Evidence for Triggered Star Formation in the Horsehead Nebula

An Honors Thesis for the Department of Physics and Astronomy

Brendan Peter Bowler Tufts University, 2007

Thesis Committee Members: William H. Waller Brian M. Patten William P. Oliver

05/11/2007 Abstract

The formation of stars in molecular clouds is subject to both internal and external influences. Triggered star formation inherently involves an external agent— usually one or more massive stars whose radiation, winds, and eventual supernova(e) can have transformative effects on neighboring clouds. Examples of this triggering process are important for constraining models of induced star formation. The Horsehead Nebula (B33) is a nearby (400 pc) “pillar” of dense gas and dust protruding from the dark cloud L1630 and is seen in silhouette against the bright HII region IC434. I aim to characterize the state and extent of star-formation in this region in order to study the interaction between the massive O9.5V star σ Orionis and this nearby pillar. I present deep IRSF/SIRIUS JHKS and Spitzer/IRAC 4-channel imagery of the Horsehead Nebula. IR color-color and color-magnitude diagrams are used to identify young stellar objects (YSOs) based on their IR excesses and positions indicating youth. When possible, spectral energy distributions and spectral indices of bona fide and candidate YSOs are used to determine IR evolutionary classes. Four young stars are identified in this region. Two flat spectrum protostars are located at the surface of the western limb of the pillar. Two probable transition disk/Class III sources are detected near the base of the pillar and are probably not directly associated with B33. The evolutionary ages of the protostars are similar to the formation timescale of B33 (∼0.2 Myr), suggesting a common formation mechanism as the dense pillar first developed. The radiation-driven implosion model is the likely triggered formation mechanism of the two protostars, as this model satisfies the observed ages, locations, and luminosities of the YSOs in this irradiated cloud. Contents

1 Introduction 4

2 The Horsehead Nebula 11 2.1 Previous Work ...... 13

3 Observations 17 3.1 InfraRed Survey Facility Near-Infrared Observations ...... 17 3.1.1 The IRSF and SIRIUS ...... 17 3.1.2 Data Reduction and Analysis ...... 18 3.2 SST and Mid-IR Observations ...... 26 3.2.1 IRAC Observations, Data Reduction, and Analysis ...... 26

4 Results 28 4.1 Rejection of Extragalactic Contaminants ...... 28 4.2 Identiﬁcation of Young Stellar Objects ...... 31 4.3 Color-Color Diagrams ...... 32 4.3.1 Near-IR Color-Color Diagram ...... 32 4.3.2 Mid-IR Color-Color Diagrams ...... 36 4.3.3 IRSF and IRAC Combined Color-Color Diagrams . . . 39 4.3.4 Summary of Color-Color Diagrams ...... 42 4.4 Color-Magnitude Diagrams ...... 45 4.4.1 Near-IR Color-Magnitude Diagrams ...... 45 4.4.2 Mid-IR Color-Magnitude Diagram ...... 51 4.4.3 Summary of Color-Magnitude Diagrams ...... 51 4.5 Spectral Energy Distributions ...... 53

1 5 Discussion and Conclusion 61 5.1 The Horsehead as a Site of Triggered Star Formation . . . . . 61 5.2 Notes on Individual Sources ...... 64 5.3 Future Work ...... 66

2 List of Figures

1.1 NGC 602 in the SMC ...... 7

2.1 The Horsehead Nebula ...... 12

3.1 Linear Magnitude Correlation ...... 24 3.2 Magnitude Transformation Fits ...... 25

4.1 Rejection of Extragalactic Objects: Mid-IR CMD ...... 29 4.2 Rejection of Extragalactic Objects: Mid-IR CCDs ...... 30 4.3 The Near-IR Color-Color Diagram ...... 33 4.4 Horsehead Near-IR Color-Color Diagram ...... 35 4.5 IRAC Color-Color Diagrams ...... 38 4.6 IRSF and IRAC Color-Color Diagrams: A ...... 41 4.7 IRSF and IRAC Color-Color Diagrams: B ...... 43 4.8 Near-IR CMD: J vs J − KS ...... 49 4.9 Near-IR CMD: KS vs. H − KS ...... 50 4.10 Mid-IR CMD: (3.6) vs. (3.6)-(4.5) ...... 52 4.11 YSO Spectral Energy Distributions ...... 57 4.12 Uncertain Member Spectral Energy Distributions ...... 58 4.13 Sources in B33 ...... 60

3 Chapter 1

Introduction

The formation and early evolution of stars and stellar clusters is currently one of the most active ﬁelds in astrophysics. Observational and theoretical research primarily over the last thirty years has established a general theory of star formation that covers nearly all stages of a star’s pre-main sequence (PMS) life. However, while the broad notions of the collapse and early evolutionary phases of stars of varying masses are thought to be understood, most of the details remain obscure. For example, the fragmentation and accretion of dense cores in molecular clouds is still being investigated. The shape of the initial mass function, the physics of outﬂows from young stars in the form of Herbig-Haro objects, and disk dissipation mechanisms continue to remain somewhat enigmatic; grain growth and planetesimal formation are also just beginning to be observationally constrained. These topics are but a small portion of the unanswered questions relating to the local star-forming properties of our Milky Way galaxy. There remain untold numbers of related topics in extragalactic star formation studies. One subject that has long been speculated on but has only recently gained the serious attention of researchers involves the interaction between massive young stars and the nearby parental giant molecular clouds from which they were born. There is evidence for mass segregation in young open clusters that cannot be accounted for through dynamical interactions (see Clarke et al. 2000 and Lada & Lada 2003 for brief reviews). Rather than have mi- grated to the centers of clusters, it is thought that massive stars are born in the densest, centralized cores of molecular clouds and remain there throughout their relatively brief lives (a few Myr for O stars). It has also been shown that some young clusters are not in fact coeval, meaning their members were

