AASTeX Accretion signatures in magnetic HAeBes

Linking signatures of accretion with magnetic field measurements – line profiles are not significantly different in magnetic and non-magnetic Herbig Ae/Be stars

Megan Reiter University of Michigan 311 West Hall, 1085 S. University Ave Ann Arbor, MI 48109-1107, USA Nuria Calvet University of Michigan 311 West Hall, 1085 S. University Ave Ann Arbor, MI 48109-1107, USA Thanawuth Thanathibodee University of Michigan 311 West Hall, 1085 S. University Ave Ann Arbor, MI 48109-1107, USA Stefan Kraus School of Physics, Astrophysics Group, University of Exeter, Stocker Road, Exeter EX4 4QL, UK P. Wilson Cauley Arizona State University John Monnier University of Michigan 311 West Hall, 1085 S. University Ave Ann Arbor, MI 48109-1107, USA Adam Rubinstein University of Michigan 311 West Hall, 1085 S. University Ave Ann Arbor, MI 48109-1107, USA Alicia Aarnio University of Colorado Tim J. Harries School of Physics, Astrophysics Group, University of Exeter, Stocker Road, Exeter EX4 4QL, UK
August 3, 2017November 3, 2017July 15, 2019
August 3, 2017November 3, 2017July 15, 2019
August 3, 2017November 3, 2017July 15, 2019

Herbig Ae/Be stars are young, pre-main-sequence stars that sample the transition in structure and evolution between low- and high-mass stars, providing a key test of accretion processes in higher-mass stars. Few Herbig Ae/Be stars have detected magnetic fields, calling into question whether the magnetospheric accretion paradigm developed for low-mass stars can be scaled to higher masses. We present He i 10830 Å line profiles for 64 Herbig Ae/Be stars with a magnetic field measurement in order to test magnetospheric accretion in the physical regime where its efficacy remains uncertain. Of the 5 stars with a magnetic field detection, 1 shows redshifted absorption, indicative of infall, and 2 show blueshifted absorption, tracing mass outflow. The fraction of redshifted and blueshifted absorption profiles in the non-magnetic Herbig Ae/Be stars is remarkably similar, suggesting that the stellar magnetic field does not affect gas kinematics traced by He i 10830 Å. Line profile morphology does not correlate with the luminosity, rotation rate, mass accretion rate, or disk inclination. Only the detection of a magnetic field and a nearly face-on disk inclination show a correlation (albeit for few sources). This provides further evidence for weaker dipoles and more complex field topologies as stars develop a radiative envelope. The small number of magnetic Herbig Ae/Be stars has already called into question whether magnetospheric accretion can be scaled to higher masses; accretion signatures are not substantially different in magnetic Herbig Ae/Be stars, casting further doubt that they accrete in the same manner as classical T Tauri stars.

stars: formation — others…

Megan Reiter

1 Introduction

Herbig Ae/Be stars (HAeBes) are the intermediate-mass analogs of low-mass T Tauri stars – young stars that are sufficiently evolved to be studied in the optical, but with IR colors that suggest remnant circumstellar material and spectroscopic signatures of accretion (Herbig, 1960; Hillenbrand et al., 1992; Waters & Waelkens, 1998). With masses between  M, HAeBes sample conditions of formation and evolution intermediate between low- and high-mass stars. Their stellar structure, multiplicity, pre-main-sequence (PMS) evolution times, and magnetic properties provide insight into the physical parameters that guide the formation of higher-mass stars. At the same time, they are more numerous than the highest-mass stars and evolve more slowly, allowing them to be studied in the optical and near-IR with techniques similar to those used for T Tauri stars (e.g., Mendigutía et al., 2012).

Extensive studies of nearby, low-mass PMS stars led to the development of a magnetically-controlled accretion paradigm (e.g., Calvet & Hartmann, 1992; Hartmann et al., 1994; Muzerolle et al., 1998b, 2001, 2004; Hartmann et al., 2016). In this picture a strong, predominately dipolar stellar magnetic field interacts with and truncates the circumstellar disk at a few stellar radii. Magnetic field lines loft material from the disk, guiding it toward the stellar surface where it splashes down at high latitudes. Models coevolved with observations to explain (1) high-velocity wings in absorption lines (e.g., Edwards et al., 1994; Bouvier et al., 1999; Muzerolle et al., 1998b), (2) excess continuum emission attributed to the accretion shock at the stellar surface (e.g., Calvet & Gullbring, 1998; Muzerolle et al., 1998b; Gullbring et al., 2000; Muzerolle et al., 2001), and (3) the detection of strong magnetic fields on T Tauri stars (e.g., Johns-Krull et al., 1999a; Johns-Krull, 2007; Johns-Krull et al., 2013).

Changes in the stellar structure above M make it unclear whether intermediate-mass stars can generate magnetic fields of sufficient strength to support magnetospheric accretion like their low-mass counterparts. In particular, stars above  M develop radiative envelopes before they disperse their disks, so they may not maintain strong, ordered magnetic fields throughout their evolution (e.g., Hussain et al., 2009; Gregory et al., 2012). Weaker fields may not be able to disrupt the disk, leading to smaller infall velocities or possibly a different accretion pathway altogether. Surveys of HAeBes find a low magnetic incidence, with fields detected in % of sources measured (see, e.g., Wade et al., 2007; Alecian et al., 2013b; Hubrig et al., 2013; Bagnulo et al., 2015). To make matters worse, derived upper limits on the magnetic field strength are smaller than the minimum field strength required for magnetospheric accretion in both Herbig Ae and Be stars (as derived from Johns-Krull et al., 1999b).

Despite the paucity of strongly magnetic HAeBes, observations suggest that a smooth scaling of the disk-mediated accretion models for low-mass T Tauri stars (e.g., Calvet et al., 2004; Muzerolle et al., 2004) can be applied to A-type stars (up to M). It is less clear that magnetospheric accretion models can be applied to higher-mass B-type stars. Recent spectroscopic studies testing whether empirical correlations between emission line luminosities and more direct tracers of the accretion rate extend to intermediate-mass stars find conflicting results. Donehew & Brittain (2011) find a break in the correlation between A- and B-type stars while both Mendigutía et al. (2011) and Fairlamb et al. (2015) find that fitting stellar parameters to each star individually produces a correlation that extends smoothly to B-type stars.

The line profiles themselves point to a different accretion geometry for earlier spectral types, if not a different physical mechanism altogether. Cauley & Johns-Krull (2014, 2015) report a smaller fraction of HAeBes showing redshifted absorption, indicative of accretion, or blueshifted absorption, tracing outflows, compared to classical T Tauri stars. Maximum velocities in the line profiles are smaller than free-fall, suggesting a more compact accretion geometry. Muzerolle et al. (2004) argued that magnetospheric accretion could proceed for A-type stars through small magnetospheres, since the faster rotation rates of higher-mass stars force corotation to smaller radii. Fewer redshifted absorption profiles in the Herbig Be stars compared to the Herbig Aes led Cauley & Johns-Krull (2014, 2015) to suggest that B-type stars may accrete via a boundary layer (e.g., Bertout et al., 1988; Basri & Bertout, 1989; Popham et al., 1993) rather than through the magnetosphere.

