strengthening the 3 conclusion that they form over a substantially different volume relative the formation region of the . narrow component of the HeI lines. 1 0 Subject headings: accretion,accretiondisks —line: profiles—stars: atmospheres—stars: formation 3 — stars: magnetic fields — stars: pre–main-sequence — 1 : v 1. INTRODUCTION Basri 2003, Jayawardhana et al. 2003). Independent of i X TTauristars(TTSs)areyoung(∼< 10Myr),lowmass the exact constraint imposed for defining a CTTS, this moniker has become synonymous with a low mass pre- ar (∼< 2.5 M⊙) stars that have only recently emerged from main star that is actively accreting material from a cir- theirnatalmolecularcloudcorestobecomeopticallyvis- cumstellardisk. Indeed,the vastmajorityofstarswhich ible. These young, low mass stars are generally subdi- fit the criteria for CTTSs show some kind of additional vided into categories such as classical and weak TTSs. evidence [e.g. inverse P-Cygni line profile shapes, opti- ThedesignationofaclassicalTTS(CTTS)wasoriginally calveiling (see below), infraredexcess]indicative of disk based on a purely observational distinction: the equiva- accretion. lent width of the Hα emission line. Classical TTSs are It is now generally accepted that accretion of circum- TTSswhichhaveanHαequivalentwidthW (Hα)>10 eq stellar disk material onto the surface of a CTTS is con- ˚AasdistinguishedfromtheweaklineTTSs(WTTSs)de- trolledby a strong stellar magnetic field (e.g. see review finedbyHerbigandBell(1988);however,Bertout(1989) byBouvieretal. 2007). Thesemagnetosphericaccretion suggeststhatabreakpointvalueof5˚Aismoreappropri- modelsassertthatstrongstellarmagneticfieldstruncate ate. More recently, investigatorshave tied the definition the inner disk, typically near the corotation radius, and totheshape(width)oftheHαlineprofile(e.g. White & channel the accreting disk material onto the stellar sur- face,mostoftenathighstellarlatitude(Camenzind1990; [email protected],[email protected] K¨onigl 1991; Collier Cameron & Campbell 1993; Shu et [email protected] jeff[email protected] al. 1994; Paatz & Camenzind 1996; Long, Romanova, [email protected],[email protected], vitaly.makagani&[email protected],0e0ri5c.)[email protected] [email protected],[email protected],[email protected] 2 simulations find that outflows launched from near the present even when the net flux weighted line of sight region in which the stellar field interacts with the sur- field integrated over the stellar surface (the net longitu- rounding accretion disk can also spin the star down to dinal magnetic field, B ) is zero. Observations of time z observedrotationrates(e.g. Ferreira2008;Romanovaet series of Stokes V spectra can be used to track changes al. 2009),thoughsomerecentworkchallengesthenotion in the amount of net field visible on the star as it ro- that these outflows can actually balance the spin-up ac- tates, ultimately allowingthe large scale field of the star cretion torques in CTTS systems (e.g. Zanni & Ferreira to be mapped using various tomographic imaging tech- 2009). niques (e.g. Donati et al. 2007 and references therein; Despite the successes of the magnetospheric accretion Kochukhov et al. 2004 and references therein). model, open issues remain. Most current theoretical In addition to potentially mapping the surface field models assume the stellar fieldis a magnetic dipole with on accreting young stars, information can be obtained the magnetic axis aligned with the rotation axis. How- on the large scale field controlling the interaction of the ever,recentspectropolarimetricmeasurementsshowthat star with its disk and the accretion flow by measuring thefieldsonTTSsareprobablynotdipolar(Johns–Krull time series of Stokes V profiles in emission lines formed et al. 1999a; Valenti & Johns–Krull 2004; Daou, Johns– in the accretion flow and shock. The first accretion line Krull, & Valenti 2006; Yang, Johns–Krull, & Valenti for which circular polarization was detected is the He I 2007; Donati et al. 2007, 2008, 2010a, Hussain et al. line at 5876 ˚A (Johns–Krull et al. 1999a), and time se- 2009). Few studies of accretion onto CTTSs have taken ries of the polarization variations in this line have been into account non-dipole field geometries. The earliest used to estimate the latitude of accretion spots on sev- of these by Johns–Krull and Gafford (2002) found that eral CTTSs (e.g. Valenti & Johns–Krull 2004; Yang et abandoning the dipole assumption reconciled observed al. 2007;Donati et al. 2008,2010b,2011a,2011b). This trends in the data with model predictions; however,this line is observedin most CTTSs and is often found to be studydidnotconsiderthetorquebalanceonthestarand composed of two components: a narrow core component whether an equilibrium rotation rate could actually be and a broad component extending out to several hun- achieved. Johns–Krull and Gafford (2002) argued that dred km s−1 (e.g. Edwards et al. 1994; Batalha et al. while the field on the stellar surface may be quite com- 1996; Alencar & Basri 2000). Based on the similarity in plex, the dipole component of the field should dominate shape between the observedline profiles of some CTTSs at distance from the star where the interaction with the and model profiles calculated in the context of magne- disk is taking place. This assumption appears to gener- tospheric accretion, Hartmann et al. (1994) suggested ally hold true in severalrecent studies (e.g. Johns–Krull that the He I 5876 ˚A line (broad and narrow compo- & Gafford 2002; Mohanty & Shu 2008; Gregory et al. nents) might form throughout the accretion flow, with 2008;Longetal. 2008;Romanovaetal. 2011;Cauleyet the narrow component primarily coming from the lower al. 2012). However,the complex nature of the field near velocity regions near the disk truncation point. Beris- the surfacehas significantimplications forthe size ofac- tain et al. (2001) instead argue that the narrow core cretion hot spots, making them smaller than would be of the He I line arises in decelerating post-shock gas on predicted by pure dipole models (Mohanty & Shu 2008; thestellarsurfaceatthebaseoftheaccretionfootpoints. Gregory et al. 2008; Long et al. 2008); and also has im- Beristain et al. (2001) argue that the broad component portant consequences for disk truncation radii and the observed in many CTTSs has a dual origin in the mag- computation of the torque balance on the star by the netospheric flow and in a high velocity wind in the most disk (Gregory et al. 2008; Long et al. 2008; Romanova strongly accreting stars. et al. 2011). Thestrong,orderedfieldsobservedinthenarrowcom- Two approaches are generally used to measure mag- ponent of this line component (∼ 2.5 kG, e.g. Johns– netic fields on low mass stars, both utilizing the Zeeman Krull et al. 1999a) argue for a formation region close to effect. Magnetic fields can be measured from the broad- the stellar surface instead of several stellar radii above ening of magnetically sensitive lines observed in inten- thestarwherethe fieldinteractswiththedisk. TheHeI sity spectra (e.g. Johns–Krull 2007; Yang et al. 2008). 