4 not all formed at the same time. Instead there appears to be an age gradient across some clusters, with the younger members generally distributed in the clusters’ outskirts and older members located near the centers (Lee & Chen 2007; Lee et al. 2005; Chen, Lee, & Sanchawala). This age spread is probably dependent on the initial size and density distribution of the original molecular cloud. However, it is still not clear whether the formation of young clusters occurs in a bimodal fashion (with massive stars formed first and lower-mass stars following) or in a more complicated framework. For example, Hernándezet al. (2007) find a higher circumstellar disk fraction toward the center of the young 3 Myr cluster σ Orionis. The birth of massive young stars and the formation of open clusters with hundreds to thousands of members is the subject of much current observational and theoretical research. Recently Charles and Elizabeth Lada performed a statistical study on nearby (≤ 2 kpc) embedded clusters in molecular clouds (Lada & Lada 2003). While it has been known for decades that stars form in molecular clouds, this study quantitatively described the properties of these newly-formed clusters. They found that most stars are born in embedded clusters of 100 or more members, but that less than 7% of these clusters survive without complete dispersal to an age of 100 Myrs. These remaining clusters have total masses that are at least 50M and are dominated by low-mass members with masses between 0.1–0.6M . This peak of low-mass stars is in good agreement with other observationally derived initial mass function (IMF) studies in star forming regions, although the IMF turnover appears to be more closely centered in the 0.1M region (see Bonnell, Larson, & Zinnecker 2007 for a review of recent IMF results). There is also evidence for mass segregation in these nearby embedded clusters, with the most massive stars having formed in the clusters’ centers. Generally by 5 Myrs young clusters are no longer associated with the molecular clouds from which they were born (Lada & Lada 2003). A scenario that partly accounts for this cluster formation and its age spread (with older stars in the centers of clusters) involves a process of induced star formation caused by ionizing shock fronts at the surfaces of these eroding molecular clouds. Although this does not explain the results by Hernándezet al. (2007) regarding the disk fraction of the σ Orionis cluster being higher in the cluster’s center, it should be noted that projection effects of triggering star formation can still account for some of this distribution. For example, molecular cloud material in our line of sight to the cluster’s center may have undergone induced star formation and would therefore create the

5 appearance of disk-bearing young stars in the cluster’s central regions. Stars with masses between roughly 10–50M emit most of their radiation in the ultraviolet portion of the spectrum. These O and B stars make up OB associations, the younger counterparts to older open clusters. The environments surrounding these massive stars are characterized by ionized atomic hydrogen caused by photoionization and are called HII regions. As young clusters are generally still associated with their parental molecular clouds, these HII regions often extend to bright-rimmed clouds (BRCs) which mark the spatial extent of young clusters. Evidence for triggered star formation is partially based on the identiﬁcation of young stars in these BRCs. NGC 602 is a young star forming region in the Small Magellanic Cloud that ex- empliﬁes the scenario of a massive star producing an HII region marked by BRCs at its boundaries (Figure 1.1). These star-forming regions are often characterized by shells, bubbles, and cavities carved out and eroded by the ionization, winds, and supernovae of massive young stars. There are two leading methods by which triggered star formation is plaus- able in these environments. Chen, Lee, and Sanchawala (2006) outline these cases and describe current observational evidence for induced star formation. There is an interesting interplay that massive stars have in the triggering of further star formation in young clusters. On the one hand they have the destructive ability to both photoevaporate and dissipate molecular clouds, thus ridding of the material from which potential stars are formed and possibly regulating the masses of a newly-formed stars (Whitworth & Zinnecker 2004). On the other hand they may provide the necessary pressure to collapse a dense region that otherwise would not spontaneously have done so. In this way many cases of cluster formation are thought to proceed with a propagation of star formation from the centers to the outer reaches of young clusters. The “collect and collapse” model of triggered star formation (Elmegreen & Lada 1977; Hosokawa & Inutsuka 2006) involves the ionized HII regions surrounding massive stars expanding and sweeping up material to create a dense shell at the boundaries of molecular clouds. This shell of dense gas and dust then fragments and collapses to form new stars. This model has been invoked to account for many observational features of star forming regions. Zavagno et al. recently (2006) observed the RCW 79 Galactic HII region in mm, near-IR and Spitzer (GLIMPSE survey). They present evidence for a second-generation of triggered massive star formation by the layered collection and subsequent collapse of massive fragments. Sanchawala et al. (2007)

6 Figure 1.1: NGC 602 in the Small Magellanic Cloud as imaged by Hubble’s Ad- vanced Camera for Surveys in 2004. The bright blue-white stars in the center of this cluster are massive stars that have eroded the inner regions of the molecular cloud that they were born from. The remaining cloud regions are being photoevaporated by the intense UV radiation from the central sources (NASA, ESA, and the Hubble Heritage Collaboration). also ﬁnd observational evidence for this collect and collapse model of star formation. They present multiwavelength observations of the Trumpler 16 star-forming region in the Carina Nebula and, based on the spatial distribution of young stellar groups in the region, conclude that an expanding HII region probably accrued enough matter to gravitationally collapse and form a new wave of star formation. The other triggered star formation model involves the photoevaporation of a BRC. An ionization shock-front will develop at the surface of the cloud whose pressure causes the formation of clumps that then gravitationally collapse to form new stars. This so-called radiation-driven implosion (RDI) model (Bertoldi 1989) has been invoked to account for the observational properties of many star-forming regions. Surveys of IRAS sources associated with BRCs in the northern (Sugitani, Fukui, & Ogura 1991) and southern (Sugitani & Ogura 1994) hemispheres provide evidence for the RDI model

7 as a common mechanism of triggered star formation. Lee and Chen (2007) use this model to explain the properties of the λ Ori, Ori OB1, and Lac OB1 associations. Getman et al. (2007) use this model to explain the age gradient of young stars extending to the cometary globule IC 1396N (a case of “small-scale sequential star formation”). Urquhart et al. (2007) find that the RDI model can account for high-mass triggered star formation in the BRC SFO 75, demonstrating that any mass regulation from RDI does not exclude the formation of intermediate- to high-mass stars. The RDI model of triggered star formation is thought to account for star formation commonly seen at the illuminated surfaces of BRCs, Bok globules, cometary globules, and gaseous pillars at the edges of HII regions (Lefloch & Lazareff 1994; Gritschneder et al. 2006). Bok globules were first observed as “holes in the sky” by William Herschel (Bok 1977; Hoskin 1963) and were seen as optically dark features in an otherwise homogeneous background of stars. Although ranging in size and shape, Bok globules are generally small (r ' 0.35 pc), low-mass (11M ) (Clemens, Yun, & Heyer 1991) clumps of gas and dust that are common in star-forming regions (e.g., NGC 281, NGC 3603, and Thackeray’s Globules in IC 2944), but which are also seen in isolation (e.g., Barnard 68). Clemens, Yun, & Heyer (1991) extrapolate a survey of 248 nearby Bok globules to galactic scales and estimate that there are roughly 3.2×105 Bok globules in our galaxy that make up roughly 0.1% of it’s total molecular mass. Cometary globules, on the other hand, are generally seen in the HII regions near massive young stars. They appear su- perficially as comets generally with round “heads” and elongated “tails” and point radially towards the high-mass weathering sources. Zealey et al. (1983) document 29 cometary globules in the Gum-Vela star-forming regions, while Ogura & Sugitani (1998) present several dozen in the Orion OB1 complex. Star-formation has previously been observed in cometary globules (Reipurth 1983); the RDI model of triggered star formation may account for this phe- nomenon (Getman et al. 2007). Pillars (or elephant trunks) are commonly observed in star-forming regions and, although differing in shape and size, are thought to all have similar formation mechanisms. As a molecular cloud is photoionized by a massive young star, any pre-existing dense clumps will photoevaporate at slower rates than the surrounding material. A pillar will form when such a clump pro- tects the material behind it from photoionization (Gritschneder et al. 2006). Although these pillars are still being eroded away, they do so at slower rates than the cloud material they protrude from. However, the combination of