H spectropolarimetry also hints at different disk geometries around A- and B-type stars (Vink et al., 2002, 2005; Mottram et al., 2007; Ababakr et al., 2017). Line polarization is consistent with a gap in the inner disk of Herbig Ae stars, as expected if the magnetic field truncates the inner disk at a few stellar radii. Unlike the Herbig Aes, observations cannot rule out disks around B-type stars that extend to the stellar surface, permitting direct accretion onto the pre-main-sequence star.

Together, these observations illustrate a number of impediments to applying the standard magnetospheric accretion paradigm to HAeBes. Comparing the magnetic properties to the accretion behavior provides a key test of magnetospheric accretion. Whereas magnetospheric accretion has been examined in the context of the stellar magnetic fields measured in T Tauri stars (e.g., Symington et al., 2005; Johns-Krull, 2007), no such comparison has been made for a large sample of HAeBes. Fortunately, recent work by Alecian et al. (2013b) provides measurements and upper limits on the longitudinal magnetic fields of 70 HAeBes. In this paper, we present a combination of new and archival spectra of 64 HAeBes targeted for a magnetic field measurement (63/70 of the Alecian et al. (2013b) sample plus one source from the literature) to compare the line profiles of sources with and without evidence for a magnetic field.

Detailed magnetospheric accretion models can reproduce many of the observed features of the Balmer line profiles (e.g., Calvet & Hartmann, 1992; Hartmann et al., 1994; Calvet & Gullbring, 1998; Muzerolle et al., 1998a; Gullbring et al., 2000; Muzerolle et al., 2001; Kurosawa et al., 2006). Among the results of these models is the general prediction that thermalized lines are less likely to show redshifted absorption. Line luminosity correlations have been used to argue for a scaling of magnetospheric accretion (e.g., Muzerolle et al., 2004; Mendigutía et al., 2011), although it is unclear whether accretion dominates the line luminosity in HAeBes (e.g., Mendigutía et al., 2015; Fairlamb et al., 2017). Kinematics provide a more direct assessment of the motion of the gas, and thus a better diagnostic of accretion.

We examine He i 10830 Å line profiles for redshifted or blueshifted absorption profiles that provide direct or indirect evidence of accretion. Resonant scattering of permitted lines excited near the young star can produce absorption profiles that trace the kinematics of the gas. Blueshifted absorption is produced by colder gas moving toward the observer in a wind/outflow from an accreting young star (e.g., Finkenzeller & Mundt, 1984). Redshifted absorption occurs when colder material moves toward the star, and thus away from the observer. Redshifted absorption can only be produced by infall (e.g., Walker, 1972).

Several papers have explored He i 10830 Å as a tracer of the structure and kinematics of accretion and outflow in young stars (e.g., Takami et al., 2002; Edwards et al., 2003, 2006; Fischer et al., 2008; Kurosawa et al., 2011; Cauley & Johns-Krull, 2014). The lower level of the transition is 20 eV above ground, but metastable and therefore long-lived. Around PMS stars, densities do not tend to be high enough for collisional deexcitation and the lower level is radiatively isolated from the ground state (Kwan et al., 2007). The large optical depth, high emissivity, and metastable lower level of the line make it particularly susceptible to absorption (see, e.g., Kwan & Fischer, 2011), thereby making it a useful tracer of gas kinematics over a range of physical conditions.

In this paper, we examine the line profiles of He i 10830 Å for evidence of accretion in HAeBes that Alecian et al. (2013b) targeted for magnetic field measurements. We present profiles for 64 HAeBes using a combination of new near-IR spectroscopy from the Folded-port InfraRed Echellette (FIRE) spectrograph on Magellan, X-Shooter (VLT) data from the archive, and GNIRS/PHOENIX spectra from Cauley & Johns-Krull (2014). We compare the line profiles of the HAeBes with and without a magnetic field detection confirmed by Alecian et al. (2013b). Including disk inclinations estimated from near-IR H-band long-baseline interferometry by Lazareff et al. (2017), where available, we compare the line profiles with the physical properties of the stars. Together, these data allow us to probe the role that magnetic fields play in accretion onto HAeBes.

2 Observations

2.1 New FIRE spectra

We present new, near-IR spectra of 33 HAeBes obtained with the Folded-port InfraRed Echellette (FIRE) spectrograph (Simcoe et al., 2013) on the 6.5 m Baade/Magellan telescope. A single FIRE spectrum covers 0.8-2.5 µm with spectral resolution of . Data were obtained on three separate nights. Objects observed on 2011 March 10 or 12 used a wide slit () while those observed on 2016 September 20 used a wide slit (). Targets were observed using the standard ABBA sequence except for a few bright sources where only two nods were used. Wavelength calibration was done using a ThAr lamp. Spectra were reduced using the firehose IDL pipeline which performs flat-fielding, object extraction, and flux and wavelength calibration. Details for each source are listed in Table 1.

