The Astronomical Journal, 139:27–38, 2010 January doi:10.1088/0004-6256/139/1/27C© 2010. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

ANALYSIS OF THE PICO DOS DIAS SURVEY HERBIG Ae/Be CANDIDATES∗

Marılia J. Sartori1, Jane Gregorio-Hetem

2, Claudia V. Rodrigues

3, Annibal Hetem Jr.

4, and Celso Batalha

51 Laboratorio Nacional de Astrofısica/MCT, Rua Estados Unidos 154, 37504-364 Itajuba, MG, Brazil; marilia@lna.br

2 Universidade de Sao Paulo, IAG, Rua do Matao 1226, 05508-900 Sao Paulo, SP, Brazil3 Instituto Nacional de Pesquisas Espaciais/MCT, Av. dos Astronautas 1758, 12227-010 Sao Jose dos Campos, SP, Brazil

4 Fundacao Santo Andre, FAENG, Rua Principe de Gales 821, 09060-650 Santo Andre, SP, Brazil5 Evergreen Valley College, 3095 Yerba Buena Rd., San Jose, CA, 95135, USA

Received 2009 May 28; accepted 2009 October 19; published 2009 November 18

ABSTRACT

A large sample of Herbig Ae/Be (HAeBe) candidates, distributed in different Galactic regions south to declination+30◦, were identified by the Pico dos Dias Survey (a search for young stellar objects based on IRAS colors). Mostof the candidates are nearby or associated with star-forming clouds, but several others are considered isolatedobjects. Aiming to verify the young nature of 93 HAeBe candidates, we searched for additional information thatcould be useful to confirm if they are pre-main-sequence (PMS) stars or evolved objects, which coincidentallyshow similar IRAS colors. By adopting a spectral index that is related to the amount of infrared excess and theshape of the spectral energy distribution, we have classified the sample according to three groups, which areanalyzed on the basis of (1) circumstellar luminosity; (2) spatial distribution; (3) optical polarization; (4) near-infrared colors; (5) stellar parameters (mass, age, effective temperature); and (5) intensity of emission lines. Ouranalysis indicates that only 76% of the studied sample, mainly the group with intermediate to low levels ofcircumstellar emission, can be more confidently considered PMS stars. The nature of the remaining stars, whichare in the other group that contains the highest levels of infrared excess, remains to be confirmed. They sharethe same characteristics of evolved objects, requiring complementary studies in order to correctly classify them.At least seven objects show characteristics typical of post-asymptotic giant branch or proto-planetary nebulae.

Key words: circumstellar matter – stars: fundamental parameters – stars: pre-main sequence

1. INTRODUCTION

After a search for massive pre-main-sequence (PMS) com-panions of T Tauri stars (TTs), by Herbig (1960), the HerbigAe/Be stars (HAeBes) were introduced as a new class of youngstellar objects following three basic criteria: (1) spectral typeearlier than F0, (2) spectral features similar to TTs (Balmerlines), and (3) association with the star-forming region (SFR).The PMS nature of the HAeBes was confirmed by Strom et al.(1972), who studied their position on the H–R Diagram indi-cating ages in the ranges of 0.1–1 Myr. Differences in physicalproperties, evolution, and circumstellar material allow a separa-tion into two sub-classes: HAe, spectral type later than or equalto B6 and HBe, spectral type earlier than or equal to B5 (Nattaet al. 2000).

During a search for PMS stars based on IRAS colors calledthe Pico dos Dias Survey (PDS), several new TTs were revealed,as well as previously unknown HAeBe candidates (Gregorio-Hetem et al. 1992; Torres et al. 1995; Torres 1999). Thesecandidates have Hα emission in at least one spectrum, spectraltype earlier than F5, and luminosity classes III–V. The criteriaused by the PDS group to classify these stars are described inVieira et al. (2003), who listed 108 objects, 18% of them beingpreviously classified as HAeBes (The et al. 1994; Malfait et al.1998).

All efforts focused on confirming the nature of the HAeBecandidates have shown that this class of stars does not haveexclusive observational characteristics. Many of the spectralfeatures, for example, are also present in more evolved objects,as verified by Fernandes & Araujo (2003) in a study of classical

∗ Based on observations made at the Observatorio do Pico dos Dias/LNA(Brazil).

Be stars, for example. Young stellar objects, post-asymptotic gi-ant branch (AGB), and proto-planetary nebula (PPN) candidateshave similar IRAS colors, since stars in different evolutionaryphases show far-IR excess due to circumstellar dust. In theIRAS color–color diagram PPNs may share the same positionoccupied by TTs and H ii regions (Garcıa-Lario et al. 1997),for example. Therefore, this useful tool must be applied withcaution to avoid incorrect classifications, requiring the anal-ysis of other characteristics. We are aware that some of themost embedded objects of the PDS sample may not correspondto PMS stars. Two examples are Hen 3-1475 (PDS 465) andIRAS19343+2926 (PDS 581) that were confirmed to be PPN(Riera et al. 1995; Bowers & Knapp 1989; Rodrigues et al.2003). A subsample of 27 PDS HAeBe candidates showinglarge IR excess was studied by Vieira et al. (2009). They ver-ified that at least six candidates show characteristics similarto those found for PDS 465 and 581, according to the analy-sis of color–color diagrams, spatial distribution, spectral energydistribution (SED) fitting, and spectral characteristics. Theseprobably evolved objects have the largest amounts of IR excessin our sample, which is a characteristic different from those pre-sented by the confirmed HAeBe stars. The classification of thePDS sample according to the SED shapes is just a first step inthe analysis of the actual nature of these HAeBe candidates.

Hillenbrand et al. (1992) proposed a classification of47 HAeBes in three groups based on the slope of the IR contin-uum. Their classification is comparable to the groups suggestedby Meeus et al. (2001) for a sample of 14 isolated HAeBe stars.Inspired by these examples of HAeBe groups separation, wesuggested in a previous work (Sartori et al. 2003) a similarclassification for almost a hundred of HAeBe candidates fromPDS. Although all these candidates have been selected on basisof their similar IRAS colors (F25/F12 and F60/F25 flux ratios),

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28 SARTORI ET AL. Vol. 139

their SED shapes are quite different, leading us to separate theminto three groups based on the spectral index calculated in theregion between optical and mid-IR. We adopted a simple diskmodel to estimate the circumstellar luminosity (Sc), defined bythe contribution of circumstellar emission to the total emittedflux, which is evaluated by integrating the synthetic SED in theIR band.

In the present work we analyze the SED classification basedon the optical–mid-IR spectral index aiming to confirm the PMSnature and to verify the validity of group separation to investigatethe evolutionary status of the HAeBe candidates from PDS.We also compare the IR excess to stellar characteristics ofthis sample, considering their SED classification, and using lowresolution spectra, polarimetric data, and near-IR data. In thiscontext, our study of the PDS HAeBE candidates is a follow-up of the work presented by Vieira et al. (2003). Compared toprevious studies on HAeBe classification, the present work hasthe advantage of using a larger sample of objects, that has almosta hundred of candidates.

In Section 2, we present the sample selection and dataacquisition (Section 2.1.1), the Sc estimation by SED fitting(Section 2.1.2), and the HAeBe groups classification based onthe optical–mid-IR spectral index (Section 2.2). In Section 3, theassociation with SFRs is investigated. Section 4 describes theanalysis of the adopted SED classification and the significance ofSc in relation to (1) the near-IR spectral index and the color–colordiagram using near- and mid-IR fluxes; (2) stellar parameters:effective temperature, mass, and age, (3) emission line profile,and (4) optical polarization. Finally, the results are discussedand the main conclusions are summarized in Section 5.