5876 ˚A arises from a triplet state and is composed of Thistechniqueisprimarilysensitivetothemagneticfield several closely spaced lines. The He I 6678 ˚A line arises modulus, the unsignedvalue of the field weightedby the fromtheanalogoussingletstate,andisobservedinmany intensity distribution of the light emitted over the vis- CTTSs as well where it displays both broad and narrow ible surface of the star. While this method does not components (see Beristain et al. 2001). Based on the suffer from flux cancellation due to regions of opposite strong similarity in their kinematic properties and the polarity appearing on the star, it does require that all measuredtriplet-singletfluxratio,Beristainetal. (2001) non-magnetic broadening mechanisms be accurately ac- conclude that the narrow component of both He I lines counted for in the observed spectra. As a result, this formsinthepost-shockgas. Ontheotherhand,thispic- techniqueisprimarilysensitivetorelativelystrongfields. ture is complicated by the observation of Donati et al. ObservationsofcircularpolarizationinStokesV spectra (2008)thatthe 6678˚Aline consistentlyshowsalongitu- can be much more sensitive to weak fields on the sur- dinalmagneticfieldstrengthapproximatelytwicethatof faceofstars;however,theStokesV signatureissensitive the5876˚AlineintheCTTSBPTauwhoseHeIlinesare only to the line of sight component of the magnetic field dominatedbyanarrowcomponent(Edwardsetal. 1994, and the signal can be reduced significantly due to flux Batalhaetal. 1996,Beristainetal. 2001). Thisisasur- cancellation when opposite field polarities are observed prising observation since models of accretion shocks on simultaneouslyonthe stellar surface. Doppler shifts due CTTSsfindthatthethicknessofthepost-shockregionis to stellar rotation can reduce the degree of flux cancel- typically105−106 cm(Calvet&Gullbring1998;Lamzin lationthat results, permitting Stokes V signatures to be 3 1998) which is a small fraction (∼< 10−5) of a stellar ra- et al. 2011), mounted in front of the fibers feeding dius. It would be surprising if the stellar magnetic field the HARPS spectrometer (Mayor et al. 2003). While strength varied so strongly with depth, suggesting then HARPSpol can also record Stokes Q and U spectra, for that perhaps the two He I lines do not trace the same the observations reported here, only Stokes V spectra regions on the stellar surface. were obtained. As mentioned above, linear polarization Tobetterclarifythemagneticfieldpropertiesofaccre- inboththelinesandthecontinuumcanresultfromscat- tion relatedlines, more spectropolarimetricobservations tering off a circumstellar disk (e.g. Vink et al. 2005); ofCTTSs,includingthosewithsubstantialbroadcompo- however, the action of a disk does not typically produce nentstotheirHeIlines,areneeded. Here,wereportnew circularpolarizationineitherthelinesorthecontinuum. observations of two CTTSs (GQ Lup and TW Hya) us- Here, we will focus only on Stokes V in the lines mea- ing the newly commissioned polarimeter operating with suredrelativethe continuumwhichis assumedtonot be the HARPS spectrograph on the ESO 3.6 m telescope circularly polarized. With this instrumental setup, each at La Silla. TW Hya is a K7 CTTS and a member of exposure simultaneously records the right and left cir- the loose TW Hydrae association (Kastner et al. 1997). cularly polarized components of the R = 108,000 spec- TheHipparcosparallaxforTWHyaimpliesadistanceof trum. Thesetwocomponentsoftheechellespectrumare 56±7 pc (Wichmann et al. 1998), making it the closest interleaved,suchthattwocopiesofeachechelleorderare CTTS to the Earth. Based on its placement in the HR presenton the two 2148×4096CCD arrays(one for the diagram, the age of TW Hya is estimated to be 10 Myr blue portion of the spectrum and one for the red). The (Webb et al. 1999,Donati et al. 2011b). Setiawanet al. twopolarizedcomponentsofeachorderareseparatedby (2008) claimed the detection of a ∼ 10 M planet in ∼16pixelsinthecrossdispersiondirectiononthearray, Jup a very close orbit around this CTTS, making TW Hya while each spectral trace is ∼ 3.5 pixels wide (FWHM) an important benchmark constraining the timescale of in the cross dispersion direction. Each observation of a planet formation. Hu´elamoet al. (2008)instead suggest starreportedhereactuallyconsistsof4separateobserva- that the observed radial velocity variations which signal tionsofthe star,withthe angleofthe quarterwaveplate the presence of the planet are in fact caused by large in the polarimeter advanced by 90◦ between the expo- starspots on the surface of TW Hya. As a result, there sures. The result of this is to interchange the sense of is great interest in knowing as much about this star as circular polarization in the two beams. This gives sub- possible. In addition, TW Hya is still accreting material stantial redundancy in the analysis which allows us to from its circumstellar disk and is observed at a low in- remove most potential sources of spurious polarization clination (i ∼ 18◦, Alencar & Batalha 2002), making it due to uncalibrated transmission and gain differences in an excellent object for studying magnetically controlled the two beams. As described below, we use the “ratio” accretion onto young stars. The magnetic properties of method to combine the spectra from these interchanged TW Hya have been investigated a number of times pre- beams in order to form Stokes I and V spectra that are viously (Yang et al. 2005, 2007; Donati et al. 2011b). largely free of these potential spurious signals (e.g. Do- GQ Lup is also a K7 CTTS, and has also recently come nati et al. 1997; Bagnulo et al. 2009). All spectra were under a great deal of scrutiny as the result of a claimed obtainedonthenights29April2010through2May2010, planetary mass companion. Neuha¨user et al. (2005)dis- withonenight(1May)lostduetoweather. Table1gives coveredan infraredcompanion at a separationof ∼0.′′7 acompletetableofthestellarobservationsreportedhere. (corresponding to ∼ 100 AU at a distance of 150 pc). Included in the Table are continuum signal-to-noise es- Based on their infrared photometry and K band spec- timates near the two He I emission lines studied here as tra,Neuha¨useretal. (2005)constrainedthe massofGQ wellasthe emissionequivalentwidths ofthese two lines. Lup B to be between 1−42 MJup, placing it possibly Also reported is the veiling found near the He I 6678 ˚A in the planet regime. More recent spectroscopic studies line as discussed below. Along with spectra of GQ Lup have favored the upper end of this range,suggesting the and TW Hya, a spectrum of the weakly accreting TTS companion is more likely a brown dwarf (Mugrauer & V2129Ophwasalsoobtainedandisusedinthe analysis Neuha¨user 2005; Guenther et al. 2005; McElwain et al. ofthe HeI lines onthe otherstars. Inadditionto stellar 2007;Seifahrtetal. 2007;Maroisetal. 