HD number Alt. Name Spectrograph Date Exp. time (s)
BD-06 1259 BF Ori 05:37:13.3 –06:35:01 PHOENIX 2013 Mar 2 800
BD-06 1253 V380 Ori 05:36:25.4 –06:42:58 FIRE 2011 Mar 12 3
BD-05 1329 T Ori 05:35:50.4 –05:28:35 PHOENIX 2013 Mar 01 800
BD-05 1324 NV Ori 05:35:31.4 –05:33:08 FIRE 2016 Sept 20 1
BD+41 3731 20:24:15.7 +42:18:01 PHOENIX 2013 Nov 12 600
BD+46 3471 V1578 Cyg 21:52:34.1 +47:13:44 PHOENIX 2013 Nov 11 600
BD+61 154 V594 Cas 00:43:18.3 +61:54:40 PHOENIX 2013 Mar 3 800
HD 17081  Cet 02:44:07.3 –13:51:31 PHOENIX 2013 Mar 3 90
HD 31293 AB Aur 04:55:45.8 +30:33:04 FIRE 2011 Mar 10 5
HD 31648 MWC 480 04:58:46.3 +29:50:37 FIRE 2011 Mar 12 3
HD 34282 V1366 Ori 05:16:00.5 –09:48:35 PHOENIX 2013 Mar 01 800
HD 35187 05:24:01.2 +24:57:38 GNIRS 2012 Dec 14 110
HD 35929 05:27:42.8 –08:19:38 X-shooter 2009 Dec 17 5
HD 36112 MWC 758 05:30:27.5 +25:19:57 FIRE 2011 Mar 12 3
HD 36910 CQ Tau 05:35:58.5 +24:44:54 PHOENIX 2013 Feb 27 600
HD 36917 V372 Ori 05:34:46.9 –05:34:14 FIRE 2016 Sept 20 1
HD 36982 LP Ori 05:35:09.8 –05:27:53 FIRE 2016 Sept 20 1
HD 37258 V586 Ori 05:36:59.1 –06:09:18 X-shooter 2010 Jan 02 10
HD 37357 05:37:47.1 –06:42:30 X-shooter 2010 Feb 05 10
HD 37806 MWC 120 05:41:02.3 –02:43:01 FIRE 2011 Mar 10 1
HD 38120 05:43:11.9 –04:59:49 FIRE 2016 Sept 20 2
HD 38238 V351 Ori 05:44:18.8 +00:08:40 FIRE 2016 Sept 20 1
HD 50083 V742 Mon 06:51:45.8 +05:05:04 GNIRS 2012 Dec 13 90
HD 52721 GU CMa 07:01:49.5 –11:18:03 FIRE 2016 Sept 20 2
HD 53367 MWC 166 07:04:25.5 –10:27:16 FIRE 2011 Mar 10 1
HD 68695 08:11:44.3 –44:05:08 X-shooter 2009 Dec 21 15
HD 72106 08:29:35.0 –38:36:19 X-shooter 2009 Dec 19 10
HD 76534 08:55:08.8 –43:27:57 X-shooter 2010 Jan 30 15
HD 98922 11:22:31.7 –53:22:11 FIRE 2011 Mar 10 10
HD 101412 11:39:44.3 –60:10:25 X-shooter 2010 Mar 30 10 4
HD 114981 V958 Cen 13:14:40.7 –38:39:06 GNIRS 2013 Jan 16 180
HD 139614 15:40:46.3 –42:29:51 X-shooter 2010 Mar 28 5
HD 141569 15:49:57.7 –03:55:16 FIRE 2011 Mar 12 20
HD 142666 15:56:40.0 –22:01:40 FIRE 2011 Mar 12 10
HD 144432 16:06:58.0 –27:43:10 FIRE 2011 Mar 12 5
HR 5999 16:08:34.3 –39:06:18 FIRE 2011 Mar 10 1
HD 145718 V718 Sco 16:13:11.6 –22:29:07 PHOENIX 2013 Feb 27 600
HD 150193 MWC 863 16:40:17.9 –23:53:45 FIRE 2011 Mar 10 1
HD 152404 AK Sco 16:54:45.0 –36:53:17 X-shooter 2009 Oct 05 5
HD 163296 17:56:21.3 –21:57:22 FIRE 2011 Mar 10 10
HD 169142 18:24:29.8 –29:46:49 FIRE 2016 Sept 20 2
HD 174571 MWC 610 18:50:47.2 +08:42:10 FIRE 2016 Sept 20 3
HD 176386 19:01:38.9 –36:53:26 FIRE 2016 Sept 20 1
HD 179218 MWC 614 19:11:11.3 +15:47:16 FIRE 2011 Mar 12 1
HD 190073 V1295 Aql 20:03:02.5 +05:44:17 FIRE 2011 Mar 12 3
HD 200775 MWC 361 21:01:36.9 +68:09:48 PHOENIX 2013 Feb 27 300
HD 216629 IL Cep 22:53:15.6 +62:08:45 PHOENIX 2013 Nov 9 600
HD 244314 V1409 Ori 05:30:19.0 +11:20:20 X-shooter 2010 Jan 02 15
HD 244604 V1410 Ori 05:31:57.2 +11:17:41 PHOENIX 2013 Feb 28 600
HD 245185 V1271 Ori 05:35:09.6 +10:01:52 X-shooter 2009 Dec 17 15
HD 249879 05:58:55.8 +16:39:57 FIRE 2016 Sept 20 5
HD 250550 V1307 Ori 06:01:60.0 +16:30:57 PHOENIX 2013 Feb 27 700
HD 259431 MWC 147 06:33:05.2 +10:19:20 FIRE 2011 Mar 10 1
HD 275877 XY Per 03:49:36.3 +38:58:56 PHOENIX 2013 Nov 8 600
HD 278937 IP Per 03:40:47.0 +32:31:54 PHOENIX 2013 Feb 28 800
HD 287823 05:24:08.0 +02:27:47 PHOENIX 2013 Mar 02 600
HD 287841 V346 Ori 05:24:42.8 +01:43:48 PHOENIX 2013 Nov 08 600
HD 290409 05:27:05.5 +00:25:08 X-shooter 2010 Jan 02 15
HD 290500 05:29:48.0 –00:23:43 X-shooter 2009 Dec 17 75
HD 290770 05:37:02.4 –01:37:21 X-shooter 2009 Dec 26 7
HD 293782 UX Ori 05:04:29.9 –03:47:14 FIRE 2016 Sept 20 2
MWC 1080 23:17:25.6 +60:50:43 PHOENIX 2013 Nov 10 600
VV Ser 18:28:47.9 +00:08:39 FIRE 2016 Sept 20 3
LkHa 215 06:32:41.8 +10:09:34 PHOENIX 2013 Mar 1 800
Table 1: He i 10830 Å Spectra

2.2 Other spectra

In order to obtain He i 10830 Å profiles for as many sources reported by Alecian et al. (2013b) as possible, we also include profiles from Cauley & Johns-Krull (2014) and Fairlamb et al. (2015). We list the observational details for these 32 spectra in Table 1.

Near-IR spectra from Fairlamb et al. (2015) were obtained in 2009-2010 with the X-Shooter spectrograph on the VLT (Vernet et al., 2011). We obtained Phase 3 pipeline-reduced spectra from the ESO data archive. Spectral resolution for the 0.4″-0.5″ slit widths used by Fairlamb et al. (2015) is . A more complete description of the data is given by Fairlamb et al. (2015).

We also include He i 10830 Å profiles from Cauley & Johns-Krull (2014). These targets are primarily in the northern hemisphere and complement the sample of sources in the southern hemisphere obtained with FIRE and X-Shooter. Specific observation parameters for these data are described by Cauley & Johns-Krull (2014). Briefly, data were obtained with GNIRS (Elias et al., 2006) on Gemini () and PHOENIX (Hinkle et al., 1998) on the Mayall 4 m and KPNO 2.1 m (). Gratings and order-blocking filters were used to isolate emission near the He i 10830 Å line. Data were reduced using custom IDL routines.

HD number Target SpT LSD [G] Regression [G] line profile
BD-06 1253 V380 Ori B9 1.99 460 70 PC
HD 36982 LP Ori B1.5 3.22 22050 O
HD 190073 V1295 Aql A1 1.92 11113 9181 PC
HD 101412 B9/A0 1.36 78555 46527 IPC
HD 35929 F1 2.12 O
see discussion of reduction methods in Section 2.3
HD 35929 is a marginal detection and a -Scuti pulsating variable (Marconi et al., 2000)
whose polarization signature may not be magnetic in origin.
B-field discovery papers: Wade et al. (2005), Petit et al. (2008), Catala et al. (2007),
Wade et al. (2007), Alecian et al. (2013b), Hubrig et al. (2011), Wade et al. (2016),
Hubrig et al. (2013)
Table 2: Herbig Ae/Be stars with a detected B-field