8 stellar winds and an expanding ionization front probably influences the formation of these structures as well (Schneps, Ho, & Barrett 1980). These pillars point radially in the direction of the weathering source and are often characterized by bright rims in Hα emission. The compression of molecular gas and dust induced by the photoevaporation of material on the surface of BRCs can trigger star formation through the RDI mechanism. Star formation at the tips of these pillars is a common indicator of the RDI triggered star formation scenario and has been observed in many pillars including the Elephant Trunk Nebula (Reach et al. 2004), the Eagle Nebula (Fukuda et al. 2002), and other analagous pillars (Sugitani, Tamura, & Ogura 1995). Evidence for triggered star formation usually comes from the spatial distribution and age gradient of young stars near a BRC (RDI model) or the observation of a dense, fragmented shell of gas with newly formed stars surrounding a massive young star (collect and collapse model). Dale, Clark, and Bonnell (2007) discuss the difficulties of determining whether young stars are actually induced to form or whether they are merely being exposed by the photoevaporation of their parental material. They perform simulations of both cases and conclude that there are no observational properties of the two populations that can be used to differentiate between these formation scenarios, although their low-number statistics may have prevented any trends from standing out. Indeed it is a difficult task to distinguish between these two scenarios and there are few ways to determine the mechanism by which a Class II (optically-thick disk with high accretion) or Class III (optically- thin disk with little or no accretion) source formed a posteriori. However a method to distinguish between spontaneous and triggered star formation has recently been proposed through the observational characteristics of protostars and their parental clouds. Dobashi et al. (2001) discovered a prominent difference between the luminosities of protostars in clouds near HII regions and those not associated with HII regions. Protostars near HII regions have consistently higher luminosities than otherwise over a wide range of cloud masses. Motoyama, Umemoto, and Shang (2007) perform numerical simulations to test whether this property can be accounted for through the RDI model of triggered star formation. They find that RDI can increase accretion rates in protostars by up to two orders of magnitude, substantially increasing the total (star plus accretion) luminosity of protostars in these regions. The luminosities of protostars can therefore be used in conjunction with the other observational properties of triggered star formation as a diagnostic tool for the formation mechanism of stars in that region.

9 In this paper I present near- and mid-IR observations of the Horsehead Nebula. The Horsehead is a dense pillar of gas and dust protruding from the giant molecular cloud L1630 and sits at the irradiated boundry of the HII region IC 434. At a distance of 400 pc it is the closest elephant trunk to earth and, still being physically connected to its parental molecular cloud, represents an ideal laboratory to study the evolution and the star-forming properties of such structures as well as the interaction between massive stars and molecular clouds. This paper is organized in the following manner. In Chapter 2 I present the Horsehead Nebula and previous work in the region. Chapter 3 introduces the IRSF and Spitzer/IRAC observations, including data reduction, calibration, and analysis. The initial extragalactic rejection criteria are also intro- duced here. Chapter 4 describes the results of the near- and mid-IR color- color diagrams as well as the color-magnitude diagrams. Finally, Chapter 5 discusses the implications from this study and the likely model of triggered star formation in this region.

10 Chapter 2

The Horsehead Nebula

The Horsehead Nebula in Orion was first described extensively by E.E. Barnard in 1913 (Barnard 1913) as a curiously dark object against a brighter background (see Pound, Reipurth, and Bally 2003 for an historical account). He subsequently included it in his catalogue of dark nebulae as B33 (Barnard 1919). Today the Horsehead is recognized as a relatively nearby (d ' 400 pc, Anthony-Twarog 1982) and in many ways unique case of a so-called elephant trunk or pillar commonly observed in star-forming regions. Protruding from the dark cloud L1630 (Orion B), the Horsehead has also been referenced as an emerging Bok globule (Reipurth & Bouchet 1984), although its evolutionary history has been debated (Pound, Reipurth, & Bally 2003; Hily-Blant et al. 2006). It is seen in silhouette against the bright Hα emission (HII) region IC434, itself photoionized by the massive O9.5 V star σ Orionis. The Horsehead is being photoevaporated at its western limb by the UV radiation of σ Ori, creating a BRC at its surface (Figure 2.1). σ Ori is at a projected distance of roughly 0.5◦ or about 3.5 pc from the Horsehead. σ Ori A is in a quintuple-system with another massive B0.5 star and several lower-mass companions. Together this system makes up the most massive member of the young 3–5 Myr σ Ori cluster. Its low-mass population has been extensively studied (e.g. Oliveira et al. 2006) down to the planetary- mass regime (Zapatero Osorio 2000) over several square degrees reaching close to the L1630 molecular cloud. Jeffries et al. (2006) used radial velocities and known distances to stars towards the cluster to find two superimposed but kinematically unique subgroups, each having different mean ages. Recently Hernándezet al. (2007) presented Spitzer Space Telescope observations of the region. They identify 336 members, many of which have circumstellar disks.

11 Figure 2.1: The Horsehead Nebula in Orion is a pillar of dense gas and dust protruding from the dark cloud L1630 and is seen in silhouette against the HII region IC434. It points radially towards the source of photoevaporation, the massive O9.5 star σ Ori (not seen in this image). North is left in this ﬁgure and West is up. Image Credit: T.A. Rector (NOAO/AURA/NSF) and the Hubble Heritage Team (STScI/AURA/NASA).