2007;Neuha¨user spectra,standardcalibrationobservationswereobtained et al. 2008). The formation of such an object presents including bias frames,spectra ofa ThoriumArgonlamp challenges to theories of companion formation in a disk, for wavelength calibration, and spectra of an incandes- and has sparked continued study of this system to bet- centlampforthepurposeofflatfielding. Thecalibration ter pin down the properties of both of its members. GQ spectrawereobtainedwiththepolarimeterinfrontofthe Lup is known to show clear signs of variable accretion fibers. (Batalha et al. 2001), making it a good target to study All spectra were reduced with the REDUCE package the role of magnetic fields in the accretion process. To of IDL echelle reduction routines (Piskunov & Valenti our knowledge, no studies of the magnetic properties of 2002)which builds onthe data reductionproceduresde- GQ Lup exist to date. In §2 we describe our observa- scribed by Valenti (1994) and Hinkle et al. (2000). The tions and data reduction. The magnetic field analysis reduction procedure is quite standard and includes bias and results are described in §3, and in §4 we discuss the subtraction, flat fielding by a normalized flat spectrum, implications of our findings. scatteredlightsubtraction,andoptimalextractionofthe spectrum. Theblazefunctionoftheechellespectrometer 2. OBSERVATIONSANDDATAREDUCTION is removed to first order by dividing the observed stel- AllspectrareportedherewereobtainedattheESO3.6 lar spectra by an extracted spectrum of the flat lamp. m telescope on La Silla using the newly comissionedpo- Final continuum normalizationwas accomplished by fit- larimeter, HARPSpol (Snik et al. 2008, 2011; Piskunov 4 ting a 2nd order polynomialto the blaze correctedspec- nent properties (e.g. Alencar & Basri2000). The result- tra in the regions around the lines of interest for this ing equivalent with of the various features then depedns study. Special care was taken to apply a consistent con- atsomelevelonhowonechoosestodotheanalysis. This tinuumnormalizationproceduretothespectraextracted is illustrated in Figure 1. The top panel shows the He I from all four sub-exposures. Occasional small difference 5876˚AlineofTWHyafromthefirstnight. Thesmooth in normalization of the two orthogonalspectra are com- solid curve shows a line profile fit employing 3 Gaussian pensated by using the “ratio” method (e.g. Bagnulo et components. Thedash-dotlineshowsafitusingonlytwo al. 2009,and below) to combine the rightand left circu- Gaussian components. There is a clear difference in the larlypolarizedcomponents. Thewavelengthsolutionfor two fits (the 3 Gaussian fit uses two Gaussians to fit the each polarization component was determined by fitting NC which is not really Gaussian as mentioned above). a two-dimensional polynomial to nλ as function of pixel The bottom panel of Figure 1 zooms in on the line and order number, n, for approximately 1000 extracted to show the recovered BC profiles. The BC from the 3 thorium lines observedfrom the internal lamp assembly. Gaussian fit is shown in the smooth solid line and that The resolution as determined by the median FWHM of from the two Gaussian fit is shown in the dash-dot line. these thorium lines was R=107,660. Also shown is a BC fit (dash-triple dot line) following As mentioned above, each subexposure obtained of a Batalha et al. (1996) where a single Gaussian is used to given star contains both the right and left circularly po- fit the region outside the NC. Finally, the solid straight larizedcomponentofthespectrum. Inordertogetafinal line connecting the two large squares shows a by eye es- measurement of the mean longitudinal magnetic field, timate of the point on both the blue and red side of the Bz, these individual measurements of the two circular line(asseeninStokesI)wheretheNCandBCjoinwith polarizationcomponentsmustbecombinedinsomeway. a linear interpolation between these points to define the We used the “ratio” method (e.g. Bagnulo et al. 2009, separation of the NC and the BC which can be used to Donati et al. 1997)to combined the rightand left circu- separately determine their equivalent widths. This then larlypolarizedcomponentsofthespectraformtheStokes gives 4 different ways to estimate the equivalent width V spectrum as well as a null spectrum, with each being of the BC (and also the NC). The two extremes for the renormalizedto the continuum intensity. We also added BC equivalent width are the single Gaussian fit (1.193 all the components together to form the Stokes I spec- ˚A) following Batalha et al. and that (1.089 ˚A) from the trum. With Stokes V and I determined, the continuum two Gaussian fit, corresponding to a difference of 9%. normalized right-hand circularly polarized (RCP) com- Clearly, none of the Gaussian fits exactly follow the red ponent of the spectrum is then R =I +V and the con- sideoftheBC,soitisimpossibletopredictjustwhatthis tinuum normalized left-hand circularly polarized (LCP) component does under the NC. Given this uncertainty component of the spectrum is L = I −V. Computing and the fact that using the linear interpolation between these fromI andV inthis wayensuresbothcircularpo- theblueandredsidesoftheapparentboundarybetween larization states have been normalized to the same con- theNCandBCgivesequivalentwidthvaluesbetweenthe tinuum. two extremes, we choose to use this method to separate bothcomponentsandmeasuuretheequivalentwidthsre- 3. ANALYSIS portedinTable1. WenotethattheBCequivalentwidth 3.1. HeI Line Equivalent Widths for the profile shown in Figure 1 computed this way dif- fers from that resulting from the 3 Gaussian fit by only Table1givestheequivalentwidthofthetwoHeIlines 4.9%. We therefore estimate that the systematic unc- studiedhereforallourtargetstars. Asmentionedbefore, etainty resulting from the choice of just how to separate previous investigators have noted that these lines often the two components likely leads to a 5% uncertainty in appear to have two distinct components (e.g. Batalha the reported equivalent widths which is not included in et al. 1996; Beristain et al. 2001): a narrow component the Table. In most cases, the boundary between the BC (NC) and broadcomponent (BC). It is thoughtthat the and the NC is clear and repeated measurements with two components may form in different physical regions slightly different choices yield results with a difference of the accretionflow onto CTTSs (Beristain et al. 2001) less than 1σ for the quoted uncertainties. There are a and their polarization properties also appear to be dif- few cases where the boundary between the NC and the ferent with the narrow compnent showing significantly BC, or the BC and the continuum, are less clear and we stronger polarization (Daou et al. 2006; Donati et al. repeated the measurements with a larger distinction in 2011b). We therefore report the equivalent width of the ourchoicesof these points. These are notedTable 1 and NC and the BC separately for the two He I lines, the we use our different measurement trials to estimate the sumgivingthetotallineequivalentwidth. Decomposing the equivalent width uncertainty for these profiles. For thelinesinthismannerrequiressomeassumptionstobe the other measurements,the uncertaintiesare computed made about how to separate the two components. Since bypropagatingtheuncertaintiesintheobservedspectra. the NC often appears asymmetric (e.g. Figure 1) with a very steep blue edge and shallower red edge, Gaussian 3.2. The Photospheric Mean Longitudinal Field fitting tothelines requiresparticularchoicestobemade on just how to do the analysis. For example, Batalha et For eachof the T Tauristars, we measuredthe photo- al. (1996)define(byeye)aregionoutsidetheNCandfit sphericB usingapproximately40magneticallysensitive z a single Gaussianto the resulting BC and subtract it off absorptionlines(Table2),whichformprimarilyoverthe in order to measure the NC equivalent width. Another portions of the stellar surface that are at photospheric procedureistofittheentirelinewithmultipleGaussians temperatures. These lines mayhaverelativelylittle con- and use the resulting fit parameters to estimate compo- tribution from the cool spots that are likely present on 5 these stars. Due to the wavelength dependence of the spheric lines. As a result, using equation (1) is strictly Zeeman effect and the fact that the signal-to-noise ra- valid only for weak lines, but give good results for real tio achieved in the observations of these late-type stars spectral lines (Mathys 1991). For