2. SPECTRAL ENERGY DISTRIBUTION

2.1. Circumstellar Luminosity

2.1.1. Sample Selection and Observational Data

Our sample was extracted from the PDS list of 108 HAeBecandidates, except for 14 objects that are visual binaries (ap-parent or intrinsic) for which optical and/or far-IR individualphotometries are not available. PDS 141 was not included inour sample due to the lack of optical photometry for this IRobject. Table 2 lists 93 selected stars that are indicated by theirPDS number and the Two Micron All Sky Survey (2MASS)source identification. Several stellar and circumstellar param-eters are also given, most of them determined in the presentwork.

For the construction of the SEDs we used: (1) the UBV(RI)cphotometry obtained during the PDS data acquisition (Gregorio-Hetem et al. 1992; Torres et al. 1995; Torres 1999); (2) near-and far-IR data extracted respectively from 2MASS and IRASPSPC catalogues; and (3) mid-IR photometry from the MSXcatalogue when available (∼47% of the sample), mainly in bandA (8.28 μm).

Optical and near-IR data were corrected for interstellarextinction by using the relation AV = 3.1E(B − V ) given bySavage & Mathis (1979) and the extinction law Aλ/AV fromCardelli et al. (1989). Considering the UV excess presentedby HAeBe stars, the (B − V) excess was estimated fromE(B −V ) = E(V − I )/1.6 (Schultz & Wiemer 1975). Intrinsiccolors and bolometric corrections from Siess et al. (2000)were adopted in the estimation of AV and absolute bolometricmagnitude.

Vieira et al. (2003) suggest the association of some ofthe HAeBe with SFRs. For these stars, we have adopted the

distances of the respective clouds (see details in Section 3).Hipparcos parallax is available for some of the stars leading toa better determination for distance in this case. Temperatureswere estimated from the Teff–spectral type relations given byde Jager & Nieuwenhuijzen (1987) by adopting spectral typesestimated by Torres (1999) and, when that was not available,from the literature.

2.1.2. SED Fitting

Different circumstellar geometries have been suggested toexplain the SED of HAeBes (see the reviews by Perez & Grady1997; Waters & Waelkens 1998; Natta et al. 2000). Althoughspherical envelope models and accretion disk models have bothsuccessfully fitted the SEDs of some HAeBe stars, severalobservations have pointed to the presence of disks showinglow accretion rates. Passive irradiated circumstellar disk modelshave been used to reproduce the SEDs of HAe stars (e.g.,Dullemond et al. 2001; Dominik et al. 2003). As the HAeBesshow many different SED shapes, it is difficult to predict asingle model that could explain the circumstellar structure of allHAeBe stars.

In order to analyze the circumstellar matter distribution ofthe PDS HAeBes, we adopted a simple model assuming ageometrically flat, passive disk surrounded by a tenuous dustenvelope (Gregorio-Hetem & Hetem 2002). The passive diskre-irradiates the energy absorbed from the star, leading to IRexcess even in the absence of any accretion luminosity (Adamset al. 1987; Kenyon & Hartmann 1987; Strom et al. 1993).

Optical depth, disk radius, and envelope radius are theparameters of the SED fitting that provide the contribution fromeach circumstellar component. The inner radius of the diskis constrained by adopting the location of grain sublimationpoint. Although the adopted geometry in this model seemsunreal for most of the objects, the calculated disk luminosity(Ld) can be interpreted in terms of warm dust emission, sincethis circumstellar component is closer to the star and theadopted disk temperature and density laws provide reprocessedemission mainly in near-IR. Furthermore, the spherical envelopecontribution (Le) plays the role of the cold dust emission. Thetemperature and density laws adopted for the thick disk haveslopes that are steeper than those for optically thin dust shell.The circumstellar luminosity compared to the total luminosityis expressed by Sc = (LT −L�)/LT , where LT = Ld +Le +L�.Two examples of SED fitting are shown in Figure 1.

Aiming to compare SED fittings obtained with a differentenvelope geometry, we have also fitted the SED using the flareddisk configuration suggested by Dullemond et al. (2001) fora passively irradiated circumstellar disk with an inner hole.A code based on genetic algorithms, developed by Hetem &Gregorio-Hetem (2007), was used to optimize the parameterestimation. The flared disk configuration provided good SEDfittings mainly for objects showing low- to moderate IR excess,but the same cannot be said about embedded or non-isolatedobjects (increasing SED slope in the IR), whose SEDs werenot well fitted. Only about 50% of our sample could have theirSEDs fitted using the flared disk model. For this part of thesample, it was possible to compare the circumstellar luminosityestimated by both methods showing that compatible values ofSc are achieved, independently of the adopted disk geometry(see Section 4.1). By comparing the flat disk model (Gregorio-Hetem & Hetem 2002) to the flared configuration (Hetem &Gregorio-Hetem 2007) it can be said that roughly the flat diskplays the role of the warm and/or hot flared disk component

No. 1, 2010 ANALYSIS OF THE PDS HERBIG Ae/Be CANDIDATES 29

Figure 1. Examples of HAeBe candidates classified as group 1. The SED fitting distinguishes contributions from individual components: star (dashed line), disk(dotted line), and envelope (dash-dotted line).

and the tenuous thin envelope mimics the cold component ofthe flared model.

The incompleteness of the SED fittings obtained with theflared disk leads us to use the results derived from the flatdisk model. The geometry of the disk/envelope components isnot discussed here, since they may not have physical meaning.Due to this uncertainty on derived parameters, the SED fittingwas used only for the estimation of circumstellar luminosity(Sc), required in the statistical analysis of our large sample ofHAeBe stars. Table 2 gives the estimated Sc. Regardless of diskparameters, the SED fitting calculations provide good estimationof both, warm and cold, dust emission components.

2.2. HAeBe Groups

The SED classification adopted here is based on the spectralindex (SED slope) measured in the visible to mid-IR range—from the V band (0.55 μm) to the IRAS 12 μm band, given byβ1 = 0.75 log (F12/FV ) − 1 (Torres 1999). This index is relatedto the amount of IR excess that was used to classify the stars ofour sample in three groups.

Objects in group 1 (β1 � 0) have far-IR emission greater thanthe optical emission, while in the other extreme of this sequenceare the objects in group 3 (β1 < −1) with far-IR emission muchsmaller than the optical emission. Group 2 (−1 � β1 < 0)shows far-IR emission corresponding to intermediate values ofIR excess.

Table 1 summarizes the criteria of group classification,listing the adopted ranges of spectral index (β1), circumstellarluminosity, and the number of objects in each group. We verifiedthat stars in group 3 have the lowest levels of circumstellaremission (Sc < 10%), group 2 shows intermediate levels ofIR-excess (10 < Sc � 70%), and group 1 presents the highestlevels of Sc, that are greater than 70%. Figures 1 and 2 show someexamples of SED fittings that illustrate the SED shapes of thethree HAeBe groups. Different curves are used to represent thesame model components as defined in Figure 1. The integrationof the theoretical SED curves gives an estimation of individualcontributions, which were used to determine the circumstellarluminosity.

The suggestion of these three groups follows the idea of anevolutionary scenario for the HAeBe stars, suggested by Malfaitet al. (1998), similar to what is seen in TTs. In this scenario,group 1 corresponds to embedded objects, such as R CrA, and

Table 1Classification Based on the SED Shapes

Group β1 Sc HAeBe

1 (II/–) > 0 > 70% 46 (3)2 (I/I) −1 to 0 10% to 70% 43 (13)3 (III/II) < −1 < 10% 4 (0)

Notes. A comparison with HAeBe groups suggested by Hillenbrand et al. (1992)and Meeus et al. (2001) is respectively presented in the first column. The numberof PDS objects in each group and the number of previously classified HAeBes(between parenthesis) are also indicated in the last column.

group 3 represents the stars with more evolved dust disks, suchas Vega. The objects of group 2 are in between these bothextreme categories, showing SEDs similar to several knownHAeBe stars.