12 They brieﬂy comment on the discovery of a spatial extension of disk-bearing young stars towards the NGC 2024 region.

2.1 Previous Work

The Horsehead is referenced as a site of ongoing star formation throughout the literature, yet few studies have been conducted to directly determine the number and evolutionary stages of young stars associated with this dusty pillar. The unique structure of the Horsehead has elicited considerable debate in the literature regarding the nature of its formation. Most pillars in star- forming regions have bright-rimmed surfaces from the photodissociation, ex- citation and ionization of dense gas. Physically these pillars have relatively featureless “trunks” extending back to a parental molecular cloud, as is seen in the Eagle Nebula (M16), NGC 3603, and NGC 6357. Some pillars exhibit Herbig-Haro flows from embedded young stars in their tips, including pillars in the Pelican Nebula (Bally & Reipurth 2003) and the Carina Neb- ula. The Horsehead, however, has two physical separations on either side of the horse’s neck region representing the “jaw” and “mane” of the horse analogy. The characterization and formation of these regions has been the subject of several scientific papers. Reipurth and Bouchet suggest the jaw region was in fact formed from the outflow of a young star, specifically the IR source B33-6 (Reipurth & Bouchet 1984). Two subsequent papers refute this hypothesis based on polarimetric evidence (Warren-Smith, Gledhill, and Scarrott 1985) and spectra of the jaw region (Neckel and Sarcander 1985). The spectra by Neckel and Sarcander were nearly identical to that of IC434 and did not appear to be shock-excited as would be expected in the case of a collimated outflow. Pound, Reipurth, and Bally (2003) find a CO “U”- shaped feature that appears to wrap around the horses head with a northern leg tracing the “nose” of B33 and the southern leg following closely the shape of the jaw cavity. They discuss possible causes for this structure including an outflow from a possibly undetected Class 0 source. Hily-Blant et al. (2005) investigated the velocity field of the Horsehead and discovered a north-east velocity gradient and a rotation about the trunk connecting the horse’s head to L1630. They discuss the possibility of a pre-existing velocity field through which centrifugal forces created the “nose” and “mane” separations from the main pillar. However the formation history of the Horsehead is still far from

13 clear and there remains much work that can be done on the subject, including numerical simulations incorporating this velocity field. The physical properties of the Horsehead have been studied at longer wavelengths by several groups. Zhou et al. (1993) infer a gas density of roughly 104 cm−3 from molecular CS observations, a value confirmed by Kramer, Stutzki, and Winnewisser (1996). This latter group finds a temperature range of 20K in L1630 to 9K in B33 to <0.3K within 0.4pc of the western edge of B33 based on 12CO(2 − 1) observations. Lada et al. (1991) derive a virial mass for the Horsehead of about 35M while Pound, Reipurth, and Bally (2003) estimate the H2 mass at 27M . This latter group derives a survival timescale of the pillar using the estimated mass and mass-loss rate caused by photoevaporation. They find that the Horsehead is a relatively long-lived structure and will evaporate in roughly 5 Myr. Two sub-mm clumps (SMM J054089−02271 and SMM J054091−02278) at the western bright ridge and one (SMM J054112−02273) at the base of B33 are identified from 450µm and 850µm mapping by Johnstone, Matthews, and Mitchell (2005). Ward-Thompson et al. (2006) map the region in the same sub-mm wavelengths and find similar results, although the emission near the western limb (labeled B33-SMM1) is not resolved into two components. The emission near the base is labeled B33-SMM2 and is interpreted as being a pre-existing clump, unlike B33-SMM1 which is thought to have been compressed by the ionization front at the western surface of B33. These sub-mm continuum dust emission regions are often sites of current or eventual star-formation and are important tracers of such activity in star-forming regions. Star formation in B33 has been discussed by several authors (see Table 2.1). Reipurth and Bouchet (1984) identify B33-1 as an IR excess source. Three other regions of proposed disruptive activity related to star formation include the jaw region, the cavity region in the nose, and the center of reflec- tion nebulosity near B33-14. Pound, Reipurth, and Bally (2003) tabulate the fluxes of the three point sources detected by IRAS, including B33-1. IRAS 05386-0229 is noted as having a steeply-rising spectral energy distribution between 12µm and 60µm. Two X-ray detections with ASCA have been re- ported (Nakano et al. 1999). ASCA Source C-18 (B33-10) wasn’t found to have an IR excess by Reipurth and Bouchet (1984), and Source C-20 wasn’t previously detected in the near-IR. Evidence for triggered star-formation at the western surface of B33 has grown rapidly with the increase in observations of the Horsehead and with a better understanding of triggering processes in general. Reipurth and

14 – 1 –

Table 1. Previously Identiﬁed Candidate YSOs in B33

αJ2000.0 δJ2000.0 Source (h m s) (◦ ’ ”) ID References

05 40 51.72 −02 26 48.9a B33-1, IRAS 05383−0228 1, 2 05 40 56.8 −02 27 46 IRAS 05384−0229 2 05 40 57.3 −02 25 26 B33-10, [NYS99] C-18 1, 3 05 41 04.2 −02 23 37 [NYS99] C-20 3 05 41 11.1 −02 28 11 IRAS 05386−0229 2

Note. — References are numbered as follows: 1)IR excess in Reipurth & Bouchet (1984), 2)IRAS detection quoted in Pound, Reipurth, & Bally (2003), 3)ASCA detection in Nakano et al. (1999) aCoordinates from 2MASS Point Source Catalogue

Bouchet (1984) first suggested this scenario as a model for the formation of B33-1. The notion of triggered star formation in the Horsehead has since been echoed by other authors (Pound, Reipurth, and Bally 2003; Ward- Thompson et al. 2006). The B33-SMM1 clump is thought to have formed from the pressure of an ionization shock front at the surface of this pillar, a scenario that is important in the RDI model of induced star formation. The illuminated surface at the western edge of B33 is likely the sharpest IR bright rim in our Galaxy as seen by ISOCAM, according to Abergel et al. (1999) and Abergel et al. (2003). Timescales are also relevant in the determination of whether any young stars predate the formation of B33. Gritschneder et al. (2006) recently demonstrated that pillars of roughly 1 pc in length can form in only 3×105 yrs from the ionizing UV flux of massive young stars. Therefore in a pillar such as the Horsehead (∼0.7 pc in length) any young stars formed through a triggering scenario shouldn’t be older than the pillar formation timescale (∼2×105 yrs), roughly corresponding to the evolutionary ages of Class I sources (White et al. 2007). If a young star is significantly older than this timescale then it is likely that the star was in fact formed before the development of the pillar and is simply being exposed by photoevaporative erosion. Hernández et al. (2007) briefly comment on a spatial extension of

15 IR excess sources to the northeast of σ Ori in the direction of NGC 2024 and possibly B33 as well. Megeath et al. (2005) present the preliminary results of their Spitzer survey of the Orion molecular clouds and find candidate young stars that appear to be associated with the B33. These observations together follow the predicted consequences of RDI-triggered star-formation possibly for the western illuminated interface of the larger L1630 region. Specifically, the formation of dense clumps which then gravitationally collapse to create new stars will leave a “trail” of more evolved young stars in the direction of BRCs as the original clump is photoevaporated away. However more studies near the boundary between IC434 and L1630 are necessary to confirm this scenario. In this study I aim to determine the state and extent of star formation in the Horsehead Nebula using a combination of near-IR (JHKS) and mid- IR (Spitzer/IRAC) imagery. These wavelengths are particularly sensitive to disk-bearing objects, as thermal disk emission becomes important over photospheric emission in the near-IR and dominates in the mid-IR. The spatial distribution of confirmed young stars in the region, their IR classes, and their spectral energy distributions will be used to asses the evolutionary states of the YSOs as well as the RDI model of triggered star formation in this region.