2.3 B field measurements and source properties

We list the stellar parameters – M, L, R, , – for each source in Table 3 using data from Alecian et al. (2013b) and Fairlamb et al. (2015). Longitudinal field strengths for the sources with a detection are listed in Table 2. Deriving a field topology from the longitudinal field strength requires time-series measurements and models of the surface magnetic field (e.g., Donati et al., 2007). We do not include sources where the magnetic field is likely associated with a low-mass companion, i.e., HD 72106 and HD 200775 (Folsom et al., 2008; Alecian et al., 2008, respectively). A few of these objects also have magnetic field detections reported by Hubrig et al. (2009, 2011); we include these in Table 2 where available. Both groups use spectropolarimetric data to determine the longitudinal field strength. Alecian et al. measure magnetic fields using the Least-Square Deconvolution (LSD) technique with higher-resolution () data from ESPaDOnS (on CFHT) and/or Narval (on Télescope Bernard Lyot). Hubrig et al. use the regression method of Bagnulo et al. (2002) on lower-resolution () FORS (on VLT) data. The two groups often come to different conclusions about the magnetic field strength (e.g., Hubrig et al., 2011). Alecian et al. (2013b) apply a more stringent detection criterion by computing a false alarm probability that the observed Stokes V profile could be produced in the absence of a magnetic field. In contrast, Hubrig et al. assume that photon counting statistics dominate their uncertainties. However, Bagnulo et al. (2012) demonstrated that photon noise is not the only source of uncertainty even in high signal-to-noise data. Re-reducing the FORS data with a more complete treatment of the uncertainties, Bagnulo et al. (2012) do not confirm many of the field detections claimed in the literature. For our analysis, we prefer the more conservative criteria employed by Alecian et al. (2013b), although we also include field strengths for those stars measured by Hubrig et al. in Table 2.

Spectral types, mass accretion rates, , and disk inclinations, , taken from the literature are listed in Table 3. Spectral types are from Table 2 in Cauley & Johns-Krull (2014) or Simbad. Mass accretion rates are taken from Fairlamb et al. (2015) where available, or from the literature via Table 2 in Cauley & Johns-Krull (2014). Most disk inclination angle estimates are taken from Lazareff et al. (2017) who use PIONIER/VLTI to obtain H-band ( µm) long-baseline interferometric observations of HAeBe disks. Near-IR visibilities are fit with either an ellipsoidal or ring-like profile. For the values listed in Table 3, we report inclinations derived from the ring-like (labeled “rl” in Table 3) model where available, but include those derived from the ellipsoidal (“el” in Table 3) model for sources where no value for the ring model is reported. Following the inclination angle conventions in that paper, a source viewed edge-on will have while a source viewed pole-on will have .

3 Line Profiles

Figure 1: He i 10830 Å line profiles of sources targets for a magnetic field measurement by Alecian et al. (2013b).
Figure 2: Same as Figure 1, but for HAeBes with a magnetic field detection (see Table 2).

The sensitivity of various spectral lines to the kinematics of gas near young stars depends on the optical depth and thermalization of the line (see, e.g., Calvet & Hartmann, 1992; Hartmann et al., 1994; Muzerolle et al., 2001). HAeBes are hotter and brighter than T Tauri stars and the Balmer lines tend to be more optically thick, making it difficult to trace the kinematics of the gas from the line profile (Muzerolle et al., 2004). This is especially true if targeting redshifted absorption, as we do here, as an unambiguous indicator of infalling gas.

Instead of the Balmer lines, we examine the profile of He i 10830 Å to look for signatures of accretion (i.e. redshifted absorption) in HAeBes targeted for a measurement of the magnetic field by Alecian et al. (2013b). He i 10830 Å has not yet been calibrated as a measure of the accretion rate. However, the metastable lower level makes the transition particularly prone to absorption, and therefore a sensitive tracer of the kinematics of gas near the star. The line is primarily populated through recombinations following excitation by high-energy photons near the star (e.g., Dupree et al., 2005; Kurosawa et al., 2011). The lower level lies 20 eV above ground, restricting the region of its formation – and therefore the gas kinematics it traces — closer to the star.

Combining our new FIRE spectra with those from the literature, we have He i 10830 Å profiles for 63/70 (90%) of the HAeBes targeted for a magnetic field measurement by Alecian et al. (2013b). We also include a previous confirmation from the literature (HD 101412) for a total 64 HAeBes with spectra and 5/64 (8%) with a detected magnetic field. We present the He i 10830 Å profiles for all 64 sources in Figures 1-2. Profile shape classifications for each source are listed in Table 3. Since we are primarily interested in tracers of accretion, we follow a simplified version of the line classification scheme used by Cauley & Johns-Krull (2014), sorting profiles into one of three categories: (1) blueshifted absorption (P Cygni profiles – PC; 19 or 30%), (2) redshifted absorption (inverse P Cygni profiles – IPC; 15 or 23%), or (3) other profile type (O; 30 or 47%).

Altogether, 53% of the line profiles show redshifted or blueshifted absorption, either directly or indirectly (see Section 4) indicative of disk accretion. A similar fraction of T Tauri stars display PC or IPC profiles, with half (18/38; 47%) of the sources presented by Edwards et al. (2006) showing one of the two profiles. However, considering all profiles that show redshifted absorption (even if blueshifted absorption is also present), Edwards et al. (2006) find that half the sample shows redshifted absorption in He i 10830 Å. A more complete analysis of redshifted absorption in HAeBes requires detailed modeling of the excitation and radiative transfer driving He i 10830 Å line formation (see, e.g., Fischer et al., 2008; Kurosawa et al., 2011).

Among the 5 HAeBes with a reported magnetic field detection (see Table 2), roughly half (2/5) show blueshifted absorption (PC), one has redshifted absorption (IPC), and the remaining two have other profiles shapes. Differences in the line profiles between the magnetic HAeBes and the non-magnetic HAeBes are difficult to quantify, given the small number of HAeBes with a detected magnetic field.

Magnetic fields are notoriously difficult to measure (Shorlin et al., 2002). We therefore conduct a parameter study to explore whether line profile morphology correlates with any of the stellar parameters, indirectly indicating the influence of the magnetic field. Figure 3 shows the distribution of redshifted absorption, blueshifted absorption, and other line profiles as a function of the stellar parameters listed in Table 3.

Magnetic braking has been invoked to explain the slow rotation of some stars. Indeed, the magnetic HAeBes appear to have much slower rotation rates than the non-magnetic sample (Alecian et al., 2013a). At the same time, magnetic fields are more difficult to measure in stars with faster rotation rates, as typically observed in HAeBes (Shorlin et al., 2002). Faster rotation rates may obscure Zeeman broadening due to more modest magnetic fields, leading to the preferential detection in sources where fields are anomalously large or the rotation rates are particularly slow. Nevertheless, if slower rotators are more likely to be associated with redshifted absorption, this may be interpreted as indirect evidence for a magnetic field. We show a histogram of different line profiles versus in Figure 3. Neither the distribution of redshifted absorption nor blueshifted absorption are statistically different from the overall distribution of .

Figure 3: Histograms showing He i 10830 Å line profiles observed as a function of stellar parameters (see Table 3). The dotted line shows all sources observed by Alecian et al. (2013b), the solid line indicates the subset with He i 10830 Å data. The red histogram shows those sources with redshifted absorption; blue dashed histogram shows the sources with blueshifted absorption; gray histogram shows the distribution of all other line profiles classifications.

The true rotation rate is likely larger than the measured since most sources will not be observed edge-on. Uncertainty in the viewing angle will redistribute fast and slow rotators in a plot of . Some of this uncertainty can be mitigated where inclination constraints exist. Disk inclinations for 25 of the targets considered here were recently derived from H-band interferometry by Lazareff et al. (2017). We list these in Table 3, together with an inclination estimate for HD 101412 from the literature. We assume that the disk axis and the rotation axis of the star are aligned. This allows us to compute the rotation period for stars with disk inclination estimate using the radii listed in Table 3. Rotation periods for these 27 stars are listed in Table 3 and the resulting distribution of line profiles as a function of period is shown in Figure 3.