each measurement of is considerably higher in the red regions of spectrum, thecenterofgravityofaspectralline,weusedthelocally we focus the analysis here only on the spectra from the measuredsignal-to-noiseinthespectrumtoestimatethe red CCD of HARPSpol. Lines for the analysis were se- uncertainty in the spectrum and then used standard er- lected by visual inspection of all the orders on the red ror propagation to find the uncertainty in the line shift CCD recorded with HARPS. Lines were deemed good bewteen the two polarizationstates and the implied B . z for the analysis if they appeared relatively strong (cen- Our final estimate of B is a weighted mean of the in- z tral depth ∼> 0.15) in the observed spectrum (though dividual line estimates, and these means are reported in mostwereconsiderablystronger),appearedfreeofblend- Table 3. ing by other photospheric lines, and were not contami- Werepeatedthemeasurementsofthelinewavelengths natedby telluric absorption. Lines passingthese criteria in the two polarization components and the resulting fi- were then checked in the Vienna Atomic Line Database nalvalueofBz anumberoftimes,makingslightlydiffer- (VALD,Kupkaetal. 1999,2000)andifthey arepresent entchoices onintegrationlimits for the center of gravity and have a value for the effective Land´e g-factor, the estimate of the wavelength of each line. However, for line was used in the analysis. In a few cases, the VALD eachtrial we alwaysused the same integrationlimits for data indicated that an apparently good line is actually a given line when analyzing the RCP and LCP compo- a very close blend of two lines. In this case, we used nentsofthespectra. Ineachcase,weachievedconsistent the line but estimated a new effective Land´e g-factor by resultswithinourquoteduncertainties. Generally,wedi- calculating the weighted mean of the effective Land´e g- videdourtrialsintwogroups. Inthefirstcase,wechoose factors of the lines in the blend. The weights used are integrationlimits veryclose to where the lines appearto the central depth of each component line as predicted reach the local continuum. This was primarily done as by VALD for the atmospheric parameters typical of K7 an effort to exclude any potential weak line blends that TTS (T = 4000 K, logg = 3.5). The initial line list might appear as a small distortion in the line wings. In eff was constructed using a visual examination of the spec- thesecondgroup,wechooseintegrationlimitsclearlyout trumofGQLupobtainedon29April2010. Fortheother inthelocalcontinuum,butwhichinsomecaseslikelyin- TTSs some lines were affected by blending with telluric cluded some weak line blends. In many cases, choosing absorption or by strong cosmic ray hits (as is also the the wider limits produced higher values of Bz, though case for later observations of GQ Lup). In these cases, in some cases the measured field went down. On aver- the lines were not included in the determination of the age, the wider bins resulted in fields stronger by ∼35%. photospheric B values. Lines so affected are noted in Choosing the wider limits does generally result in some- z Table 2. what larger uncertainty estimates, typically by a factor Once the line list was determined, the mean longitu- of 1.8. Again, the two groups of results are consistent dinal magnetic field, B , canbe estimated by measuring within these uncertainties (the difference typically being z the wavelength shift of each line, ∆λ = λ −λ , where ∼ 1.5σ). In Table 3 we quote the values for the wider R L λ is the wavelength of the line observed in the RCP integrationlimits with their correspondinglygreater un- R component of the spectrum and λ is the wavelength certainties. L measured in the LCP component of the spectrum (Bab- Examiningthephotosphericfieldsandtheiruncertain- cock 1962). The shift of the line observed in the two ties as reported in Table 3, apparently significant fields polarization states is related to B by arefoundonbothTWHyaandGQLupeachnightthey z wereobserved. However,thevalueofB foundonV2129 z ∆λ=2 e λ2g B =9.34×10−7λ2g B m˚A Oph is less than 1σ and does not represent a real detec- 4πm c2 eff z eff z tion. As a test of our measurement techniques and in e (1) order to gain confidence in our uncertainty estimates, where g is the effective Land´e g-factor of the transi- two different null tests were performed on the observa- eff tion, B is the strength of the mean longitudinal mag- tionsofeachtarget. Eachtestshouldreturna measured z netic field in kilogauss, and λ is the wavelength of the value of B = 0, thus serving to test how accurately we z transitioninAngstroms (Babcock 1962;see also Mathys can recoverthis value and whether the uncertainties are 1989, 1991). In order to measure the wavelength shift, realistically estimated. As described above, the spectra ∆λ, we measured the wavelength of each line of interest from the 4 subexposures were also combined in such a inthetwocircularpolarizationcomponentsusingtheso- way as to produce a null Stokes V profile (e.g. Donati calledcenterofgravitymethod(e.g. Mathys1989,1991; et al. 1997; Bagnulo et al. 2009), which can be used to Plachinda&Tarasova1999). Thismethodforestimating calculatenullRCPandLCPspectra. Weanalyzedthese B is mathematically equivalent to estimating B from null spectra in the same as described above. This has z z the firstorder momentof the Stokes V profile,assuming the advantage of using exactly the same lines as used to the same integration limits are used in the two methods measureB ,withexactlythe samewavelengthlimitsfor z (e.g. Borra&Vaughn1977;Mathys1989;LandiDeglIn- computingtheshiftofeachline,butwithdatacombined nocenti 2004). Determining B requires knowledge of inadifferentwaythanintherealmeasurement. Thesec- z the effective Land´e g-factor, g , for the transition. The ond null test we employed used 15 strong telluric lines eff weights for individual π and σ components of a given from the atmospheric B band of oxygen between 6883 - spectrallinewhichgointothedefinitionofgeff assumean 6910 ˚A. Since these lines should show no significant po- optically thin line, soequation(1) is only approximately larization of their own, we can analyze them in exactly true in the case of moderately strong (saturated) photo- 6 the same fashion as we do the stellar lines from which of gravity technique to measure the wavelength shift of we derive B (that is we take ∆λ=λ −λ as done for the line as observed in the two polarization components z R L the stellar measurements) combining the data from the (RCPandLCP).OurmeasurementsofB anditsuncer- z sub-exposures in exactly the same way as done for the tainty for both HeI emission lines are reported in Table stellarmeasurements. Inordertotranslatethemeasured 3 for each star on each night. When measuring the field shiftstoavalueofB weassignaLand´eg-factortoeach intheHeIformationregion,caremustbegivenwhense- z telluric line equal to the mean value of the photospheric lecting the wavelengthlimits for determing the center of lines usedfor the givenobservation. The weightsarethe gravitywavelengthofthelineandalsoindecidinghowto uncertainty on the value of B derived from each of the separate the narrow (NC) and broad (BC) components z photospheric lines. These null test field values are also oftheline(§3.1). LookingattheHeI5876˚AlineofTW reported in Table 3. HyainFigures1and2,polarizationisclearlyseeninthe NC of the line, but is not obviously apparent in the BC 3.3. B in the Accretion Shock Emission z extending off to the red