16 Chapter 3

Observations

3.1 InfraRed Survey Facility Near-Infrared Observations

Observations at the 1.4m InfraRed Survey Facility (IRSF) in South Africa were kindly obtained on March 13th, 2002, by Dr. Motohide Tamura during an observing run. To account for bad pixels in the infrared array the observations were taken in a dithered pattern. In each filter three sets of ten 30s dithered images were obtained in the following manner. An image was first taken while centered on B33, then nine more were obtained by moving the telescope in a circular manner about the central image, each new position offset by only a dozen arcseconds or so from the central one. This creates an effective integration time of 900s for the coadded images in each filter. See Section 3.1.2 for details on the data reduction process.

3.1.1 The IRSF and SIRIUS Located in Sutherland, South Africa, the IRSF 1.4m telescope is a joint collaboration between Nagoya University, Kyoto University, the National As- tronomical Observatory of Japan (NAOJ), and the South African Astronom- ical Observatory (SAAO). Its primary detector is the Simultaneous-3color Infrared Imager for Unbiased Survey (SIRIUS), where focused light passes through two dichroic mirrors to the three J (λ = 1.25 µm), H (λ = 1.63 µm), and KS (λ = 2.14 µm) near-infrared filters. These are the standard ground-based near-infrared filters used in astronomy and allow specific wave-

17 bands to pass through nearly unimpeded while blocking out all wavelengths outside that bandpass. Each corresponds to a “window” in the earth’s atmo- spheric transmission. Unlike many other telescope instrument designs that have filter wheels to interchange individual filters, SIRIUS allows for the simultaneous imaging of astronomical objects in all three near-IR filters. Each HgCdTe (HAWAII) infrared array contains 1024×1024 pixels with a pixel scale of 0.45” per pixel. The field of view is 7.8’ (Nagashima et al. 1999, Nagayama et al. 2003).

3.1.2 Data Reduction and Analysis The reduction of data from infrared arrays is similar to that for CCDs. There is a dark current that is created in the readout process from thermal electrons in the material. There is also a bias frame that is inherent in infrared arrays and forms a “pedestal” on which the source electrons accumulate. Both of these can be subtracted from the raw data, but there is inherent noise in these frames that add to the total noise of the image. The sky background is another additive feature that can be accounted for by imaging another “blank” region of the sky and subtracting the two frames after taking into account any sources that might be present in that “blank” region. There are also pixel to pixel variations which can be normalized by uniformly illu- minating the array through the imaging a flat field screen or the telescope dome. This “flat field” must be divided out to normalize the response of all of the pixels (Howell 2000; Glass 1999; Romanishin 2002).

Pipeline and Manual Reduction

The dark current subtraction, sky subtraction (centered at αJ2000.0 = 05:41:40.9 and δJ2000.0 = −02:26:34) and the flat field division were performed on site at the IRSF through a pipeline reduction procedure, outlined in Nagashima et al. 1999. Although the pipeline software also shifts the individual images, com- pensates for bad pixels, and coadds each of the dithered frames, the resulting coadded images for each of the three filters have nonuniform effective integration times across the images, resulting from the dithering process and affecting only the sides of the images. To get a uniform image the reduced frames were manually shifted, trimmed and coadded using the IMAGES/IMMATCH/IMALIGN and IMAGES/IMMATCH/IMCOMBINE

18 packages in the Image Reduction and Analysis Facility (IRAF).1 The resulting coadded images have an eﬀective integration time of 900s over the entire 7.8’ ﬁeld of view.

Source Detection The source detection algorithm daofind in the NOAO/DIGIPHOT/APPHOT package in IRAF was run separately for each of the three coadded images. The daofind task searches an image for local intensity maxima at user-defined values above the local background. It also computes roundness characteristics and filters out nonastronomical objects in the output. A visual inspection of the images was also done to ensure that no sources were missed and that obvious contaminations were excluded. The criteria for source detection included a 5σ detection in at least two bands and final J −H and H −KS color errors less than 0.2 mag (see below for details). The former criterion ensures sources are indeed real and is used as an initial screening. A total of 160 sources satisfied this requirement. The latter criterion was used after color transformations and associated errors were determined and applied to all sources so that the J −H and H −KS colors would have final errors less than 0.2 mag. A total of 134 near-IR sources satisfy both criteria. The corresponding limiting magnitudes from these criteria are roughly 19 in the J band, 18 in the H band, and 17.5 in the KS band. For reference note that at the distance of B33 (roughly 400 pc), the 1 Myr brown dwarf limit of 0.08M is about 13 in the KS band, well within reach of this study even for very reddened stars.

Stellar Photometry Stellar photometry is the process of determining source magnitudes from the photon flux detected in an instrument. This process is complicated by many factors. The response of any detector is never perfect, both in the photoelectric effect that generates the stored electrons and in the readout process itself. Bright sources saturate pixels and their accumulated electrons can “spill out” or “bleed” onto other pixels. Diffraction spikes from secondary

1IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

19 mirror supports often contaminate pixels. Some pixels don’t work properly; cosmic rays can ruin single pixels or rows of pixels. Other artifacts can arise, but most of these problems are accounted for in the reduction of the images. The shape of the radial profile of point sources on infrared arrays and CCDs is called the point spread function (PSF). The PSF is roughly Gaus- sian in shape and has wings that quickly decay, but that extend far from the source (to “infinity”). Therefore any attempt to determine the total number of photon counts from a point source must take into account the extent of the wings.2 There is also a background flux of photons that must be subtracted from the counts of a stellar or extragalactic source. These photometric problems are generally solved using one of several methods, including aperture photometry and PSF modeling. The method employed for the extraction of photometry in the IRSF images involves summing the total number of photons in a circle of a given radius centered on a source. Then the averaged pixel value of background counts in a given annulus surrounding the aperture circle is multiplied by the number of pixels within that circle and subtracted from the total photon count. In this way most of the source’s counts are included in the inner circle while the counts in the annulus are mostly those due to the sky background and noise. It is also useful for regions where the background sky is not constant throughout the image, as is the case for molecular clouds. This technique is also useful for uncrowded fields where the PSFs of sources don’t significantly overlap. A common measure of the PSF is the full width at half-maximum (FWHM) value. The FWHM is the width of the full PSF at half the maximum PSF value. A nice feature of PSF profiles is that both bright and faint stars have the same PSF shape and size. The only difference is in the intensity of the source. Moreover, the net flux at some radius of a point source includes the same percentage of total stellar flux independent of source intensity (Ro- manishin 2002). So an 11th magnitude star and a 16th magnitude star have the same percentage of flux contained in a given source aperture. As long the same source aperture is used for every star in an image, and because all these sources are later calibrated to an external filter system, the relative magnitudes are the important factors for the determination of calibrated

2In this discussion I will refer to photons in an image, rather than the real electrons that are read out to determine pixel values which make up the image. The photon response is assumed to be linear with the number of photons, so whatever the eﬃciency of the detector might be it will apply equally to all pixels and over a large range of magnitudes.