We also consider line profiles as a function of the source inclination. In the magnetospheric accretion paradigm, redshifted absorption will only be observable for certain viewing geometries, i.e. if the accretion column intersects the line of sight (see, e.g., Figure 1 in Hartmann et al., 2016). For sources that are viewed nearly pole-on, the line of sight is most likely to intersect the outflowing gas, and thus the line may be more likely to show blueshifted absorption. We plot line profiles as a function of the inclination angle in Figure 4. Among the small number of sources with both an inclination estimate and red/blueshifted absorption, there is no clear trend with source inclination.

For all stellar parameters that we consider, neither the redshifted nor the blueshifted absorption profiles are distributed in a manner that is statistically distinguishable from the overall sample. Probabilities returned from a two-sided Kolmogorov-Smirnov test for each parameter support the null hypothesis (that the two distributions have the same parent population). With the caveat that sample sizes are small (especially for parameters like ), none of the line profiles differ from the overall distribution with likelihood %.

Figure 4: Top: Histograms showing line profiles observed in sources with an estimated disk inclination angle. As in Figure 3, the dotted line shows all sources with an inclination estimate, the solid line indicates the subset with a spectrum, the red histogram shows those sources with redshifted absorption, the blue dashed histogram shows blueshifted absorption profiles, and the gray histogram shows all other line profiles classifications. Bottom: Histogram showing the source inclination angles of magnetic HAeBes (gray) compared to the overall sample with estimated inclination angles. The source with is HD 101412 (see Section 4).

Lastly, we compare the inclination angles of the sample as a whole with the inclination angles of the magnetic HAeBes. The resulting histogram is shown in Figure 4. Few of the magnetic HAeBes also have a measured inclination. However, those that do tend to be observed nearly pole-on (). The only magnetic HAeBe observed with a more edge-on orientation is HD 101412, with an estimated from Fedele et al. (2008). However, this result derives from longer wavelength (mid-IR) observations obtained with significantly fewer baselines than used by Lazareff et al. (2017). Marginally resolved observations of colder dust may sample different structures in the disk (i.e. flaring) at larger radii, biasing the inferred inclination angle.

4 Discussion and Conclusions

We present He i 10830 Å profiles of 64 HAeBes targeted for a magnetic field measurement. More than half of the sources in our sample display either redshifted (15/64; 23%) or blueshifted (19/64; 30%) absorption, directly or indirectly indicating accretion. Cauley & Johns-Krull (2014) find a remarkably similar proportion of IPC (redshifted) and PC (blueshifted) profiles (20% and 30%, respectively) in a similarly sized sample of HAeBes. While there is some overlap between the Cauley & Johns-Krull (2014) sample and the one presented here (see Table 1), the two studies differ in their target selection criterion. Cauley & Johns-Krull (2014) required only that a previous survey identified an object as a HAeBe whereas we present HAeBes that Alecian et al. (2013b) targeted for a magnetic field measurement. By comparing the profiles of sources with and without a detected magnetic field, we examine the role of strong, ordered magnetic fields in accretion in HAeBes.

Of the 64 HAeBes presented here, 5 (8%) have a detected magnetic field (see Table 2). Most (3/5; 60%) of the magnetic HAeBes show redshifted (1/5; 20%) or blueshifted (2/5; 40%) absorption in He i 10830 Å. Redshifted absorption can only be created by infall and is the most direct indication that these HAeBes (magnetic or not) are accreting.

Blueshifted absorption tracing outflowing gas may be interpreted as indirect evidence for accretion. Accretion energy powers outflows and indeed sources with an enhanced accretion rate (i.e. FU Ori-like outbursts) also have an elevated mass-loss rate (e.g., Croswell et al., 1987). However, the existence of an outflow cannot be taken as evidence of magnetospheric accretion. Jet rotation studies (e.g., Bacciotti et al., 2002; Coffey et al., 2004; Ferreira et al., 2006) suggest that jets may launch from a range of disk radii that lie outside the influence of the stellar magnetic field (although see Coffey et al., 2015). Cauley & Johns-Krull (2014) argue that a lack of narrow, blueshifted absorption in He i 10830 Å points to the absence of inner disk winds. Instead, the broad blueshifted absorption profiles are more consistent with stellar winds (e.g., Catala et al., 2007, Aarnio et al., submitted).

The He i 10830 Å line profiles of the 64 HAeBes presented here are not significantly different between sources with and without a detected magnetic field. Roughly the same fraction of non-magnetic HAeBes show blueshifted absorption (17/59; 29%) or redshifted absorption (14/59; 24%) as the few sources with a detected field (40% or 20%, respectively). A more robust comparison is difficult given the small number of sources with a field detection.

Despite the similarity of their line profiles, Alecian et al. (2013a) report markedly different rotation rates between the magnetic and non-magnetic HAeBes. Magnetic braking (e.g., Stȩpień, 2000) may lead to slower rotation rates even in sources with fields too weak and/or too disordered to be detected. Coupling between the star and disk may be more difficult in this case; nevertheless, if the slowest rotating non-magnetic stars all show redshifted absorption, this might hint at an underlying magnetospheric accretion process. To test this hypothesis, we compare the He i 10830 Å profiles to the projected rotational velocity measured at the stellar surface, (from Alecian et al., 2013b, see Figure 3). The distribution of redshifted and blueshifted profiles are statistically indistinguishable from the overall distribution of for the full sample.

Projection effects may intermix faster and slower rotators in a plot of , obscuring any underlying trend. Assuming that the disk and rotation axes are aligned, we correct for projection and compute the rotation period for 26/64 HAeBes. Again, the distributions of redshifted and blueshifted profiles do not indicate any trend with rotation period (see Figure 3). Altogether, this argues against accretion onto HAeBes being mediated by a dipolar magnetic field.

None of the line profiles appear to correlate with direct or indirect indications of a magnetic field. Among the comparisons between stellar parameters and orientation effects that we explore in this paper, only the relationship between a magnetic field detection and the disk inclination of the source hints at a correlation. Of the three magnetic HAeBes with an estimated disk inclination (see Table 2), all but one have . The only source seen nearly edge-on is HD 101412, although we consider this estimate less reliable than those obtained by Lazareff et al. (2017) (see Section 3).

If inclination affects field detectability, this suggests that magnetic sensitivity is viewing angle dependent. Zeeman splitting of magnetic-sensitive lines will be harder to separate from Doppler broadening in stars with faster rotation rates. Stars seen nearly pole-on will have slower projected surface rotation velocities, and thus Doppler and Zeeman broadening may be more readily disentangled. Weak fields may be more difficult to detect among the faster rotating HAeBes.

Gregory et al. (2012) have argued that the magnetic field evolves during the pre-main-sequence evolution of stars that develop a radiative envelope. All of the HAeBes in the Alecian et al. (2013b) sample lie in a portion of the H-R diagram where stars should have substantial radiative envelopes (see Figure 4 in Gregory et al., 2012). Whereas convective interiors can support the strong dipolar fields observed on some T Tauris (e.g., Johns-Krull et al., 1999a; Johns-Krull, 2007; Johns-Krull et al., 2013), once the radiative envelope has developed, the dipole component gets weaker and the complexity of the field increases.