A comparison with other classifications shows that ourgroup 2 is similar to the group I of both Hillenbrand et al.(1992) and Meeus et al. (2001), where most of the knownHAeBes are found. Hillenbrand et al. (1992) classified as groupI the objects that have flattened SEDs in the near-IR (64% oftheir sample), whose IR excess follows a power law representedby λFλ ∼ λ−4/3. The authors suggest that an SED in theirgroup I seems to be well reproduced by a flat circumstellardisk model. Sources classified in Meeus et al.’s (2001) group Ihave continuum that can be reconstructed by a power law and ablack body, reproducing the suggested circumstellar geometry:an optically thin, flared region surrounding an optically thickdisk.

On the other hand, Meeus et al. (2001) suggest that sources intheir group II have a flat disk, whose SED shape is comparableto those of Hillenbrand et al.’s (1992) group III and our group3. In this case, the objects are similar to classical Be stars andshow the lowest values of IR excess.

Part of the objects in our group 1 show similarities whencompared to the objects of Hillenbrand et al.’s (1992) group II(23% of their sample), which have a gap in the near-IR (double-peaked SED) which is suggested to be due to dust distribution inan optically thin spherical envelope. A comparison of effectivetemperatures and (V −[12]) color excess shows that Hillenbrandet al.’s (1992) group II contains mainly Ae and late Be stars,having the highest values of color excess. This kind of SED isnot considered in the classification made by Meeus et al. (2001).

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Table 2Program Stars

PDS 2MASS ST AV β1 SED αIR Sc dedge Mass Age P Wλ(Hα) [V](mag) (%) (◦) (M�) (Myr) (%) (Å)

Group 1

018d 05534254−1024006 B7 2.3 1.26 ci 0.7 90.3 1.7 . . . . . . . . . 40019 05574947−1405336 B5 1.6 0.27 ci − 0.11 83.9 4.4 . . . . . . 0.18 ± 0.04 4023 06265390−1015349 OB 2.3 0.76 ci 0.6 87.5 −0.2 . . . . . . . . . 40024 06484168−1648056 A 0.6 0.1 dp −0.01 66.8 6.9 . . . . . . . . . 36027d 07193593−1739180 B2 2.4 1.09 ci 0.83 95.6 5.5 . . . . . . . . . 88034 08495851−4553055 B 1.4 0.34 ci −0.12 75.3 1.0 . . . . . . 2.34 ± 0.42 50037 10100032−5702073 B2 2.7 1.07 ci 0.44 93.2 −0.3 . . . . . . 2.78 ± 0.10 105067 13524285−6332492 B 2.5 1 ci 0.83 96.6 0.1 . . . . . . 2.73 ± 0.07 154[1.7]130 06495854−0738522 A 0.6 0.23 ci 0.36 76.3 3.9 . . . . . . . . . 40133 07250495−2545496 B? 1 −0.02 ci 0.01 71.1 −0.1 . . . . . . . . . 95139 11531324−6205210 B0 2.2 0.7 dp 0.62 88 1.6 . . . . . . . . . 2140 12174750−5943590 F2 0.7 0.23 ci −0.19 66.5 4.4 . . . . . . 1.11 ± 0.10 9[0.6]168 04305028+2300088 F0 3.4 2.25 ci 0.98 75.1 −0.2 . . . . . . 5.84 ± 0.28 8174d 05065551−0321132 B3 1.7 1 dp 0.59 96.7 0.1 . . . . . . . . . 65187 05320030−0455539 A 0.2 0.4 ci 0.53 83.1 1.3 . . . . . . . . . 18193a 05380931−0649166 B9 2.2 0.8 ci 0.22 70.3 −1.2 . . . . . . 2.25 ± 0.13 12198a 05385862−0716457 F0 1.5 0.35 ci −0.12 85.1 −0.9 . . . . . . . . . 9204d 05501389+2352177 B1 2.1 0.75 ci 0.6 95.2 6.5 . . . . . . . . . 245207 06071539+2957550 B? 1.5 0.94 ci 0.39 89 −0.4 . . . . . . . . . 5211 06101732+2925165 A 0.9 0.47 ci 0.36 75.1 0.5 . . . . . . . . . 35216d 06235631+1430280 B2 2.3 0.98 ci 1.03 91.8 3.8 . . . . . . . . . 200234 06584435−0341099 B? 2.4 1.6 ci 1.02 98.3 1.2 . . . . . . . . . 400241 07083879−0419048 B0 1.7 0.79 dp 0.2 98.6 2.5 . . . . . . 3.45 ± 0.04 65249 07241754−2616053 A 0.7 0.49 dp 0.17 84.4 0.4 2.0–3.0 1–5 2.25 ± 0.19 9250 07243697−2434474 B? 2.3 0.87 ci 0.28 85.2 −0.3 . . . . . . . . . 25255 07314884−1927334 B1 1.7 0.72 dp 0.77 93.5 3.0 . . . . . . . . . 300257 07414105−2000134 A 1 0.74 ci 0.53 88.6 3.1 . . . . . . . . . 15290b 09261450−5242048 A 0.7 0.78 dp 0.11 94.2 3.1 1.5–2.0 5–10 1.74 ± 0.23 7322 10520871−5612066 B2 0.8 0.71 dp 0.69 98.6 2.3 3.0–5.0 1–5 2.00 ± 0.07 −4324 10572427−6253132 B1 1.7 1.17 dp 0.82 91 3.6 . . . . . . . . . 4344 11403284−6432058 B5 0.7 0.21 dp −0.48 68.9 0.6 . . . . . . 1.78 ± 0.06 35353d 12222318−6317167 B5 1.8 1 ci 0.75 88.9 3.7 . . . . . . 2.17 ± 0.06 200361 13032149−6213262 B3 1.2 −0.01 dp −0.13 88.3 0.2 3.0–5.0 1–5 0.61 ± 0.03 12364 13200359−6223540 B? 1.1 0.4 ci 0.12 85.9 0.2 . . . . . . 0.33 ± 0.03 120389 15144705−6216597 A3 2 0.42 dp −0.17 58.7 0.3 0.5–1.0 1–5 5.28 ± 0.22 9[1.0]394 15351712−6159041 F0 0.4 0.79 dp 0.35 85.9 2.4 1.0–1.5 10 2.25 ± 0.17 −2406 16050392−3945034 A5 0.5 0.4 ci −0.28 82.4 −0.3 . . . . . . 3.20 ± 0.08 8[0.6]431 16545918−4321496 A0 0.8 1.04 dp 0.51 93.8 1.4 2.0–3.0 1–5 0.23 ± 0.18 −2[−0.2]465c 17451419−1756469 B 2 0.81 ci 0.43 92.2 0.6 . . . . . . 7.88 ± 0.25 110477 18003031−1647259 B1 2.3 0.97 ci 0.37 87.3 −0.02 . . . . . . 1.40 ± 0.09 120518a 18273952−0349520 OB 4.3 1.61 ci 0.92 90.6 0.03 . . . . . . 3.81 ± 0.10 374[0.6]520 18300616+0042336 F3 1.7 0.94 ci 0.29 84.3 −1.4 . . . . . . 3.66 ± 0.11 33530 18413436+0808207 A5 0.5 0.62 ci 0.48 81.7 1.3 . . . . . . 11.06 ± 0.65 28543 18480066+0254170 B0 2.9 0.54 dp 0.07 93.7 −3.4 . . . . . . 1.14 ± 0.04 1551 18552297+0404353 B0 2.9 1.59 ci 0.89 95.8 0.2 . . . . . . 11.00 ± 0.63 50581c 19361890+2932500 B1 1.2 0.74 ci 0.42 93.6 1.5 . . . . . . 12.22 ± 0.15 200