20 magnitudes. The FWHM values for the IRSF images were determined from an average of roughly a dozen bright, typical, and faint sources. They are 2.7 pixels for the J filter, 2.6 pixels for H, and 2.5 pixels for KS (at 0.45” per pixel this corresponds to FWHMs of roughly 1.22”, 1.17”, and 1.13” for the three respective bands). There is an optimum aperture radius for which to extract photometry. A radius which is too small will not include sufficient stellar flux to precisely determine the source’s magnitude, despite the relative magnitudes being the important quantity for this work. The reason is that the greater the percentage of source flux the more accurate is the flux determination. Too wide a source aperture will include more and more sky noise, which cannot be subtracted from the net flux. A simple way to find the optimum radius is by keeping the background annulus (and hence sky value) constant while varying the inner source aperture. The signal to noise ratio will vary throughout this procedure and it’s peak value indicates the optimum source radius for that star (Romanishin 2002). This process was performed for 14 sources ranging from moderate to low values in apparent magnitude (low magnitudes corresponding to brighter sources). Optimal radii were 6-8 pixels for moderately bright sources and 10-11 pixels for bright sources. As most of the sources in the image are relatively faint, an aperture radius of 4 pixels was chosen (equivalent to 1.8”). It was often the case that the flux from one source “contaminated” the background sky counts inside the annulus of the other source. Sources which were close together were accounted for by creating a 3σ threshold above the median sky value over the entire image. Any pixels above this value were discarded in the determination of the local sky value. Based on a visual inspection of the images, no two sources significantly contaminated the flux inside the aperture radius of any other sources. Photometry was extracted with the APPHOT package in IRAF. An aperture radius of 1.8” was used with an annulus of inner radius 4.5” and an outer radius of 7.2”. A default arbitrary zero-point offset of 25 was used. As the photometry was later calibrated to bright 2MASS sources within the IRSF field of view there was no need to perform an aperture correction to account for flux lost in the wings of the PSFs. The PSF shape was constant over the course of the observations, as the airmass varied only slightly (from 1.18 to 1.25). Two bright sources were saturated in all three bands as was apparent in their surface plots (2MASS J05405609-0224021 and 2MASS J05404959- 0228561). One source corresponding to a previously identified candidate

21 YSO near the tip of the Horsehead pillar was saturated in the KS band (2MASS J05405172-0226489). 2MASS magnitudes are used for these three sources.

2MASS Calibration The Two Micron All Sky Survey (2MASS) conducted a near-IR survey in the J, H, and KS bandpasses covering 99.998% of the sky down to 10σ limiting magnitudes of roughly 15 for all three photometric bands (Skrutskie et al. 2006). Their data have been made available online for public and scientific use.3 Because of the extent and depth of its survey the 2MASS filter system has become yet another standard IR filter set commonly used and referred to in near-IR astronomy. It is therefore an appropriate system to calibrate near-IR observations to when lacking data of standard stars. The criteria for selecting 2MASS point sources in the 7.8’ IRSF field of view included the best photometric quality flags (labeled “A” in the 2MASS Point Source Catalogue, with a SNR ≥ 10) in at least two bands and a flag of at least second-best photometric quality in the third band (labeled “B,” with a SNR ≥ 7). An additional criterion of sources having magnitudes between roughly 12 and 15 in the KS bandpass was imposed, as 2MASS was most sensitive to those magnitudes in that band. A total of 23 2MASS sources satisfying these criteria were found and were used to calibrate the IRSF magnitudes. To transform colors and magnitudes from the measured IRSF instrumental magnitudes to the standard 2MASS filter system, the basic transformation equations for photometry were used. The filters, detectors, observing locations, altitude, and weather at the IRSF inevitably differed from the 2MASS filters and observing conditions, so there may be a slight color dependence on the transformed magnitudes (Romanishin 2002; Zissell 1998). The transformation equations are below.

J = cJ (J − H) + j + zJ

(J − H) = cJ−H (j − h) + (j − h) + zJ−H

(H − K) = cH−K (h − k) + (h − k) + zH−K Here capital letters denote the 2MASS magnitudes while lower-case letters denote the measured IRSF instrumental magnitudes. The zX term denotes 3www.ipac.caltech.edu/2mass

22 the zero-point offsets for the X -band magnitude or the X color index. These equations are linear, as is expected if the measured response to stellar flux is the same in both IR arrays. To test this linearity between the two systems I’ve plotted the 2MASS vs. IRSF instrumental magnitudes for all 23 sources. A linear least-squares fit is indeed a good fit to the data (Figure 3.1). However, if the filters, detectors, and all other conditions were exactly the same for both the IRSF and the 2MASS observations then we would expect the cJ coefficient to equal zero, as there would be no color dependence on the j magnitude. Similarly, cJ−H and cH−K would also be zero. Rearranging these equations, we get the following:

[J − j] = cJ (J − H) + zJ (3.1)

[(J − H) − (j − h)] = cJ−H (j − h) + zJ−H (3.2)

[(H − K) − (h − k)] = cH−K (h − k) + zH−K (3.3) These points are computed for all 23 2MASS and corresponding IRSF sources and are plotted in Figure 3.2. A linear weighted least squares fit is applied to determine the coefficients and zero-point offsets. This fit uses the 2MASS errors on the y-axis as these are generally appreciably larger than the IRSF measured errors. These errors are derived from the above equations and are propagated using the normal quadrature method. The transformation coefficients are listed below:

cJ = 0.03927 ± 0.02406

zJ = −7.97462 ± 0.01807

cJ−H = 0.07013 ± 0.02924

zJ−H = −0.23701 ± 0.02709

cH−K = −0.05466 ± 0.04502

zH−K = 0.75172 ± 0.02327

Clearly there is very little color dependence, as all cX coeﬃcients are nearly zero. The transformation equations (3.1), (3.2), and (3.3) are applied with the above coeﬃcients to all 131 IRSF sources with extracted photometry (the three saturated sources in the IRSF images used the original 2MASS photometry).

23 Figure 3.1: 2MASS vs. IRSF instrumental magnitudes for all 23 selected sources. Because there may be a slight dependence on color in this relationship, these plots are merely included to show the linear trend between these magnitude systems. See Figure 3.2 for the derivation of transformation equations between the two systems.

24 Figure 3.2: A weighted linear least squares fit is used to find the transformation coefficients between the IRSF instrumental magnitudes and 2MASS magnitudes for all 23 sources.