Even strongly magnetic low-mass stars are not well-described by a pure dipole (e.g., Johns-Krull et al., 1999a). Nevertheless, the strength of non-dipolar components decreases rapidly with distance from the stars (Valenti & Johns-Krull, 2004). While the standard magnetospheric accretion paradigm for low-mass stars assumes strong dipolar fields, several authors have argued that magnetic-mediated accretion can proceed even with realistically complex fields (e.g., Gregory et al., 2006, 2008; Mohanty & Shu, 2008; Adams & Gregory, 2012; Johnstone et al., 2014). Weak, complex fields may exist on HAeBes, below the detection limits of Alecian et al. (2013b). The absence of strong dipolar fields does not preclude the possibility that accretion may proceed along the higher-order field lines that direct material closer to the stellar equator.

Comparing the He i 10830 Å emission of HAeBes with and without a detected magnetic field suggests that line profile morphologies are insensitive to the magnetic field. However, few sources in our sample have a detected magnetic field; the majority have upper limits on the field strength. Both Muzerolle et al. (2004) and Cauley & Johns-Krull (2014) argue that magnetically-mediated accretion may be possible through a more compact accretion geometry where higher-order field components create a smaller magnetosphere. Additional work is needed to model the excitation and radiative transfer of the He i 10830 Å line. Detailed theoretical models are essential to determine if the magnetospheric accretion paradigm can be modified to allow accretion through weaker, more topologically complex fields.

MR would like to thank Lee Hartmann and Chris Miller for useful discussions. JDM and AA were supported by NSF AST. 1311698. SK acknowledges support from an ERC Starting Grant (Grant Agreement No. 639889) and a STFC Rutherford fellowship (ST/J004030/1). This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. This is research has made use of the services of the ESO Science Archive Facility. Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere under ESO programme(s) 088.C-0218(A). Based on data obtained from the ESO Science Archive Facility under request numbers 231792, 231794, 231804, 231805, 231808, 231809, 231810, 255461, 231795, 231796, 231797, 231799, and 231802. \facilitiesMagellan(FIRE), VLT(X-Shooter, Gemini(GNIRS), KPNO(PHOENIX)