Group 2

002 01174349−5233307 F3 0 −0.84 dp −0.87 24.9 46 . . . . . . 0.16 ± 0.05 3[−0.4]004 03390056+2941457 A0 0.5 −0.43 dp −0.4 33.1 0.5 1.5–2.0 1–5 0.47 ± 0.04 9[−0.8]016 05431188−0459499 B9 0.3 −0.29 dp −0.11 45.7 1.0 2.0–3.0 <1 0.14 ± 0.05 37[1.4]020 05585578+1639573 B8 0.2 −0.37 dp −0.33 49 0.2 3.0–5.0 1–5 0.61 ± 0.03 41[1.0]021 06021488−1000595 B5 0.8 −0.14 dp 0.14 51.1 3.2 . . . . . . 0.87 ± 0.03 95022 06033705−1453025 A0 0.3 −0.42 dp −0.15 40.8 6.0 2.0–3.0 1–5 0.18 ± 0.02 25025 06542787−2502158 A3 0 −0.37 dp −0.48 44 6.3 . . . . . . . . . 13033 08484565−4048210 A 0.4 −0.26 dp −0.43 41.3 0.4 2.0–3.0 1–5 0.86 ± 0.08 18057a 11394445−6010278 A0 0.3 −0.53 dp −0.42 29 2.5 3.0–5.0 1–5 0.39 ± 0.03 9[0.6]061a 12000511−7811346 A0 0.4 −0.68 cd −0.61 28.4 1.7 2.0–3.0 <1 1.71 ± 0.11 24076a 15564002−2201400 F0 0 −0.38 dp −0.51 38.6 5.3 1.0–1.5 <1 0.73 ± 0.05 2[1.2]078a 16065795−2743094 A5 0.2 −0.57 cd −0.63 25.4 2.4 3.0–5.0 <1 0.09 ± 0.08 10[1.4]080b 16131158−2229066 F1 0 −0.39 cd −0.33 39.1 1.5 1.5–2.0 5–10 0.59 ± 0.08 −7 [0.8]095 17310584−3508292 B4 1.3 −0.09 dp −0.13 43.5 1.0 >5 <1 2.05 ± 0.10 27[0.7]096 17340462−3923413 B 0.9 −0.24 dp −0.19 58.1 0.3 . . . . . . 0.35 ± 0.03 54 [1.1]114 05270547+0025075 B9 0.3 −0.59 dp −0.46 27.8 3.4 2.0–3.0 1–5 0.13 ± 0.04 4[0.6]

No. 1, 2010 ANALYSIS OF THE PDS HERBIG Ae/Be CANDIDATES 31

Table 2(Continued)

PDS 2MASS ST AV β1 SED αIR Sc dedge Mass Age P Wλ(Hα) [V](mag) (%) (◦) (M�) (Myr) (%) (Å)

124 06065848−0555066 A 0.7 −0.05 dp −0.41 44.3 −0.9 2.0–3.0 1–5 . . . 20126 06133726−0625017 A7 0.3 −0.41 cd −0.21 50 0.6 2.0–3.0 <1 . . . 5134 07322662−2155358 B6 0.8 −0.32 dp −0.29 42.8 0.5 3.0–5.0 1–5 . . . 15172a 04584626+2950370 A3 0 −0.66 cd −0.68 25.4 3.6 1.0–1.5 <1 0.47 ± 0.07 22[1.0]176a 05160047−0948353 A0 0.3 −0.85 dp −0.82 41.6 4.6 . . . . . . 0.23 ± 0.03 1[0.1]179b 05240804+0227468 A0 0.3 −0.87 dp −0.86 25.3 4.5 2.0–3.0 1–5 0.11 ± 0.03 1[−0.1]180a 05244279+0143482 A8 0 −1.02 dp −1.03 25.2 5.2 2.0–3.0 <1 0.25 ± 0.03 2183a 05302753+2519571 A8 0 −0.63 dp −0.8 38.4 6.2 1.0–1.5 <1 1.90 ± 0.09 23184 05301903+1120199 A1 0.3 −0.62 cd −0.54 25 4.1 1.5-2.0 <1 0.21 ± 0.02 23 [0.7]191a 05374708−0642301 A1 0.1 −0.79 cd −0.77 19.1 −1.1 3.0–5.0 1–5 0.18 ± 0.01 5192 05380526−0115216 F0 0 −0.84 dp −0.85 43.3 0.5 1.0–1.5 <1 0.32 ± 0.04 11194a 05381450−0525132 A5 0 −0.84 dp −0.76 14.9 0.2 . . . . . . 2.33 ± 0.31 7201b 05441880+0008403 A7 0.1 −0.96 dp −0.98 39.6 −0.2 3.0–5.0 <1 0.46 ± 0.04 −2272 08114457−4405087 A0 0.2 −0.85 dp −0.77 37.5 2.4 2.0–3.0 <1 0.15 ± 0.03 7[0.4]277 08231185−3907015 F3 0 −0.82 dp −0.88 33.1 0.9 1.5–2.0 5–10 0.54 ± 0.05 3286 09060000−4718581 OB 6.4 −0.41 dp −1.1 20.9 0.3 . . . . . . 6.65 ± 0.06 70[1.2]297 09424032−5615341 A7 0.1 −0.39 dp −0.38 32 2.2 1.5–2.0 10 1.18 ± 0.06 −8303b 10025146−5916547 A0 0.2 −1 dp −1.05 35.2 1.7 3.0–5.0 1–5 0.17 ± 0.05 −7315 10330498−6019513 O8 0.9 −0.24 dp 0.17 80.3 1.3 . . . . . . 1.41 ± 0.04 1[0.7]339 11330559−5419285 A9 0 −0.71 dp −0.67 40 7.8 1.5–2.0 1–5 0.21 ± 0.06 0.5[0.1]340a 11332542−7011412 B9 0.2 −0.3 dp −0.13 52.3 4.1 2.0–3.0 1–5 3.41 ± 0.13 8[0.2]395 15404638−4229536 A8 0 −0.7 dp −0.71 30.4 5.0 1.5–2.0 5–10 1.81 ± 0.15 4[0.6]453 17205612−2603307 F2 0.1 −0.08 dp −0.31 66.9 1.8 . . . . . . 4.21 ± 0.06 −4[−1.5]469 17505809−1416118 A0 0.8 −0.05 dp −0.55 66.3 0.5 2.0–3.0 1–5 0.66 ± 0.03 2[1.2]473a 17562128−2157218 A0 0.2 −0.67 cd −0.57 30.7 2.0 1.5–2.0 5–10 0.74 ± 0.07 21[1.2]514 18242978−2946492 A0 0.5 −0.89 dp −1.04 27.1 3.9 1.5–2.0 5–10 0.80 ± 0.04 9[1.1]564a 19111124+1547155 A3 0 −0.44 dp −0.28 34.0 0.2 2.0–3.0 1–5 0.94 ± 0.09 10[0.8]

Group 3

185S 05321456+1703217 B8 0.2 −2.01 dp −1.92 1.6 9.6 . . . . . . 0.57 ± 0.11 −4225 06514575+0505037 B3 0.7 −1.72 dp −1.84 2 0.1 . . . . . . 0.95 ± 0.22 43281b 08554594−4425140 B5 0.8 −1.29 dp −1.52 9.7 0.1 3.0–5.0 <1 1.40 ± 0.05 −2[−2.4]327 11034056−5925590 B1 0.7 −1.3 dp −1.36 4.4 2.0 . . . . . . 0.34 ± 0.04 0

Notes. Columns description: (1) PDS studied stars; (2) 2MASS source identification; (3) spectral type; (4) visual extinction; (5–7) spectral indices (β1 and αIR) andSED shape; (8) circumstellar luminosity (Sc); (9) distance to closest cloud (dedge); (10 & 11) stellar mass and age; (12) intrinsic V-band polarization (P%); (13)equivalent width of Hα (absorption lines are represented by Wλ(Hα)< 0) and Hα variability [V].a Previously classified HAeBes (The et al. 1994).b Possible evolved HAeBe (Vieira et al. 2003).c Confirmed PPN objects.d Possible PPNs (Vieira et al. 2009).