25 Source Coordinate Determination The IRSF does not have the ability to apply a world coordinate system to the header of the Flexible Image Transport System (FITS) images it outputs. However it is possible to determine the right ascension (RA) and declination (DEC) of sources in these images by using the accurate 2MASS sources previously used for the photometric calibration. The Astrom Fortran program4 is a basic astrometry calculator that uses known positions in RA and DEC with x- and y-pixel coordinates to compute the RA and DECs of other pixel coordinates using diﬀerent interpolation models. Astrom was applied to all 134 detected sources in the IRSF images to determine their coordinates in the J2000.0 epoch.

3.2 Spitzer Space Telescope Mid-Infrared Observations

The Spitzer Space Telescope is a 0.85m telescope that can perform imaging and photometry from 3.6–160µm, spectroscopy from 5.2–38µm, and spec- trophotometry from 51–106µm5. It was launched into an earth-trailing he- liocentric orbit in August 2003 and has an estimated cryogenic lifetime of 5.5 years. The Infrared Array Camera (IRAC) on board Spitzer is a four- channel camera that obtains simultaneous broad-band images in 3.6, 4.5, 5.8, and 8.0µm (Fazio et al. 2004a).

3.2.1 IRAC Observations, Data Reduction, and Analysis Spitzer/IRAC observations of the Horsehead were obtained as part of the program An IRAC Survey of the L1630 and L1641 (Orion) Molecular Clouds (PI: S.T. Megeath; PID43)6. Two epochs of observations separated by about six months consisted of two 12s “high dynamic range frames.” These frames consist of a 12s frame with 10.4s of integration time and a 0.6s frame with

4Distributed through the Starlink Software Collection which is currently run by the Joint Astronomy Centre of the UK Science and Technology Facilities Council. 5See www.spitzer.caltech.edu for details. 6The information regarding IRAC observations of the Horsehead as well as stellar photometry of sources in the region have been kindly provided by S. Thomas Megeath

26 0.4s of integration time. The total integration time is 41.6s for the long frames and 1.6s for the short frames. Photometry of point sources was initially obtained from the Basic Cal- ibrated Data (BCD) from the Spitzer Science Center S14 data pipeline. This BCD data were mosaiced using the WCSmosaic IDL package developed by Rob. Gutermuth. This software rejects cosmic rays by identifying deviant pixel values in overlapping frames, applies a frame-by-frame distor- tion correct, derotation, and subpixel offsetting in a single transformation, and matches background by comparing signal levels in overlapping images. The photometry was extracted using PHOTVIS, an IDL photometry and visualization tool developed by Rob. Gutermuth. This program performs a spatial filtering of the data, computes a noise map from the filtered data, and then searches for point sources that exceed a source detection threshold. Aperture photometry is obtained for each of the sources using an aperture of 2 pixels and a sky annulus of 2-6 pixels. Each pixel is 1.2” × 1.2”. The data were then bandmerged and sources were considered coincident when they were within 1”. If the two sources satisfied that criterion then the closest sources were then identified as coincident. The zero points used to obtain the magnitudes (in DN/S per BCD image) were 19.6642, 18.9276, 16.8468, and 17.3909, for the 3.6, 4.5, 5.8, and 8.0µm bands, respectively. The limiting magnitudes are [3.6] ' 17.0, [4.5] ' 16.6, [5.8] ' 15.8, and [8.0] ' 13.5 for errors ≤ 0.1 in the first two IRAC channels and ≤ 0.2 in the last two channels.

27 Chapter 4

Results

4.1 Rejection of Extragalactic Contaminants

Various methods have been developed to differentiate young stars from background galaxies or older foreground stars using IR photometry. Objects that one researcher finds interesting might be considered contamination by another. In this case background galaxies may be contaminating the field and need to be rejected from our sample. Often galaxies do not show the same extended structure in the IR as they do in the optical, and their bright nuclei may resemble stellar sources. This is especially true of galaxies with active galactic nuclei (AGN). Moreover, galaxies with high star formation rates may occupy the same regions on near-IR color-color diagrams that young stars do (Geller et al. 2006). Hernándezet al. (2007) outline two methods for detecting extragalactic objects from IRAC data. A source is considered to be an extragalactic contaminant if its position on either the CCD or CMDs are indicative of a background contaminating source. Roughly half of the sources in the [3.6] vs. [5.8]−[8.0] diagram above [3.6] = 14 are extragalactic in nature (Fazio et al. 2004b). Figure 4.1 plots this color-magnitude diagram for all sources detected with errors ≤ 0.1 in channels [3.6] and [4.5], and errors ≤ 0.2 in channels [5.8] and [8.0]. There is one source above the [3.6] = 14 limit that corresponds to a source that had previously been rejected in the near-IR rejection criteria based on large color errors. The other rejection scheme involves identifying sources in a particular region of the [3.6]−[5.8] vs. [4.5]−[8.0] and the [4.5]−[5.8] vs. [5.8]−[8.0]

28 Figure 4.1: The initial rejection of extragalactic objects includes sources with 3.6µm band magnitudes greater than 14 in the [3.6] vs. [5.8]−[8.0] diagram. One source is rejected based on this criterion.

29 Figure 4.2: The [3.6]−[5.8] vs. [4.5]−[8.0] (top) and [4.5]−[5.8] vs. [5.8]−[8.0] (bottom) color-color diagrams are used to distinguish reddened stars from background galaxies. The regions redward of the solid lines in the longer-wavelength colors for each plot are dominated by extragalactic sources (Gutermuth et al. 2007; Hern´andez et al. 2007). No sources from this sample occupied these regions.

30 color-color diagrams. Gutermuth et al. (2007) identify this region as being dominated by polycyclic aromatic hydrocarbon (PAH)-rich galaxies with active star formation. This characteristic extragalactic region was confirmed by Hernández et al. (2007). These plots are shown in Figure 4.2, where the regions redward of the solid lines in the longer-wavelength colors are dominated by background galaxies. No sources from this sample occupied these regions, and the source which was identified as a probable contaminant in Figure 4.1 did not appear in this region. However, because this source is fainter than would be expected for a non-embedded (based on visual inspection), low-mass young star, and because it failed the first criterion, it is removed from the source list.