side of the line. As a result, we As described above, Johns–Krull et al. (1999a) dis- focussed in on the NC of the line when measuring the covered that the He I 5876 ˚A emission line can be cir- value of B in the 5876 ˚A line of HeI. z cularly polarized in spectra of CTTS, implying coherent In addition to the specific wavelength region chosen, magnetic fields at the footpoints of accretion columns. care must also be taken when defining the local contin- They measured Bz = 2.46±0.12 kG for BP Tau. Since uum to be used when measuring the center of gravity this original discovery, polarization in the He I 5876 ˚A wavelength for the line. The reason for this is that the emission line has now been reported for a number of center of gravity technique (as well as the first moment CTTSs by a number of investigators (Valenti & Johns– of Stokes V) is an intensity weighted mean wavelength, Krull 2004; Symington et al. 2005; Smirnov et al. 2004; where the intensity used is that above the continuum in Yang et al. 2007; Donati et al. 2007, 2008, 2010b; Chen the case of an emissionline. In the case of the HeI 5876 & Johns–Krull 2011). Polarization has since been de- ˚Aemissionline,the NC ofthe linesits ontopofaBC in tected in other emission lines (notably He I 6678 ˚A and manycasesasdiscussedabove. Inthiscase,significantly theCaIIIRTlines)thoughttobe associatedwithaccre- different results are obtained if the stellar continuum is tion shock emission as well (Yang et al. 2007; Donati et usedcomparedto whatisobtainedifasomewhathigher al. 2007, 2008, 2010b; Chen & Johns–Krull 2011). The continuum defined by the BC is used. We proceed in spectra obtained here contain both the HeI 5876 ˚A and an effort to isolate the emission from the NC and mea- 6678 ˚A lines, so we analyze them with a focus on trying sure Bz in this component. Interpreting the NC of the He I emission as an excess line emitted from a distinct to see how well the fields derived from each line agree regionthat adds its emission to that from both the stel- with one another as this could provide clues to the loca- lar continuum and the BC of the He I emission, these tionandgeometryoftheaccretionshocksonthestar. As additional sources should be subtracted off when mea- mentioned earlier, the 5876 ˚A line of He I is composed suring the center of gravity wavelengthof the NC of the of several components (6) which are closely spaced in emission line. In §3.1 we described several methods of wavelength, and as a result is subject to the Paschen- separating the NC and BC when measuring their equiv- Back effect (e.g. Yang et al. 2007; Asensio Ramos et alentwidth, showingthat eachmethod is subject to cer- al. 2008). Therefore, the exact splitting pattern of the tain biases, but that the resulting systematic differences lines can vary considerably, depending on the strength weresmall(∼5%forthe methodwesettle on). We used of the magnetic field. However, since most of the level the same methods to remove the BC from the line (each crossing and merging has occurredby the time the local BC is shown in the bottom of Figure 1) and measured field strength reaches 2 kG (e.g. Asensio Ramos et al. the resulting B . For the profile shown in Figure 1, us- 2008), the HeI 5876 ˚A line should be in or very close to z ing the single Gaussian fitted to the region outside the thecompletePaschen-Backlimitgiventhefieldstrengths NC gave the largest magnitude field (−2.57±0.11 kG), we recover below. As a result, we set g =1.0 for this eff while the two Gaussian fit and the linear interpolation line as done in Yang et al. (2007). We can then test both gave −2.38±0.10 kG, and the 3 Gaussian fit gave the validity of this assumption once we have our field −2.40±0.11 kG. All 4 methods give results consistent measurements. to within 2σ and three of the 4 differ by 0.02 kG or less. Figure 2 shows the right- and left-circularly polarized As in the case for the equivalent widths above, we again componentsoftheHeI5876˚AemissionlinesofGQLup adoptthelinearinterpolationundertheNCasthewayof and TW Hya as observed on 29 April 2010. Also shown removingthe BC and use the resulting NC to determine inthefigurearetheStokesV profilesofthelines. Nearby the B values given in Table 3. z photospheric absorption lines are also seen in each star. As mentioned earlier, there have been reports in the Since the polarization signal in the photosphere is quite literature that the He I 6678 ˚A line shows substantially weak due to the low value of B present there (Table z stronger polarization than the 5876 ˚A indicative of a 3),these individualphotosphericabsorptionlines do not stronger local field in the line formation region of this show obvious polarization. They do serve to show that more optically thin line. We looked for evidence of this thetwopolarizationcomponentsarewellalignedinwave- effect in our stars by analyzing the the 6678 ˚A emission length,sothatthe obviousshiftoftheHeIemissionline line in all 3 of the T Tauri stars oberved here in a way between the RCP andLCP components indicates a very similar to how the the 5876 ˚A line is analyzed. How- strong field in the line formation region. ever, it became immediately apparent that additional In order to measure the value of B in the line for- mation region of the He I line we agazin use the center care needs to be taken when analyzing the 6678 ˚A line. 7 Figure 3 illustrates the situation for GQ Lup with the ionization stage, it should be possible to “fit” the 6663 spectrum obtained on the first night of our observing ˚A feature and then predict the strength of the 6678 ˚A run. Shown in the figure is the right and left circularly line. We use the spectrum synthesis code SYNTHMAG polarized components of the spectrum, along with the (Piskunov 1999). The line data for this spectral region Stokes V profile. Unlike the 5876 ˚A line where the right is taken from VALD (Krupka et al. 1999, 2000). As is and left circularly polarized components of the emission often the case,the initial predicted line strengths do not line have essentially the same shape, only shifted, the matchupwellwithobservationsoftheTTSsortheSun, 6678 ˚A line of GQ Lup has a different shape in the two so we tuned the oscillator strengths (and in a few cases circular polarization components. On the other hand, the Van der Waals broadening constants) of the strong the Stokes V profile for the two He I lines looks quite lines until a good match with the solar atlas (Kurucz et similar. The reason this is the case is the presence of an al. 1984) was obtained. In tuning the line parameters FeIphotosphericabsorptionlineatthesamewavelength tothe solarspectrum, weuse the modelatmosphereand of the 6678 ˚A line which can be seen in the spectrum of associated parameters (vsini, macroturbulence, etc.) for V2129 Oph (taken on night 2) which is also shown in theTeff =5731KscaledKuruczmodelfromValentiand Figure 4. In V2129 Oph, the He I 6678 ˚A line shows Piskunov(1996). Once the line data is set, we then syn- thesize this same spectral region for each of our CTTSs no obviousemission, but it turns outthere is some weak usingNextGenmodelatmospheres(Allard&Hauschildt emission in this line in V2129 Oph that partially fills in 1995). To do so, we must select an effective tempera- the photospheric absorption line and produces a clear ture, gravity, and metallicity for each star. We take the Stokes V signature which is illustrated in Figure 4 and gravity as logg =3.5 and metallicity as [M/H] = 0.0 for discussed below. all stars; however, we pick effective temperatures from As described above, the formation region for the nar- row components of these He I lines is thought to be at the standard NextGen grid that are as close as possible to that for each star. Donati et al. (2007) adopt an the base of the accretion columns where material from effective temperature of 4500±200 K for V2129 Oph, the disk is raining down onto the star. In addition to and we adopt the NextGen model with T = 4600 K producingsomeemissionlines,theseaccretionfootpoints eff for this star. Yang et al. (2005) find T = 4126±24 are believed to be the source of the optical continuum eff K for TW Hya, and we adopt the NextGen model with veiling seen in most CTTSs (e.g. Basri & Batalha 1990; T = 4200 K. GQ Lup has the same spectral type as Hartiganetal. 