  • Ababakr et al. (2017) Ababakr, K., Oudmaijer, R., & Vink, J. 2017, ArXiv e-prints, arXiv:1707.08408
  • Adams & Gregory (2012) Adams, F. C., & Gregory, S. G. 2012, ApJ, 744, 55
  • Alecian et al. (2013a) Alecian, E., Wade, G. A., Catala, C., et al. 2013a, MNRAS, 429, 1027
  • Alecian et al. (2008) Alecian, E., Catala, C., Wade, G. A., et al. 2008, MNRAS, 385, 391
  • Alecian et al. (2013b) Alecian, E., Wade, G. A., Catala, C., et al. 2013b, MNRAS, 429, 1001
  • Bacciotti et al. (2002) Bacciotti, F., Ray, T. P., Mundt, R., Eislöffel, J., & Solf, J. 2002, ApJ, 576, 222
  • Bagnulo et al. (2015) Bagnulo, S., Fossati, L., Landstreet, J. D., & Izzo, C. 2015, A&A, 583, A115
  • Bagnulo et al. (2012) Bagnulo, S., Landstreet, J. D., Fossati, L., & Kochukhov, O. 2012, A&A, 538, A129
  • Bagnulo et al. (2002) Bagnulo, S., Szeifert, T., Wade, G. A., Landstreet, J. D., & Mathys, G. 2002, A&A, 389, 191
  • Basri & Bertout (1989) Basri, G., & Bertout, C. 1989, ApJ, 341, 340
  • Bertout et al. (1988) Bertout, C., Basri, G., & Bouvier, J. 1988, ApJ, 330, 350
  • Bouvier et al. (1999) Bouvier, J., Chelli, A., Allain, S., et al. 1999, A&A, 349, 619
  • Calvet & Gullbring (1998) Calvet, N., & Gullbring, E. 1998, ApJ, 509, 802
  • Calvet & Hartmann (1992) Calvet, N., & Hartmann, L. 1992, ApJ, 386, 239
  • Calvet et al. (2004) Calvet, N., Muzerolle, J., Briceño, C., et al. 2004, AJ, 128, 1294
  • Catala et al. (2007) Catala, C., Alecian, E., Donati, J.-F., et al. 2007, A&A, 462, 293
  • Cauley & Johns-Krull (2014) Cauley, P. W., & Johns-Krull, C. M. 2014, ApJ, 797, 112
  • Cauley & Johns-Krull (2015) —. 2015, ApJ, 810, 5
  • Coffey et al. (2004) Coffey, D., Bacciotti, F., Woitas, J., Ray, T. P., & Eislöffel, J. 2004, ApJ, 604, 758
  • Coffey et al. (2015) Coffey, D., Dougados, C., Cabrit, S., Pety, J., & Bacciotti, F. 2015, ApJ, 804, 2
  • Cowley et al. (2010) Cowley, C. R., Hubrig, S., González, J. F., & Savanov, I. 2010, A&A, 523, A65
  • Croswell et al. (1987) Croswell, K., Hartmann, L., & Avrett, E. H. 1987, ApJ, 312, 227
  • Donati et al. (2007) Donati, J.-F., Jardine, M. M., Gregory, S. G., et al. 2007, MNRAS, 380, 1297
  • Donehew & Brittain (2011) Donehew, B., & Brittain, S. 2011, AJ, 141, 46
  • Dupree et al. (2005) Dupree, A. K., Brickhouse, N. S., Smith, G. H., & Strader, J. 2005, ApJ, 625, L131
  • Edwards et al. (2006) Edwards, S., Fischer, W., Hillenbrand, L., & Kwan, J. 2006, ApJ, 646, 319
  • Edwards et al. (2003) Edwards, S., Fischer, W., Kwan, J., Hillenbrand, L., & Dupree, A. K. 2003, ApJ, 599, L41
  • Edwards et al. (1994) Edwards, S., Hartigan, P., Ghandour, L., & Andrulis, C. 1994, AJ, 108, doi:10.1086/117134
  • Elias et al. (2006) Elias, J. H., Joyce, R. R., Liang, M., et al. 2006, in Proc. SPIE, Vol. 6269, Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, 62694C
  • Fairlamb et al. (2015) Fairlamb, J. R., Oudmaijer, R. D., Mendigutía, I., Ilee, J. D., & van den Ancker, M. E. 2015, MNRAS, 453, 976
  • Fairlamb et al. (2017) Fairlamb, J. R., Oudmaijer, R. D., Mendigutia, I., Ilee, J. D., & van den Ancker, M. E. 2017, MNRAS, 464, 4721
  • Fedele et al. (2008) Fedele, D., van den Ancker, M. E., Acke, B., et al. 2008, A&A, 491, 809
  • Ferreira et al. (2006) Ferreira, J., Dougados, C., & Cabrit, S. 2006, A&A, 453, 785
  • Finkenzeller & Mundt (1984) Finkenzeller, U., & Mundt, R. 1984, A&AS, 55, 109
  • Fischer et al. (2008) Fischer, W., Kwan, J., Edwards, S., & Hillenbrand, L. 2008, ApJ, 687, 1117
  • Folsom et al. (2008) Folsom, C. P., Wade, G. A., Kochukhov, O., et al. 2008, MNRAS, 391, 901
  • Garcia Lopez et al. (2006) Garcia Lopez, R., Natta, A., Testi, L., & Habart, E. 2006, A&A, 459, 837
  • Gregory et al. (2012) Gregory, S. G., Donati, J.-F., Morin, J., et al. 2012, ApJ, 755, 97
  • Gregory et al. (2006) Gregory, S. G., Jardine, M., Simpson, I., & Donati, J.-F. 2006, MNRAS, 371, 999
  • Gregory et al. (2008) Gregory, S. G., Matt, S. P., Donati, J.-F., & Jardine, M. 2008, MNRAS, 389, 1839
  • Gullbring et al. (2000) Gullbring, E., Calvet, N., Muzerolle, J., & Hartmann, L. 2000, ApJ, 544, 927
  • Hartmann et al. (2016) Hartmann, L., Herczeg, G., & Calvet, N. 2016, ARA&A, 54, 135
  • Hartmann et al. (1994) Hartmann, L., Hewett, R., & Calvet, N. 1994, ApJ, 426, 669
  • Herbig (1960) Herbig, G. H. 1960, ApJS, 4, 337
  • Hillenbrand et al. (1992) Hillenbrand, L. A., Strom, S. E., Vrba, F. J., & Keene, J. 1992, ApJ, 397, 613
  • Hinkle et al. (1998) Hinkle, K. H., Cuberly, R. W., Gaughan, N. A., et al. 1998, in Proc. SPIE, Vol. 3354, Infrared Astronomical Instrumentation, ed. A. M. Fowler, 810–821
  • Hubrig et al. (2013) Hubrig, S., Ilyin, I., Schöller, M., & Lo Curto, G. 2013, Astronomische Nachrichten, 334, 1093
  • Hubrig et al. (2009) Hubrig, S., Stelzer, B., Schöller, M., et al. 2009, A&A, 502, 283
  • Hubrig et al. (2011) Hubrig, S., Schöller, M., Ilyin, I., et al. 2011, A&A, 536, A45
  • Hussain et al. (2009) Hussain, G. A. J., Collier Cameron, A., Jardine, M. M., et al. 2009, MNRAS, 398, 189
  • Johns-Krull (2007) Johns-Krull, C. M. 2007, ApJ, 664, 975
  • Johns-Krull et al. (1999a) Johns-Krull, C. M., Valenti, J. A., Hatzes, A. P., & Kanaan, A. 1999a, ApJ, 510, L41
  • Johns-Krull et al. (1999b) Johns-Krull, C. M., Valenti, J. A., & Koresko, C. 1999b, ApJ, 516, 900
  • Johns-Krull et al. (2013) Johns-Krull, C. M., Chen, W., Valenti, J. A., et al. 2013, ApJ, 765, 11
  • Johnstone et al. (2014) Johnstone, C. P., Jardine, M., Gregory, S. G., Donati, J.-F., & Hussain, G. 2014, MNRAS, 437, 3202
  • Kurosawa et al. (2006) Kurosawa, R., Harries, T. J., & Symington, N. H. 2006, MNRAS, 370, 580
  • Kurosawa et al. (2011) Kurosawa, R., Romanova, M. M., & Harries, T. J. 2011, MNRAS, 416, 2623
  • Kwan et al. (2007) Kwan, J., Edwards, S., & Fischer, W. 2007, ApJ, 657, 897
  • Kwan & Fischer (2011) Kwan, J., & Fischer, W. 2011, MNRAS, 411, 2383
  • Lazareff et al. (2017) Lazareff, B., Berger, J.-P., Kluska, J., et al. 2017, A&A, 599, A85
  • Marconi et al. (2000) Marconi, M., Ripepi, V., Alcalá, J. M., et al. 2000, A&A, 355, L35
  • Mendigutía et al. (2011) Mendigutía, I., Calvet, N., Montesinos, B., et al. 2011, A&A, 535, A99
  • Mendigutía et al. (2012) Mendigutía, I., Mora, A., Montesinos, B., et al. 2012, A&A, 543, A59
  • Mendigutía et al. (2015) Mendigutía, I., Oudmaijer, R. D., Rigliaco, E., et al. 2015, MNRAS, 452, 2837