Figure 2. Examples of HAeBe candidates classified as group 2 (PDS 126 and PDS 114), and an example of HAeBe candidate classified as group 3 (PDS 225). TheIRAS 100 μm flux is an upper limit.

Double-peaked SEDs are found in all our groups, when thespectrum decreases in the near-IR range and increases from themid- to far-IR range. This SED shape is related to a visiblecentral star (first SED peak) surrounded by a significant amountof circumstellar cold matter that produces the second peak in themid- to far-IR band. It is important to note that each group showsa different intensity in the second peak, which is extremely high

in group 1, similar to the first (stellar) SED peak in group 2, andvery low in group 3.

An absence of double peak is verified in other sub-types ofSED shapes. In the adopted classification, only group 1 containsstars whose SEDs increase in a “continuous” way. This can beinterpreted in terms of an important contribution of the warmdust in the circumstellar structure that emits mainly in the near-

32 SARTORI ET AL. Vol. 139

Figure 3. Distance to the edge of nearest clouds compared to Galacticextinction estimation. The higher extinction values are found for objects showingdedge < 0◦ (dashed line). The position of PDS 002 (dedge = 46◦) is indicatedby an arrow.

IR producing an excess in this band. In contrast, only group 2 hasstars whose SED is flattened or decrease in a “continuous” waythat means a decreasing of the warm dust contribution (near-IR)to the SED, when compared to the group 1 objects. Table 2 givesthe spectral index complemented with an information about theSED shape: double peaked (dp); continuously increasing (ci);or continuously decreasing (cd).

The results of the SED analysis in this large sample showthat objects having low β1 and Sc (group 3) are quite similar toclassical Be stars or HAeBes with disappearing disk, as verifiedby Hillenbrand et al. (1992) in their group III, and by Meeuset al. (2001) in their group II. The whole group 2 is comparableto group I of the other authors, containing typical HAeBes. Partof group 1 (16/46 stars) having β1 � 0.5 seems to be HAe stars,like group II of Hillenbrand et al. (1992). The remaining objectsof group 1 cannot be compared to the groups of other authors.Most of them have high Sc that could be due to interstellarcontamination, leading us to examine the possible associationwith clouds as described in the next section.

3. SPATIAL DISTRIBUTION AND ASSOCIATION WITHCLOUDS

In a study of the Galactic distribution of PDS HAeBe stars,Vieira et al. (2003) calculated minimum distances by assumingthat the candidates were main-sequence stars. Considering theproblems with this assumption, they adopted those distancesonly as indicative values to investigate possible associationswith SFRs. They verified that 78% of the candidates seem tobe related to one of the main SFRs, such as Taurus-Auriga,Orion, Canis Major, etc. In a few cases, they adopted reliabledistances from Hipparcos to confirm the association with SFR,in particular for nearby clouds such as Taurus, Chamaeleon, andOphiuchus. However, for most of the objects, the uncertaintieson distance determination do not allow a conclusive study aboutcloud association.

In our analysis, we aim to investigate the effects of interstellarclouds on far-IR observations and consequently on SED shapes.This possible contamination can be verified by comparing theposition of the stars projected against the position of dark clouds.This means that, in principle, isolated objects should not showpositional coincidence with clouds.

Dark clouds from catalogs by Lynds (1962) and Feitzinger &Stuwe (1984) were selected by choosing those that are closest to

Figure 4. Frequency distribution of angular distances of the HAeBe candidatesto the edge of the nearest dark cloud. Negative dedge values represent objectsthat are projected against the cloud area.

the candidate’s position. An estimation of the star’s distance tothe edge of each cloud (dedge) was obtained from the differencebetween the distance to the cloud’s center and cloud’s radius,roughly defined by dedge = d − 0.5 × A1/2, where d is theprojected distance from the star to the cloud’s center and A isthe area, obtained from the dark clouds catalogs. These angulardistances are shown in Table 2, where negative values of dedgeindicate a star position projected against the cloud area. A roughestimation of error is about 0.◦5.

Galactic extinction from Schlegel et al. (1998) was usedto check the validity of using the projected distance to theclouds. An estimation of color excess E(B −V ) is compared todedge in Figure 3. Only objects in high Galactic latitudes wereanalyzed, due to unreliable extinction obtained for positionswith |b| < 5◦. As expected, a good correlation is verified withextinction decreasing as dedge increases. Most objects that aredistant from clouds (dedge > 1◦) have E(B − V ) � 0.6; whileE(B − V ) � 1 for most of objects with dedge � 0.

Figure 4 presents the frequency distribution of dedge for eachgroup showing that 22/46 objects of group 1 are found closeto clouds (projected distances are less than 0.◦5). On the otherhand, 35/43 objects in group 2 have dedge > 0.◦5. Two stars inGroup 3 can be considered isolated; the others remain close tothe clouds (less than 0.◦5).

Large distances to clouds were verified for most of the 17PDS stars without clear association with SFR, as suggested byVieira et al. (2003). However, four of these stars has dedge < 0◦,meaning that only a few objects, which we considered close toclouds, are actually isolated.

We conclude that 48% of the objects in group 1 are apparentlyclose to interstellar clouds, since they have dedge < 0.◦5. Thesame can be said for a smaller fraction (19%) of objects ingroup 2. Almost 10 objects seem to be suffering effects ofcloud contamination, in particular those showing dedge < 0◦and AV > 1.5.

4. ANALYSIS

4.1. Infrared Colors and Spectral Index

The SED shapes used in the classification of the sample aredirectly related to the IR excesses observed in both, near- andfar-IR bands. The SED inclination and color–color diagrams inthese bands are usually evaluated in the comparison of objects

No. 1, 2010 ANALYSIS OF THE PDS HERBIG Ae/Be CANDIDATES 33

Figure 5. Distribution of near-IR spectral index, αIR, as a function of Sc for thethree SED groups based on the optical-mid-IR index. The divisions of the TTclasses are indicated by dashed horizontal lines and the approximate limits ofSc for the three groups are indicated by vertical dotted lines. Filled symbols areused to represent Sc derived from the flared disk calculation (groups 2 and 3only).

in different evolutive phases. For this reason, we discuss herethe IR characteristics of our sample.

The αIR spectral index, measured in the near- to mid-IRrange, was estimated for each star using the synthetic SEDcurve. This spectral index was first proposed by Lada (1987) toclassify TTs in three evolutionary classes, which are interestingto be compared to the HAeBe groups. Horizontal dashed linesare used in Figure 5 to show the separation of TT classes,adopting the revised values by Andre & Montmerle (1994).Different symbols represent the Sc values derived from theflat disk model (Gregorio-Hetem & Hetem 2002) and the flareddisk model (Hetem & Gregorio-Hetem 2007). As expected, thedistribution of αIR as a function of Sc, for the three groups(Figure 5), shows a good correlation. A negligible dispersionof the circumstellar luminosity calculated by different models(different circumstellar geometries) is verified, providing thesame separation of objects in three groups.