1991;Valentietal. 1993;Gullbringetal. eff TW Hya (K7), so we use the same NextGen model for 1998). Thisextralineandcontinuumemissionregionef- this star as well. In order to do the final spectrum syn- fectively blocks some small portion of the stellar surface thesis and fitting to the observed profile, we must select and the light it emits, adding its own emission on top the microandmacroturbulenceaswellasthe vsinirota- of the stellar spectrum coming from the non-accreting tional velocity. Following Johns–Krull et al. (1999b) we regions of the star (e.g. Cavet & Gullbring 1998). The spectrum of the excess can then be studied by subtract- take 2.0 km s−1 for the macroturbulent velocity and set ing an appropritelyscaled(veiled) stellartemplate spec- the microturbulence to 1.0 km s−1. In order to account trum from the observed CTTS spectrum (e.g. Hartigan forrotation,wesetvsini=4.0kms−1 forTWHya(Do- et al. 1995; Gullbring et al. 1998; Stempels & Piskunov nati et al. 2011b),we set vsini=14.5 km s−1 for V2129 2003). The spectrum of the excess is sometimes stud- Oph(Donatiet al. 2007),andwe set vsini=6.8 kms−1 ied by subtracting off the spectrum of a veiled WTTS for GQ Lup (Guenther et al. 2005). or main sequence star of the same spectral type, while Using the spectrum synthesis described above, we other studies use a synthetic spectrum computed from a match the strength of the Fe I feature at 6663 ˚A and model stellar atmosphere. Here, we use a model stellar subtractthe model spectrum from the rightand left cir- atmosphere since we did not observe a suitable WTTS cularly polarized components of the observed spectrum, having no excess continuum or line emission of its own. add a pseudo-continuum of 1.0 back in, and then follow The top panelofFigure 3 showsa portionof the spec- theproceduredescribedabovefortheHeI5876˚Alinein trumofV2129OphintheneighborhoodoftheHeI6678 order to measure the field in the formationregion of the ˚A line with our best fit synthetic spectrum including a 6678˚AHeIline. The onlyfreeparameterusedtofitthe modestamountofveiling(r =0.17,thatis,acontinuum 6663 ˚A FeI is the value of the continuum veiling. These excessequalto0.17ofthelocalstellarcontinuum)added veiling values and the B values for the 6678˚A emission z in. Thisvaluefortheveilingisslightlyhigherthanprevi- linearereportedinTable1. Ingeneralourveilingvalues ouslyreportedvaluesforV2129Oph: r =0.0±0.1(Basri agreewellwithpreviouslypublisheddeterminationscon- &Batalha1990);r =0.075±0.020(Donatietal. 2011a, sidering this quantity is often quite variable in CTTSs. thoughthisisonlyarelativeveilingvalueandrepresents In addition to previous veiling measurements for V2129 a lower limit to the true value). For our purposes, the Ophdiscussedabove,veilingonTWHyainsimilarspec- exactvalue of the veiling is not important. The primary tral regions has been shown to vary, reaching as high as goal with such a fit is to predict the strength of the FeI 0.80 in the study of Alencar and Batalha (2002) and as absorption line at 6678 ˚A. To do so, we synthesize the high as ∼ 0.92 in the work of Donati et al. (2011b). In spectral region between 6660 – 6680 ˚A as shown in Fig- the case of GQ Lup, Weise et al. (2010) found a veiling ure 4. The two strongest features in this region are the of 0.5 near the Li I line at 6707 ˚A, while here we find FeIfeatureat6678˚AandanotherFeIfeatureat6663˚A. values ranging from 0.3 to 0.6. Again though, the exact BothofthesefeaturesseeninFigure4areactualyblends value ofthe veilingis ofsecondaryimportance: account- of two Fe I lines. However, since all the components of ing for the photospheric absorption which is coincident both features are from the same element in the same 8 with the He I 6678 ˚A emission makes a substantial dif- of the He I emission line in their field determinations ference in the recovered field strengths for this line. To (Yang,privatecommunication). Thefieldmeasurements gauge this effect, we repeated the analysis of the 6678 ˚A reportedhere in Table 3 for B in the HeI 5876˚A emis- z line without making any correction for the photospheric sion line generally agree well with those of Yang et al. absorption present. Generally, the fields we measure in andaresignificantlylessthanmostofthe fieldsreported this case are a factor of two larger than those reported by Donati et al. However, Yang et al. used a cross cor- in Table 3. relation technique to measure line shifts and resulting field strengths, while here we use the center of gravity 4. DISCUSSION technique to measure line shifts. Examiningboth the null StokesV field values and the In order to verify that the difference in measurement techniquedoes notintroduce aspuriousdifference inthe field values obtained from the analysis of the telluric recovered field, we reanalyzed the Yang et al. (2007) lines shows that each of these is quite low and generally spectrum of TW Hya obtained on 21 April 1999 (the equal to zero (as they should be) within the measured strongestHeI field they found) using the same center of uncertainties. Inthe caseofthenullspectrum,thevalue gravitytechniqueemployedabove. Briefly,thisspectrum reachesasignificanceof3.3σinonecase,butmostvalues arebetween1−2σ. For the telluric lines,allofthe mea- was obtained with the 2.7 m Harlan J. Smith telescope suredB valuesareequaltozero(asexpected)towithin of McDonald Observatory used to feed the Robert G. z Tull coude echelle spectrometer (Tull et al. 1995). The 1σ or less. The uncertainties from the telluric tests are starlight was split into its circularly polarized compo- also generally lower than the for the photospheric lines nents using a Zeeman Analyzer (ZA, Vogt et al. 1980) (bothactualmeasurementandnulltest). Thisisprimar- placed in front of the spectrometer slit. The ZA con- ily due to the telluric lines being very sharp and strong, tainsaBabinet-Soleilphasecompensatorusedtocorrect allowingforveryaccuratemeasurementsofanypotential for a potential reduction in sensitivity to circular polar- wavelength shift. We conclude from these tests that our ization which can be introduced by the non-azimuthally uncertainties in B are generally well characterized and z symmetric reflections in the coude mirror train. Such that our detections of polarization in the photospheric lines and resulting measurements of B on GQ Lup and bounces can introduce some linear polarization into an z originally circularly polarized beam and the phase com- TW Hya are real. pensator is used to convert this linear polarization back TW Hya has been studied with spectropolarimetryby into the original circular polarized signal (see Vogt et Yang et al. (2007) and Donati et al. (2011b). On one al. 