  • Mohanty & Shu (2008) Mohanty, S., & Shu, F. H. 2008, ApJ, 687, 1323
  • Mottram et al. (2007) Mottram, J. C., Vink, J. S., Oudmaijer, R. D., & Patel, M. 2007, MNRAS, 377, 1363
  • Muzerolle et al. (1998a) Muzerolle, J., Calvet, N., & Hartmann, L. 1998a, ApJ, 492, 743
  • Muzerolle et al. (2001) —. 2001, ApJ, 550, 944
  • Muzerolle et al. (2004) Muzerolle, J., D’Alessio, P., Calvet, N., & Hartmann, L. 2004, ApJ, 617, 406
  • Muzerolle et al. (1998b) Muzerolle, J., Hartmann, L., & Calvet, N. 1998b, AJ, 116, 455
  • Petit et al. (2008) Petit, V., Wade, G. A., Drissen, L., Montmerle, T., & Alecian, E. 2008, MNRAS, 387, L23
  • Popham et al. (1993) Popham, R., Narayan, R., Hartmann, L., & Kenyon, S. 1993, ApJ, 415, L127
  • Shorlin et al. (2002) Shorlin, S. L. S., Wade, G. A., Donati, J.-F., et al. 2002, A&A, 392, 637
  • Simcoe et al. (2013) Simcoe, R. A., Burgasser, A. J., Schechter, P. L., et al. 2013, PASP, 125, 270
  • Stȩpień (2000) Stȩpień, K. 2000, A&A, 353, 227
  • Symington et al. (2005) Symington, N. H., Harries, T. J., Kurosawa, R., & Naylor, T. 2005, MNRAS, 358, 977
  • Takami et al. (2002) Takami, M., Chrysostomou, A., Bailey, J., et al. 2002, ApJ, 568, L53
  • Valenti & Johns-Krull (2004) Valenti, J. A., & Johns-Krull, C. M. 2004, Ap&SS, 292, 619
  • Vernet et al. (2011) Vernet, J., Dekker, H., D’Odorico, S., et al. 2011, A&A, 536, A105
  • Vink et al. (2002) Vink, J. S., Drew, J. E., Harries, T. J., & Oudmaijer, R. D. 2002, MNRAS, 337, 356
  • Vink et al. (2005) Vink, J. S., Drew, J. E., Harries, T. J., Oudmaijer, R. D., & Unruh, Y. 2005, MNRAS, 359, 1049
  • Wade et al. (2007) Wade, G. A., Bagnulo, S., Drouin, D., Landstreet, J. D., & Monin, D. 2007, MNRAS, 376, 1145
  • Wade et al. (2005) Wade, G. A., Drouin, D., Bagnulo, S., et al. 2005, A&A, 442, L31
  • Wade et al. (2016) Wade, G. A., Neiner, C., Alecian, E., et al. 2016, MNRAS, 456, 2
  • Walker (1972) Walker, M. F. 1972, ApJ, 175, 89
  • Waters & Waelkens (1998) Waters, L. B. F. M., & Waelkens, C. 1998, ARA&A, 36, 233
  • Wenger et al. (2000) Wenger, M., Ochsenbein, F., Egret, D., et al. 2000, A&AS, 143, 9
HD/BD other spectral log(L) M R log() cos(i) v He i
number name type [L] [M] [R] [M yr] (km s) (km s) (km s) Profile
BD-06 1259 BF Ori A2 1.75 2.58 3.26 IPC
BD-05 1253 V380 Ori B9 1.99 2.87 3.00 [27.3,28.2] PC
BD-05 1329 T Ori A3 1.97 3.13 4.47 O
BD-05 1324 NV Ori F6 1.32 2.28 3.77 IPC
BD41 3731 B5 3.03 5.50 3.8 O
BD46 3471 V1578 Cyg A1 2.84 5.9 9.7 O
BD61 154 V594 Cyg B8 1.95 3.41 2.42 PC
HD 17081  Cet B8 2.750 4.65 4.84 30.6 [11.0,12.7] O
HD 31293 AB Aur A0 1.76 2.50 2.62 -6.85 0.91 279.8 PC
HD 31648 MWC 480 A4 1.18 1.93 1.93 155.8 PC
HD 34282 A3 1.13 1.59 1.66 115.1 IPC
HD 35187 A2 1.15 1.93 1.58 O
HD 35929 F1 2.12 4.13 8.1 117.3 O
HD 36112 MWC 758 A5 1.81 2.90 4.4 72.0 PC
HD 36910 CQ Tau F2 1.69 2.93 5.1 IPC
HD 36917 V372 Ori B9 2.39 3.98 5.2 127.1 O
HD 36982 LP Ori B1.5 3.22 6.70 3.42 O
HD 37258 V586 Ori A1 1.44 2.28 1.94 219.2 IPC
HD 37357 A1 1.72 2.47 2.83 O
HD 37806 MWC 120 B9 2.45 3.94 4.6 181.4 IPC
HD 38120 B9 1.62 2.49 1.91 O
HD 38238 V351 Ori A6 1.79 2.88 4.38 IPC
HD 50083 V742 Mon B4 4.15 12.1 10.0 O
HD 52721 GU CMa B3 3.77 9.1 5.0 O
HD 53367 MWC 166 B1 4.50 16.1 7.1 O
HD 68695 F2 1.80 2.64 3.3 PC
HD 72106 A0 1.34 2.40 1.3 O
HD72106_B 0.96 1.9 1.3
HD 76534 B2 3.75 9.0 7.7 O
HD 98922 B9 2.48 4.0 PC
HD 101412 B9/A0 1.36 2.0 3.0 IPC
HD 114981 V958 Cen B3/5 3.56 7.9 7.0 2.0 -2.0 O
HD 139614 F0 1.10 1.76 2.06 45.7 O
HD 141569 A0 1.49 2.33 1.94 O
HD 142666 V1026 Sco A5 1.44 2.15 2.82 6.1 IPC
HD 144432 A7 1.28 1.95 2.59 190.1 PC
HD 144668 HR 5999 A7 1.56 2.31 3.0 251.1 O
HD 145718 V718 Sco A4 1.29 1.93 2.25 128.5 O
HD 150193 MWC 863 A1 1.79 2.56 2.89 205.0 PC
HD 152404_A AK Sco F5 0.95 1.66 2.4 PC
HD 152404_B AK_Sco_B 0.71 1.43 1.79
HD 163296 A1 1.52 2.23 2.28 173.8 PC
HD 169142 A7 0.88 1.69 1.64 122.0 O
HD 174571 MWC 610 B3 3.58 8.0 4.7 O
HD 176386 B9 1.91 3.02 2.28 O
HD 179218 MWC 614 A0 2.26 3.66 4.8 91.6 O
HD 190073 V1295 Aql A1 1.92 2.85 3.60 [0-8.3] 0-16.3 PC
HD 200775 A MWC 361 A B2 3.95 10.7 10.4 [-23.3,8.2] IPC
HD 200775 B MWC 361 B ? 3.77 9.3 8.3 [-21.1,9.3]
HD 216629 A IL Cep A B4 2.58 [-39,31] O
HD 216629 B IL Cep B ? ? [-87,-30]
HD 244314 V1409 Ori A1 1.45 2.33 2.07 PC
HD 244604 V1410 Ori A4 1.74 2.66 3.69 120.7 PC
HD 245185 V1271 Ori A1 1.40 2.19 1.85 O
HD 249879 A2 2.31 4.0 5.9 -8.00 O
HD 250550 V1307 Mon B8 100.7 PC
HD 259431 MWC 147 B6 3.35 7.1 8.0 417.1 PC
HD 275877 XY Per A2 1.21 1.95 1.65 IPC
HD 278937 IP Per A3 1.21 1.86 2.10 PC
HD 287823 A A0 1.79 2.5 2.6 O
HD 287823 B ? 0.82 1.6 1.8 ?
HD 287841 V346 Ori A7 1.05 1.72 1.96 IPC
HD 290409 A2 1.32 2.04 1.75 O
HD 290500 A2 1.22 1.96 1.68 IPC
HD 290770 B9 1.91 2.86 2.49 PC
HD 293782 UX Ori A1 2.98 6.72 12.1 448.2 IPC
MWC 1080 B1 5.77 17.4 7.3 PC
VV Ser B7 2.51 4.0 3.1 -7.50 155.0 IPC
LkHa 215 A B7 3.08 5.8 5.9 O
LkHa 215 B ? 3.08 5.8 5.9 [12,22]
PC = P-Cygni profile (blueshifted absorption); IPC = inverse P-Cygni profile (redshifted absorption), O = other
spectral type from Simbad (Wenger et al., 2000), Fairlamb et al. (2015), Mendigutía et al. (2011), Cauley & Johns-Krull (2014), Cowley et al. (2010),
Fedele et al. (2008), Hubrig et al. (2009), Garcia Lopez et al. (2006), Donehew & Brittain (2011)
ring model from Lazareff et al. (2017), ellipsoidal model from Lazareff et al. (2017)
Table 3: Herbig Ae/Be stars targeted for magnetic field measurement
Comments 0
Request Comment
You are adding the first comment!
How to quickly get a good reply:
  • Give credit where it’s due by listing out the positive aspects of a paper before getting into which changes should be made.
  • Be specific in your critique, and provide supporting evidence with appropriate references to substantiate general statements.
  • Your comment should inspire ideas to flow and help the author improves the paper.

The better we are at sharing our knowledge with each other, the faster we move forward.
The feedback must be of minimum 40 characters and the title a minimum of 5 characters
Add comment
Loading ...
This is a comment super asjknd jkasnjk adsnkj
The feedback must be of minumum 40 characters
The feedback must be of minumum 40 characters

You are asking your first question!
How to quickly get a good answer:
  • Keep your question short and to the point
  • Check for grammar or spelling errors.
  • Phrase it like a question
Test description