The adopted separation of the three HAeBe groups is quitesimilar to the TT classes, where most of group 1 objects arefound in the Class I region (αIR > 0), which contains the moreembedded objects (less evolved TTs). Group 2 coincides withClass II (−1.5 < αIR < 0), and group 3 objects are mainly inthe Class III region (αIR < −1.5).

Figure 6 shows the distribution of the three groups in adiagram of (J −K) versus (K −[25]) colors, which respectivelydescribe the shape of the stellar spectrum in the near-IR, and therelation of stellar and dust peaks (near- to mid-IR). This diagramwas suggested by Ueta et al. (2000) to indicate morphologicalcharacteristics of detached dust shells and different SED shapesof PPNs. The expected locations of PPNs are delimited by dottedboxes in Figure 6. In principle, dipolar nebulae, surroundedby an optically thick torus, have (J − K) > 1.5, whilebelow this limit are found objects with elliptical morphology,corresponding to an optically thin dust shell, in which thestarlight passes through in all directions.

Van de Steene et al. (2000) also studied the distribution ofPPNs in this diagram. They suggested that high values of (J−K)and (K − [25]) colors are expected in objects with high levelsof 25 μm flux that can be due to the presence of thick coolcircumstellar shell, or high temperatures of central star, whichdecreases the K magnitude. On the other hand, cool central stars

Figure 6. Color–color diagram of near-IR and 25 μm band flux for the sample,using different symbols to illustrate SED groups and shapes: double peaked (dp),continuously increasing (ci), and continuously decreasing (cd). The expectedlocation of PPNs is defined by dotted boxes.

and optical thin shells are expected to show low values of bothcolors.

From the distribution in Figure 6 we note that several objects,mainly from groups 2 and 3 in the region of (J − K) < 1.5and (K − [25]) < 8, are most probably in the PMS category,since their near- to mid-IR colors do not coincide with the locusof the PPNs (indicated by dotted boxes in Figure 6) studiedby Ueta et al. (2000). Different symbols are used in Figure 6 todenote the SED shapes. A clear separation can be noted for eachsub-group: objects from group 1 “continuously increasing” SEDmostly appear in the location of dipolar nebulae, while group 1“double peaked” SED coincides with elliptical morphologyobjects. Most of the group 2 objects, as well as group 3, arefound in the region of lower colors, outside the “PPN boxes.”The (J − K) versus (K − [12]) diagram does not provide anyadditional insight, since the distribution of the sample is quitesimilar to that observed in (J − K) versus (K − [25]). We havealso inspected if objects close to clouds (dedge < 0.◦5) have theirposition in Figure 6 affected by interstellar clouds. However,no particular trend is found when comparing near-IR colors andspatial distribution.

4.2. Stellar Parameters

Stellar temperatures are compared to Sc for the three groupsin Figure 7. The distribution of Sc is not the same for theHerbig Ae (spectral types later than B6) and for the Herbig Be(spectral types earlier than B5) stars. Although the HAe starshave Sc distributed along all values, most of them are objects ofgroup 2 (10 < Sc < 70%). On the other hand, most HBe starsshow Sc � 70%. Considering that group 1 may correspond toembedded, younger objects, this result is in agreement with theexpected faster evolution toward the main sequence of the HBestars compared to the HAe stars.

We have also analyzed the position of the candidates inthe near-IR color–magnitude diagram, aiming to verify thecorrelation of age to the circumstellar properties. Figure 8 showsthe distribution of the near-IR sources compared to isochronescalculated by Siess et al. (2000). By visual inspection, weverified the number of stars in between different intervals ofmass and age, but this analysis was constrained to about half ofthe sample (42/93), because 37 objects, mainly those showing(J − K) > 1.5, are not located in between the isochrones. We

34 SARTORI ET AL. Vol. 139

Figure 7. Distribution of Sc as a function of the Teff of the stars (symbols definedin Figure 6). Some spectral types are indicated at the top of the diagram. Thevertical dashed line indicates the adopted limit between HAe and HBe stars andthe horizontal dotted lines indicate the approximate limits of Sc for the threegroups.

Figure 8. Diagram of absolute J magnitude as a function of the (J − H )color (both corrected for extinction). Filled symbols indicate objects with(J − K) > 1.5. The Zero Age Main Sequence (ZAMS) and isochrones (0.1, 1,5, and 10 Myr) are indicated by full lines. Stellar masses 0.5, 1, 1.5, 2, 2.5, 5,and 7 M� are indicated by dotted lines.

also excluded 14 objects of this analysis with large uncertaintyon distance determination.

Figure 9 shows the distribution of the objects as a functionof their masses and ages, when compared to the spectral indexβ1. A particular trend in terms of age distribution is not verified,since both groups 1 and 2 are distributed in the whole range ofages. Only one object of group 3 could be placed in this diagram.

From the comparison of groups in Figure 9 it can be verifiedthat stars with M > 2.5 M� (triangles) are younger than 5 Myr.About 47% of the low-mass stars are older than 5 Myr and 41%are younger than ∼1 Myr.

4.3. Emission Lines

The profile of Hα was evaluated in order to have additionalinformation about the nature of the sample. Aiming to verifythe occurrence of variable emission lines among the studiedstars, we obtained low resolution spectra, in the λλ400–800 nmregion, for part of the PDS sample. The observations have beendone with the Perkin-Elmer 1.6 m telescope at Observatorio

Figure 9. Distribution of stellar ages as a function of β1 index, comparing tworanges of stellar masses: lower (dots) and higher (triangles) than 2.5 M�.

do Pico dos Dias (OPD), operated by Laboratorio Nacional deAstrofısica (LNA), Brazil, from 1996 to 1998. A CassegrainSpectrograph with a 300 l mm−1 grating was used, yieldingspectra with R in the range 1200–1500. The obtained spectra inthe Hα line region for 33 objects of our sample are shown inFigure 10.

In Table 2, the last column shows the equivalent widths ofHα line (Wλ(Hα), where negative values represent absorption),and an indicator of their variation is shown between brackets.This variation is estimated by dividing Wλ(Hα) measured in ourspectra by the values previously presented in PDS catalogs.Almost no variation is verified if Wλ(Hα) ratio is near 1;lower values correspond to a decrease in intensities; Wλ(Hα)ratios higher than 1 mean an increase in line intensity; andnegative values indicate an inversion of line profile (absorption/emission). For the objects not covered by our spectroscopicobservations we adopted the Wλ(Hα) published in the PDScatalogs (Gregorio-Hetem et al. 1992; Torres et al. 1995; Torres1999). In these cases, it was not possible to verify if the line isvariable or not. Only a few objects show variation larger than40%. Some of them suffer profile inversion: PDS 002, 004, 067,179, 281, 339, 340, 431, 453.

Vieira et al. (2003) suggested that six of the PDS HAeBecandidates could be objects at the end of the PMS, as indicatedby their Hα line profile. Based on our spectra we could confirmthis for PDS 080 only, whose Hα line remains unchanged. On theother hand, we verified that the Hα line of PDS 281 is extremelyvariable, which usually indicates the accretion process in PMSstars.

We have also compared the Wλ(Hα) to the (J − K) color,which are plotted in Figure 11. It is interesting to noteagain differences among objects showing larger near-IR excess((J − K) > 1.5), mainly those of group 1 with “continuouslyincreasing” SED, which show high intensities of Hα. The over-all trend is an increase of Hα line intensity as (J −K) increases.Considering that extreme values of Hα intensity are not typicalof PMS stars, we suggest that objects showing Wλ(Hα) � 70 Åand (J − K) < 1.5 are most likely young stars, i.e., mainlygroups 2 and 3, as well as part of objects in group 1 having“double peaked” SED (except PDS 255).

