Magnetically Controlled Accretion on the Classical T Tauri Stars GQ Lupi and TW Hydrae

0.3 MB·English

by  C. M. Johns-Krull

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