1980 and Vogt 1978). The exposure of TW Hya for of six nights, Yang et al. (2007) measured polarization that night totalled 4300 s. Yang et al. (2007) measured in a dozen photospheric lines of TW Hya correspond- ing to a longitudinal magnetic field of B = 149± 33 a field of −1806±114 G for the narrow component of z G,while finding no significantpolarization(thoughwith the He I 5876 ˚A line using a cross correlation analysis, larger uncertainties) on the other five nights. Donati et whilewefindafieldof−2326±118Gusingthecenterof al. (2011b) observed TW Hya for a total of 20 nights gravitytechnique,fully consistentwithourcurrentmea- and used their least squares deconvolution (LSD) anal- surements. This measurement of the field is stronger by ysis (Donati et al. 1997) to measure photospheric fields 520±164 G than that of Yang et al. (2007), represent- ranging from B =380−700 G with uncertainties typi- ing a 3.2σ difference. The difference is driven both by z cally of 15 G. Donati et al. suggest that long term tem- the method of estimating the line shift (center of grav- poralvariability may be responsible for the variations in ity versus cross correlation) as well as in the treatment B between the measurements of Yang et al. (taken in of removing the broad component. As a result, it does z April 1999) and those of Donati et al. (taken in March notappearthatdifferencesinthemeasurementtechnique 2008andMarch2010). The results presentedhere, from canaccountforthe variationsinthe fieldstrengthinthe AprilandMay2010,agreewellwiththose ofYang etal. He I emission line formation region found between Do- and less so with those of Donati et al. (2011b). nati et al. (2011b) and this study plus that of Yang et Yang et al. (2007) measured polarization in the He I al. (2007). It is more likely that intrinsic variability is 5876˚A and6678˚A emissionlines of TW Hya,as well as at work in this accretion diagnostic, and indeed, Donati etal. (2011b)findsignificantdifferencesinthe HeI field intheCaII8498˚Aemissionline. Yangetal. reportpos- strengthfromonerotationphasetothenextinTWHya, siblerotationalmodulationofthepolarizationintheHeI while the bulk of their field measurements are larger in 5876˚Aline,withimpliedfieldsvaryingfrom∼−1450to magnitude than either Yang et al. (2007) or those here. ∼ −1800 G. Donati et al. (2011b) also measure strong One of the motivations for this study was to verify polarizationintheHeI5876˚Alinewithvariationsagain and expand on the intriguing result that the field mea- suggestive of rotational modulation. The field strength sured from the He I 6678 ˚A line is significantly different implied in this case ranges from ∼ −2000 to ∼ −3500 (stronger) than that measured in the He I 5876 ˚A line G. Donati et al. attribute the significant difference be- (e.g. Donati et al. 2008). In principle, measuring the tween their He I measurements and those of Yang et al. field in different emission line diagnostics could offer a to perceived differences in the way the fields were mea- means for probing the magnetic field structure through sured in the two studies. Donati et al. measured the the accretion shock on CTTSs. As described in §3.3, field only in the narrow component of the line for simi- we discovered that there is a strong photospheric Fe I lar reasons to those cited above in §3.3, and they state absorption line coincident in wavelength with the He I that Yang et al. (2007) used the entire HeI line in their 6678 ˚A line that can severely affect the field strength fielddetermination;however,thisisnotcorrect. Yanget measured in this emission line if the photospheric line al. (2007) also focussed only on the narrow component 9 is not properly accounted for. As described above, we strongfieldsinthislinefromafewCTTSs;however,also attempted to correct for the Fe I photospheric line by with substantial error bars. Higher precision measure- computing veiled model spectra for each of our observa- ments(withuncertaintiesfrom∼0.1−0.4kG)havepeak tionsandsubtractedthe resultingmodelfromthe obser- fields measured in the HeI 5876 ˚A line of |B |<3.4 kG z vations before measuring the field. Doing so produced (Donati et al. 2011) with peak values typically |B |∼2 z field measurements in the HeI 6678˚A line that are very kG depending on the specific star in question (Johns– consistentwith those measured in the 5876˚A line of the Krull et al. 1999a; Valenti & Johns–Krull 2004; Yang same element. The field strengths recovered from the et al. 2005; Donati et al. 2007, 2008, 2010b). Very two lines are the same to within 3σ for all observations. recently, Donati et al. (2012) present polarimetric mea- If we do not account the Fe I photospheric absorption surementsofGQLupattwodifferentepochs(July2009, line, we generally recover a field in the 6678 ˚A line that June 2011). While the photospheric fields they detect is twice as strong as reported in Table 3. In a separate doshowsomevariationfromoneepochtothe other,our test,wecorrectedforthephotosphericabsorptionlinein measurements over about half the rotation period found the observations of TW Hya and GQ Lup using a veiled by Donati et al. are consistent with the fields found at versionoftheobservedV2129Ophspectrum. Whilethis both epochs. On the ther hand, Donati et al. find a is not ideal due to the weak He I 6678 ˚A emission from substantial decrease in the large scale field controlling V2129 Oph, the results were the same as those in Ta- the accretion over the two years between their epochs. ble 3 within the uncertainties. As a result, it appears Our data (April 2010) is fully consistent with their ear- that the field in the two He I lines is essentially identi- lier epoch, with our He I Bz determinations matching cal, at least in TW Hya and GQ Lup, but that special thestrongestvaluestheyobserveinJuly2009,providing care must be taken when analyzing the 6678 ˚A line. We additional constraints on the timescale involved in the magnetic field change observed by Donati et al. also note from earlier that our analysis of the 5876 ˚A line assumed it formed in the complete Paschen-Back regime (i.e. g = 1.0). Since shock models suggest eff the field strengthin these two lines shouldessentially be We wish to thank ESO staff in Santiago and on La the same, and that is indeed what we find, these results Silla for their hospitality and help during the observing alsosuggestthat ourtreatment ofthe 5876˚A line in the run there. This research has made use of the Simbad complete Paschen-Backeffect is appropriate,at least for Astronomical database, the VALD line database, and the strong magnetic fields recoveredhere. the NASA AstrophysicsData System. CMJ-Kwishes to The magnetic field on GQ Lup appears to be quite acknowledge partial support for this research from the remarkable. The field detected in the photospheric ab- NASA Origins of Solar Systems program through grant sorption lines is fairly typical in magnitude (|B | ∼ 200 numbers NNX10AI53G made to Rice University. OK is z G) relative to the strength detected on many CTTSs. a Royal Swedish Academy of Sciences Research Fellow, However, the field (∼ 6 kG) detected in the He I supported by grants from Knut and Alice Wallenberg line formation region, the accretion shock, is substan- Foundation and Swedish Research Council. Finally, we tially stronger in magnitude than that observed in most thank an anonymous referee for helpful suggestions to CTTSs. Symington et al. (2005) reported some quite improve the presentation in the manuscript. REFERENCES Alencar,S.H.P.,&Basri,G.2000, AJ,119,1881 Camenzind, M. 1990, Rev. Modern Astron., (Berlin: Springer- Alencar,S.H.P.,&Batalha,C.2002,ApJ,571,378 Verlag),3,234 Alencar,S.H.P.,Johns–Krull,C.M.,&Basri,G.,2001,AJ,122, Cauley, P.W.,Johns-Krull,C.M.,Hamilton,C.M.,& Lockhart, 3335 K.2012,ApJ,756,68 Allard, F. & Hauschildt, P.H. 1995, in Bottom of the Main Chen,W.