4.4. Optical Polarization

This section presents the intrinsic optical polarization of75% of our sample and how it correlates with the circum-

No. 1, 2010 ANALYSIS OF THE PDS HERBIG Ae/Be CANDIDATES 35

Figure 10. Low-resolution spectra obtained at OPD. The flux is given in log (λFλ) in Watt m−2 as a function of λ in μm.

Figure 11. Equivalent width of the Hα line compared to (J −K) color (symbolsdefined in Figure 6). Most of objects in groups 2 and 3 have (J − K) < 1.5 andWλ(Hα) � 70 Å.

stellar properties and the HAeBe classification previouslypresented.

An intrinsic polarization in HAeBe stars originates by thescattering of the starlight in asymmetric structures and hencecan be used to obtain information on circumstellar material. Thepolarization, however, is dependent on inclination of the shellrelative to the observer: for instance, a disk seen face-on presentsa null integrated polarization. More than that, geometricalproperties can be estimated even if the object is not resolved.The first polarimetric works on samples of HAeBes are fromBreger (1974), Vrba (1975), and Garrison & Anderson (1978).Possibly the largest statistical study of the optical polarization

of young stellar objects was done by Yudin (2000). He founda correlation between the polarization and the (V − L) excessand an evidence of a decrease in polarization with the PMSevolution.

Table 2 presents the values of the intrinsic V-band polarizationfor our sample as estimated in Rodrigues et al. (2009)—see thisreference for details on observations and reduction. Figure 12shows the intrinsic optical polarization of our sample as afunction of the circumstellar luminosity (Sc). There is a cleartrend toward an increasing of the polarization with Sc: group 3objects have the smaller polarization values, while group 1includes the more polarized objects. This can be interpreted asa decrease of the polarization with the PMS evolution, as notedby Yudin (2000). Hashimoto et al. (2008) have also proposedan evolutionary sequence of circumstellar structures in massiveyoung stellar objects.

The majority of the objects with P > 3% are associated with“continuously increasing” SED (filled symbols in Figure 12).This kind of SED probably represents the objects having anoptically thick dust torus, because the polarization needs anon-spherical material distribution in order to be produced.For group 1 objects, they may represent the objects in whichthe central star light is blocked by the circumstellar material.The group 1 objects with “double peaked” SED tend to havesmaller polarization. This can be explained by a large portionof light from the central source seen by the observer andcausing a dilution of the scattered emission responsible forthe polarization. Among the objects with the extreme values ofpolarization (P > 7%) we found the two confirmed PPN objectsof the sample (PDS 465 and 581). It poses the question if sucha high intrinsic polarization could be connected with a post-main-sequence nature. However, PDS 530 is highly polarized

36 SARTORI ET AL. Vol. 139

Figure 12. Intrinsic optical polarization compared to the circumstellar lumi-nosity (symbols defined in Figure 6). Vertical dotted lines show the groupsseparation and a vertical dashed line shows the P > 3% level.

Figure 13. Optical polarization compared to Se/Sc , the fraction of emissionfrom the cold optically thin envelope (symbols defined in Figure 6). Dottedand dashed lines are used to show that most of the objects with P > 3% haveSe > 0.5Sc .

Figure 14. Intrinsic optical polarization compared to the effective temperatureof the central objects (symbols defined in Figure 6).

and is in the direction of an SFR. Some confirmed HAeBes fromother works (e.g., Yudin 2000) also presented high values ofpolarization. Consequently, the polarization value is not enoughto distinguish the PMS from evolved objects.

Figure 15. Intrinsic optical polarization compared to the projected distance toclouds (symbols defined in Figure 6).

In Figure 13, we plot the intrinsic optical polarization as afunction of the fraction of circumstellar luminosity that cor-respond to the cold material, represented by Se/Sc. Objectswith high P(%) values tend to show a more important con-tribution from the cold (envelope) than from the warm (disk)material, considering that 9 out of 13 objects with P(%) >3 have Se > 0.5Sc. All these objects belong to groups 1 or 2.The temperature of the circumstellar material depends on its op-tical depth and on the distance to the central object. The highlypolarized objects may represent those with the largest quantityof circumstellar material and hence having the higher opticaldepths. This material must be non-spherically distributed in or-der to produce the observed polarization. Hence the materialproducing the polarization in these objects is not in the illumi-nated region of the disk, which is heated by the stellar emission.It may be concentrated in the internal region of a flared disk orin the external portion of a flat disk. An alternative is that thisdisk is very far from the central source. In this case, however,a “continuously increasing” SED as seen in most of the highlypolarized objects is not expected.

It could also be explored if the circumstellar asymmetry candepend on the mass of the central object. Figure 14 shows apossible tendency of massive objects to be less polarized thanHAe stars. If this is really the case it could be interpreted by adecreasing degree of circumstellar asymmetry for OB objects.However, the small number of massive objects can be biasingthe results.

In Figure 15, the optical polarization is compared to theangular distances to clouds, aiming to verify a possible relationwith spatial distribution. Objects showing P > 3% are mainlyprojected against or close to clouds (dedge � 1◦), but three ofthem are found far from clouds, which is not typically observedin young massive stars, since it is not expected that an objecthaving yet a large quantity of circumstellar material be far fromthe parent cloud.

5. CONCLUSIONS

Several tests have been applied to our sample aiming toconfirm their HAeBe young nature. Our results indicate that76% of the candidates can be actually considered PMS objects(mainly group 2), while the others (mainly group 1) need tobe studied in more detail, since several of them show typicalcharacteristics of evolved objects (e.g., PPNs). A summary ofthese results is presented below.

No. 1, 2010 ANALYSIS OF THE PDS HERBIG Ae/Be CANDIDATES 37

In order to verify possible interstellar contamination withthe far-IR measured flux, the spatial distribution of the candi-dates and their association with interstellar clouds have beenevaluated. We roughly estimated the projected distance fromthe HAeBe candidate to the edge (dedge) of the nearest darkcloud, allowing us to compare dedge to Sc and SED shapes. Weconclude that almost 62% of the sample can be considered iso-lated objects, and 38% are close to or associated with clouds(dedge < 0.◦5).

Only objects in group 2 have SED shapes (flattened or de-creasing slopes) typical of HAeBe stars, showing 10% < Sc <70%. They have a near-IR spectral index (αIR) similar to Class IITT stars. Group 3 objects probably are Vega-like objects or Bestars. Most of group 1 stars have αIR typical of Class I objects,but only part of them may be embedded in clouds (dedge < 0.◦5)and the others are isolated and may have an important (torus?)circumstellar envelope. High values of optical polarization arefound mainly for group 1 objects, but at least three stars ofgroup 2 have P > 3% indicating circumstellar asymmetry.

The near- and mid-IR color–color diagram shows a clearseparation between group 2 and group 1 stars, when comparedwith the distribution of evolved (PPN) objects. Only group 1stars coincide with the PPN location and most of them are inthe region of dipolar morphological type of nebulae, suggestedby Ueta et al. (2000), which seem to have a circumstellartorus. On the other hand, most of the objects in groups 2 and3 are concentrated to the left of the PPN region, favoring aconfirmation of their young nature.

Ages could be evaluated for most of the group 2 objects,which represent almost half of the sample and have less than 10Myr. A lack of information about ages for the other half of thesample (groups 1 and 3) does not allow us to achieve statisticalresults about the correlation of circumstellar matter and age.Consequently, it could not be verified a trend of Sc decreasingwith age, which would be expected in an evolutive scenario ofcircumstellar material disappearing with age. However, it seemsthat a separation occurs between low-mass and massive objects,according to their ages: stars with M > 2.5 M� are youngerthan 5 Myr, while the lower masses are found in two ranges ofages, �1 Myr or �5 Myr.