&Johns–Krull,C.M.2011,ASPConf.Series,inpress SequenceandBeyond,ed.C.G.Tinney(Berlin:Springer),32 CollierCameron,A.&Campbell,C.G.1993,A&A,274,309 AsensioRamos,A.,TrujilloBueno,J.,&LandiDegl’Innocenti,E. Daou, A. G., Johns–Krull,C. M.,& Valenti, J. A.2006, AJ, 131, 2008,ApJ,683,542 520 Azevedo,R.,Calvet,N.,Hartmann,L.,Folha,D.F.M.,Gameiro, Donati, J.-F., Semel, M., Carter, B. D., Rees, D. E., & Collier F.,&Muzerolle,J.2006,A&A,456,225 Cameron,A.1997,MNRAS,291,658 Babcock, H. W. 1962, in Astronomical Techniques, ed. W. A. Donati,J.-F.,etal.2007,MNRAS,380,1297 Hiltner(Chicago:Univ.ofChicagoPress),p.107 Donati,J.-F.,etal.2008,MNRAS,386,1234 Bagnulo,S.,Landolfi,M.,Landstreet, J.D.,LandiDeglInnocenti, Donati,J.-F.,etal.2010a,MNRAS,402,1426 E.,Fossati,L.,&Sterzik,M.2009,PASP,121,993 Donati,J.-F.,etal.2010b,MNRAS,409,1347 Basri,G.,&Batalha,C.1990,ApJ,363,654 Donati,J.-F.,etal.2011a,MNRAS,412,2454 Basri,G.,Marcy,G.W.,&Valenti,J.A.1992,ApJ,390,622 Donati,J.-F.,etal.2011b,MNRAS,417,472 Batalha,C.,Lopes,D.F.,&Batalha,N.M.2001,ApJ,548,377 Donati,J.-F.,etal.2012,MNRAS,425,2948 Batalha,C.C.,Stout–Batalha,N.M.,Basri,G.,&Terra,M.A.O. Edwards,S.,Hartigan,P.,Ghandour,L.,&Andrulis,C.1994,AJ, 1996,ApJS,103,211 108,1056 Beristain,G.,Edwards,S.,&Kwan,J.2001,ApJ,551,1037 Feigelson,E.D.&Montmerle,T.1999, ARAA,37,363 Bertout,C.1989, ARAA,27,351 Ferreira,J.2008, NewAstronomyReview,52,42 Bertout,C.,Basri,G.,&Bouvier,J.1988, ApJ,330,350 Gregory, S. G., Matt, S. P., Donati, J.-F., & Jardine, M. 2008, Bertout, C., Harder, S., Malbet, F., Mennessier, C., & Regev, O. MNRAS,389,1839 1996,AJ,112,2159 Guenther,E.W.,Lehmann,H.,Emerson,J.P.,&Staude,J.1999, Borra,E.F.&Vaughan, A.H.1977,ApJ,216,462 A&A,341,768 Bouvier,J.,etal.1999, A&A,349,619 Guenther, E. W., Neuha¨user, R., Wuchterl, G., et al. 2005, Bouvier, J., Alencar, S. H. P., Harries,T. J., Johns-Krull,C. M., AstronomischeNachrichten, 326,958 &Romanova,M.M.2007,ProtostarsandPlanets V,479 Gullbring,E.,Hartmann,L.,Bricen˜o,C.,&Calvet,N.1998,ApJ, Calvet,N.,&Gullbring,E.1998,ApJ,509,802 492,323 10 Hartigan,P.,Edwards,S.,&Ghandour,L.1995, ApJ,452,736 Muzerolle, J., Calvet, N., Brice˜eno, C., Hartmann, L., & Hartigan,P.,Kenyon,S.J.,Hartmann,L.,Strom,S.E.,Edwards, Hillenbrand,L.2000, ApJ,535,L47 S.,Welty, A.D.,&Stauffer,J.1991,ApJ,382,617 Muzerolle,J.,Calvet,N.,&Hartmann,L.1998,ApJ,492,743 Hartmann,L.,Hewett, R.,&Calvet,N.1994,ApJ,426,669 Muzerolle,J.,Calvet,N.,&Hartmann,L.2001,ApJ,550,944 Herbig, G. H., & Bell, K. R. 1988, Third catalog of emission- Neuha¨user, R.,Guenther, E.W.,Wuchterl, G.,et al.2005, A&A, line stars of the Orion population., by G.H. gerbig and 435,L13 K.R.Bell. LickObservatory Bulletin#1111, Santa Cruz:Lick Neuha¨user, R., Mugrauer, M., Seifahrt, A., Schmidt, T. O. B., & Observatory,Jun1988,90p. Vogt,N.2008, A&A,484,281 Herbst, W.,Bailer–Jones, C. A.L., &Mundt, R.2001, ApJ, 554, Paatz,G.&Camenzind,M.1996,A&A,308,77 L197 Piskunov,N.1999, Polarization,243,515 Hinkle, K., Wallace, L., Valenti, J., & Harmer, D. 2000, Visible Piskunov,N.E.,&Valenti,J.A.2002, A&A,385,1095 and Near Infrared Atlas of the Arcturus Spectrum 3727-9300 Plachinda,S.I.,&Tarasova,T.N.1999,ApJ,514,402 Aed. KennethHinkle,LloydWallace, JeffValenti,andDianne Romanova, M. M., Long, M., Lamb, F. K., Kulkarni, A. K., & Harmer.(SanFrancisco:ASP)ISBN:1-58381-037-4,2000., Donati,J.-F.2011,MNRAS,411,915 Hu´elamo,N.,Figueira,P.,Bonfils,X.,etal.2008,A&A,489,L9 Romanova,M.M.,Ustyugova,G.V.,Koldoba,A.V.,&Lovelace, Hussain,G.A.J.,etal.2009,MNRAS,398,189 R.V.E.2004, ApJ,610,920 Jayawardhana, R.,Mohanty,S.,&Basri,G.2003, ApJ,592,282 Johns,C.M.&Basri,G.1995,ApJ,449,341 Romanova,M.M.,Ustyugova,G.V.,Koldoba,A.V.,&Lovelace, Johns–Krull,C.M.2007,ApJ,664,975 R.V.E.2009, MNRAS,399,1802 Johns–Krull,C.M.&Gafford,A.D.2002,ApJ,573,685 Seifahrt,A.,Neuha¨user, R.,&Hauschildt, P.H.2007, A&A,463, Johns–Krull, C. M., Valenti, J. A., Hatzes, A. P., & Kanaan, A. 309 1999a,510,L41 SeperueloDuarte, E.,Alencar,S.H.P.,Batalha, C.,& Lopes,D. Johns–Krull,C.M.,Valenti,J.A.,&Koresko,C.1999b,ApJ,516, 2008, A&A,489,349 900 Setiawan,J.,Henning,T.,Launhardt,R.,etal.2008,Nature,451, Johns–Krull, C. M.,Valenti, J. A., & Saar, S. H. 2004, ApJ, 617, 38 1204 Shu,F.H.,Najita,J.,Ostriker,E.,Wilkin,F.,Ruden,S.,&Lizano, Kastner, J. H., Zuckerman, B., Weintraub, D. A., & Forveille, T. S.1994,ApJ,429,781 1997,Science,277,67 Siess,L.,Dufour,E.,&Forestini,M.2000,A&A,358,593 K¨onigl,A.1991,ApJ,370,L39 Smirnov,D.A.,Lamzin,S.A.,Fabrika,S.N.,&Chuntonov,G.A. Kenyon, S. J., Hartmann, L., Hewett, R., Carrasco, L., Cruz– 2004, AstronomyLetters,30,456 Gonzalez, I., Recillas,E., Salas, L.,Serrano, A., Strom, K.M., Snik, F., Jeffers, S., Keller, C., Piskunov, N., Kochukhov, O., Strom,S.E.,&Newton, G.1994,AJ,107,2153 Valenti,J.,&Johns-Krull,C.2008,Proc.SPIE,7014 Kochukhov, O.,Bagnulo, S.,Wade, G. A.,Sangalli,L.,Piskunov, Snik, F., Kochukhov, O., Piskunov, N., et al. 2011, Solar N., Landstreet, J. D., Petit, P., & Sigut, T. A. A. 2004, A&A, Polarization6,437,237 414,613 Stempels,H.C.,&Piskunov,N.2003, A&A,408,693 Kochukhov&Wade2010,A&A,513,13 Symington,N.H.,Harries,T.J.,Kurosawa,R.,&Naylor,T.2005, Kupka,F.,Ryabchikova, T.A.,Piskunov,N.E.,Stempels,H.C.,& MNRAS,358,977 Weiss,W.W.,2000,BalticAstronomy,vol.9,590 Tull,R.G.,MacQueen,P.J.,Sneden,C.,&Lambert,D.L.1995, Kupka, F.,Piskunov, N.E.,Ryabchikova, T.A.,Stempels H.C.,& PASP,107,251 Weiss,W.W.,1999,A&AS138,119 Valenti,J.A.1994,Ph.D.Thesis,UniversityofCalifornia,Berkeley Kurucz, R. L., Furenlid, I., Brault, J., & Testerman, L. 1984, Valenti,J.A.,Basri,G.,&Johns,C.M.1993,AJ,106,2024 National Solar Observatory Atlas, Sunspot, New Mexico: Valenti,J.A.&Johns–Krull,C.M.2004,Ap&SS,292,619 NationalSolarObservatory,1984 Valenti,J.A.,&Piskunov,N.1996, A&AS,118,595 Lamzin,S.A.1998,AstronomyReports,42,322 Vink,J.S.,Drew,J.E.,Harries,T.J.,Oudmaijer,R.D.,&Unruh, Landi DeglInnocenti E. & Landolfi, M., 2004, Polarisation in Y.C.2003, A&A,406,703 Spectral Lines. Vol. 307 of Astrophysics and Space Science Vink,J.S.,Harries,T.J.,&Drew,J.E.2005,A&A,430,213 Library,KluwerAcademicPublishers Vogt, S. S. 1978, Ph.D. Thesis, The Univ. of Texas at Austin, Long,M.,Romanova,M.M.,&Lovelace,R.V.E.2005,ApJ,634, Austin,TX 1214 Long,M.,Romanova,M.M.,&Lovelace,R.V.E.2008,MNRAS, Vogt,S.S.,Tull,R.G.,&Kelton,P.W.1980,ApJ,236,308 386,1274 Webb,R.A.,Zuckerman,B.,Platais,I.,etal.1999,ApJ,512,L63 Marois,C.,Macintosh,B.,&Barman,T.2007,ApJ,654,L151 Weise,P.,Launhardt,R.,Setiawan,J.,&Henning,T.2010,A&A, Mathys,G.1989,Fund.CosmicPhys.,13,143 517,A88 Mathys,G.1991,A&AS,89,121 White,R.J.,&Basri,G.2003,ApJ,582,1109 Mayor,M.,etal.2003,TheMessenger,114,20 Wichmann,R.,Bastian,U.,Krautter,J.,Jankovics,I.,&Rucinski, McElwain,M.W.,Metchev,S.A.,Larkin,J.E.,etal.2007,ApJ, S.M.1998, MNRAS,301,L39 656,505 Yang,H.,&Johns-Krull,C.M.2011, ApJ,729,83 Mohanty,S.,&Shu,F.H.2008,ApJ,687,1323 Yang,H.,Johns–Krull,C.M.,&Valenti,J.A.2005,ApJ,635,466 Morin,J.,Donati,J.-F.,Petit,P.,etal.2008, MNRAS,390,567 Yang,H.,Johns–Krull,C.M.,&Valenti,J.A.2007,AJ,133,73 Mugrauer,M.,&Neuha¨user,R.2005,AstronomischeNachrichten, Yang,H.,Johns-Krull,C.M.,&Valenti,J.A.2008,AJ,136,2286 326,701 Zanni,C.,&Ferreira,J.2009,A&A,508,1117