Most of the HAe stars of group 2 show Sc smaller than 50%of the total emitted flux. On the other hand, the HBe stars haveshown both configurations: some stars (mainly group 1) haveSc greater than 50% and others (groups 2 and 3) have smallervalues of Sc. These results are consistent with the fast evolutiontoward the main sequence of the HBe stars.

We verified a trend of the Hα line correlated with (J − K),in particular for objects of group 1 that have “continuouslyincreasing” SED shape, which show Wλ(Hα) � 70 Å. Theselarge emission of Hα has also been observed in PPN ob-jects, such as Hen3-1475 (PDS 465) and IRAS19343+2926(PDS 581), while most objects of group 2 show Wλ(Hα) � 50 Å,typical values found in several known HAeBes.

All the tests lead us to suggest that part of the HAeBecandidates listed by PDS require a more detailed analysis to havetheir PMS nature definitively confirmed, since several of themprobably are post-main-sequence objects. The work by Vieiraet al. (2009) is particularly dedicated to elucidating the numberof PPN candidates that are contaminating the HAeBe sample ofstars detected by PDS. According to their results for 25 stars ofour sample (mainly group 1 with β1 > 0.7): including PDS 465and 581, at least 8 objects show characteristics of PPNs; 16objects are more probably Class I PMS stars; and 1 object has

inconclusive results. From the results of the present work, wesuggest that other three objects (PDS 234, 241, and 255) ofgroup 1, with high Wλ(Hα) and large dedge, seem to be PPNcandidates, which were not analyzed by Vieira et al. (2009). Onthe other hand, 11 stars of group 1, mainly those with β1 < 0.5,having late spectral types or dedge < 0.5, and the Hα line typicalof HAeBes, are more probably PMS stars.

In total, we suggest that 71 objects of our sample (24 of themfrom group 1) can be more confidently considered young stars.Among the remaining objects of group 1 there are probably7 PPNs and 15 stars that deserve a more detailed analysis aimingto have their actual nature elucidated.

The authors thank Dr. C. A. O. Torres and Dr. G. Quast forhelpful discussions. M.J.S. thanks the Brazilian agency FAPESPfor the post-doc fellowship (No.00/06954-6). J.G.H. and C.V.R.thank partial support from FAPESP (Proc. No. 2005/00397-1and Proc. No. 2001/12589-1, respectively). This research hasmade use of the SIMBAD database, operated at CDS, Stras-bourg, France.

Facilities: LNA:1.6m (), LNA:BC0.6m ().

REFERENCES

Adams, F. C., Lada, C., & Shu, F. H. 1987, ApJ, 312, 788Andre, P., & Montmerle, T. 1994, ApJ, 420, 837Bowers, P. F., & Knapp, G. R. 1989, ApJ, 347, 325Breger, M. 1974, ApJ, 188, 53Cardelli, J. A., Cleyton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245de Jager, C., & Nieuwenhuijzen, H. 1987, A&A, 177, 217Dominik, C., Dullemond, C. P., Waters, L. B. F. M., & Walch, S. 2003, A&A,

398, 607Dullemond, C. P., Dominik, C., & Natta, A. 2001, ApJ, 560, 957Feitzinger, J. V., & Stuwe, J. A. 1984, A&AS, 58, 365Fernandes, M. B., & Araujo, F. X. 2003, in Open Issues in Local Star Formation,

ed. J. R. D. Lepine & J. Gregorio-Hetem, AP&SS, Library, Vol. 299(Dordrecht: Kluwer), CD-ROM

Garcıa-Lario, P., Manchado, A., Pych, W., & Pottasch, S. R. 1997, A&AS, 126,479

Garrison, L. M., Jr., & Anderson, C. M. 1978, ApJ, 221, 601Gregorio-Hetem, J., & Hetem, A. Jr. 2002, MNRAS, 336, 197Gregorio-Hetem, J., Lepine, J. R. D., Quast, G. R., Torres, C. A. O., & de la

Reza, R. 1992, AJ, 103, 549Hashimoto, J., et al. 2008, ApJ, 677, L39Herbig, G. H. 1960, ApJ, 43, 37Hetem, A., Jr., & Gregorio-Hetem, J. 2007, MNRAS, 382, 1707Hillenbrand, L. A., Strom, S. E., Vrba, F. J., & Keene, J. 1992, ApJ, 397, 613Kenyon, S. J., & Hartmann, L. W. 1987, ApJ, 323, 714Lada, C. J. 1987, in IAU Symp., Star Forming Regions, ed. M. Peimbert & J.

Jugaku (Dordrecht: D. Reidel), 1Lynds, B. T. 1962, ApJS, 7, 1Malfait, K., Bogaert, E., & Waelkens, C. 1998, A&A, 331, 211Meeus, G., Waters, L. B. F. M., Bouwman, J., van den Ancker, M. E., Waelkens,

C., & Malfait, K. 2001, A&A, 365, 476Natta, A., Grinin, V. P., & Mannings, V. 2000, in Protostars and Planets IV, ed.

V. Mannings, A. P. Boss, & S. S. Russell (Tucson, AZ: Univ. Arizona Press),559

Perez, M. R., & Grady, C. A. 1997, Space Sci. Rev., 82, 407Riera, A., Garcıa-Lario, P., Manchado, A., Potash, S. R., & Raga, A. C. 1995,

A&A, 302, 137Rodrigues, C. V., Jablonski, F. J., Gregorio-Hetem, J., Hickel, G. R., & Sartori,

M. J. 2003, ApJ, 587, 312Rodrigues, C. V., Sartori, M. J., Gregorio-Hetem, J., & Magalhaes, A. M.

2009, ApJ, 698, 2031Sartori, M. J., Gregorio-Hetem, J., & Hetem, A., Jr. 2003, in Ap&SS Library

299, Open Issues in Local Star Formation, ed. J. R. D. Lepine & J. Gregorio-Hetem (Dordrecht: Kluwer), 133

Savage, B. D., & Mathis, J. S. 1979, ARA&A, 17, 73Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525Schultz, G. V., & Wiemer, W. 1975, A&A, 43, 133Siess, L., Dufour, E., & Forestine, M. 2000, A&A, 358, 593Strom, K. M., Strom, S. E., & Merrill, K. M. 1993, ApJ, 412, 233

38 SARTORI ET AL. Vol. 139

Strom, S. E., Strom, K. M., Yost, J., Carrasco, L., & Grasdalen, G. 1972, ApJ,173, 353

The, P. S., de Winter, D., & Perez, M. R. 1994, A&AS, 104, 315Torres, C. A. O., Quast, G. R., de la Reza, R., Gregorio-Hetem, J., & Lepine,

J. R. D. 1995, AJ, 109, 2146Torres, C. A. O. 1999, Publication of CNPq/Observatorio Nacional (Brazil),

10, 1Ueta, T., Meixner, M., & Bobrowsky, M. 2000, ApJ, 528, 861

Van de Steene, G. C., van Hoof, P. A. M., & Wood, P. R. 2000, A&A, 362, 984Vieira, S. L. A., Corradi, W. J. B., Alencar, S. H. P., Mendes, L. T. S., Torres,

C. A. O., Quast, G. R., Guimaraes, M. M., & da Silva, L. 2003, AJ, 126,2971

Vieira, R. G., Gregorio-Hetem, J., Jr., & Hetem, A. 2009, AJ, submittedVrba, F. J. 1975, ApJ, 195, 101Waters, L. B. F. M., & Waelkens, C. 1998, ARA&A, 36, 233Yudin, R. V. 2000, A&AS, 144, 285