4.2 Identiﬁcation of Young Stellar Objects

A source is considered a YSO if there is evidence for an excess of IR emission which cannot be accounted for by photospheric emission alone for those sources not determined to be an AGN or other extragalactic contaminants (see section 4.1 for mid-IR AGN rejections). The positions of sources on color-color diagrams can be used to identify physical properties of the sources themselves. The color of a star is independent of its distance (neglecting reddening), where this color index is simply the difference of magnitudes (the logarithm of the ratio of fluxes) between the magnitudes of two different wavebands. Furthermore, color-magnitude diagrams can be used to assess the masses and ages of stars. However, reddening from starlight scattered by dust between the earth and the source as well as a warm disk surrounding a young star can complicate matters on these diagrams. These sources will nevertheless appear redward of the zero-age main sequence (ZAMS) tracks in color-magnitude diagrams and can be used to constrain the properties of candidate young stars. Yet because the reddening to these sources is unknown, they cannot be dereddened and are not selected as candidate members as an initial selection criterion. After the removal of probable background contaminating AGN from the mid-IR data (see section 4.2.3), young stellar objects (YSOs) are identified using the following criteria: i) position on the near-IR color-color diagram (CCD) and on the combination of near- and mid-IR CCDs indicating an excess of IR emission; ii) positions on the mid-IR CCDs indicating an excess of mid-IR emission; iii) positions on color-magnitude diagrams (cmds) being

31 redward of the ZAMS. The criteria for YSO determination are as follows: iii must be satisfied in all cases; if i is satisfied then ii must also be; if ii is satisfied i does not necessarily have to be. In other words, all sources which are young and are at the correct distance (roughly 400 pc) will show up in the appropriate regions of color-magnitude diagrams (redward of main-sequence tracks). If a source has an excess of near-IR emission and it is young, it will also have an excess at mid-IR wavelengths, as disks with near-IR excesses are thought to be sufficiently close to the central star that they thermally emit at shorter wavelengths than the outer, cooler portions of the same disk (Class I/II sources; CTTSs). Both IR regimes will therefore emit strongly. However, if the disk has a substantial inner gap then there will be little near- IR emission, but it is still possible to have mid-IR emission from a warm outer disk. This is commonly observed in slightly more evolved young stars (Class III sources; wTTSs). The reasoning for this criteria is to further select against AGNs and interacting galaxies, as they occupy a similar region on the near-IR color-color diagram as do young stars with excesses (Geller et al. 2006). These criteria eliminate contaminants which may have been missed in the mid-IR AGN rejections. To summarize, all YSOs must satisfy i, ii, and iii, or ii and iii. The 3.6-8.0µm spectral indices, the spectral energy distributions, and the positions on CCDs are then used in combination with the [3.6]−[4.5] vs. [5.8]−[8.0] CCD to assess the evolutionary IR class of each bona fide YSO.

4.3 Color-Color Diagrams

4.3.1 Near-IR Color-Color Diagram The near-IR CCD has long been recognized as an important tool in the iden- tiﬁcation of young stars with disks. In 1988 Bessell and Brett demonstrated that the positions of stars on the J −H vs. H −KS diagram are located on characteristic loci that depend on the spectral types and evolutionary phases of the stars. For example, the locus of main sequence dwarf stars on this diagram sit on the solid track in Figure 4.3, with the positions of various spectral types indicated. This is a dereddened track and shows where main sequence dwarf stars would be located if there were no interstellar dust to scatter the starlight. Because it is almost never the case (except for close sources) that there is no reddening, stars will not sit exactly on this locus.

32 Figure 4.3: The near-IR color-color diagram is used to distinguish reddened stars from stars with IR excesses. The empirically deﬁned loci of dwarf stars is plotted as a solid track, giants as a dashed track (Bessel & Brett 1988), and classical T Tauri stars as a dot-dashed line (Meyer, Calvet, & Hillenbrand 1997). Dotted lines parallel to the reddening vector represent reddening envelopes. A reddening vector of AV = 2 mag is shown as an arrow.

33 Instead they will be displaced by some amount, in the direction of the reddening vector, depending on the dust-column density to the star. In Figure 4.3 a reddening vector of 2 magnitudes in the V band is shown (derived from Tokunaga 2001). So a star embedded in a molecular cloud will be displaced in the direction of the reddening vector to a greater degree than one at the same distance which is not embedded. There is a degeneracy in this locus such that a slightly reddened late-type star located redward of the mid-M position in both colors is indistinguishable from a highly reddened early-type star. If the amount of reddening to a star is known then its colors can be dereddened and its spectral type can be determined. Similarly, if the spectral type of a star on this diagram is known, then an estimate of the reddening to that star is attainable. Similar to the dwarf locus, the locus of giant stars is indicated by a dashed track (Bessel & Brett 1988) and the locus of classical T Tauri stars is indicated by the dot-dashed line (Meyer, Calvet, & Hillenbrand 1997). The dotted lines show the reddening envelopes of these tracks and are parallel to the reddening vector. Sources outside of these reddening envelopes have “abnormal” colors. Sources located the redward of the reddening line in the H − KS color that correspond to the late-M part of the dwarf locus have infrared excesses and are considered candidate YSOs from the near-IR data. The positions of these sources cannot be explained by photospheric emission alone and are interpreted either as AGN or star-forming galaxy contaminants (Geller et al. 2006) or as young stars with warm disks or disks plus envelopes. In this study I consider both the high-confidence candidate young stars (with H − KS colors ≥ 1σ to the right of the right-most reddening envelope) and low-confidence candidate young stars (with H − KS colors < 1σ to the right of the right-most reddening envelope). Sources with these characteristics constitute the initial sample of candidate YSOs, although the positions of many of these sources are likely caused by photometric scatter. All 134 near-IR sources from the IRSF images are plotted on the J −H vs. H −KS diagram in Figure 4.4. The positions of sources to the left of the left- most reddening envelope are caused by photometric scatter with relatively large photometric errors; these sources do not have normal colors. Error bars for near-IR excess sources are displayed, while typical errors for all sources are represented in the lower right corner. High confidence candidate YSOs are plotted as solid pentagons while low-confidence candidate YSOs are plotted as crosses. All other sources are plotted as open squares and may be either

34 Figure 4.4: All 134 IRSF sources are plotted on the near-IR color-color diagram. See Figure 4.3 for track descriptions. Error bars are plotted for low-confidence near-IR excess sources (< 1σ excess in H − KS; crosses) and high-confidence near- IR excesses (≥ 1σ excess in H − KS; filled pentagons).

35 field stars, background galaxies, classical T Tauri stars with only a slight IR excess, or weak T Tauri stars with little or no IR excess. There are 18 candidate young stars in Figure 4.4; 6 of these are high- confidence sources and 12 are low-confidence. This constitutes the initial candidate YSO selection. The near-IR color magnitude diagrams, the Spitzer mid-IR and IRSF plus Spitzer color-color plots, and the IR spectral energy distributions are then used to determine which candidate young stars are bona fide YSOs. If these sources are indeed young then they will lie to the right of the main sequence in color-magnitude diagrams. They will also exhibit a near-IR excess in IRAC photometry. However, there may be sources which were detected with Spitzer but were not seen to have an excess with IRSF. More evolved disks are thought to lack the warmer inner disk that would otherwise emit in the near-IR, but the mid-IR emission from the outer disk would still be present. Only 1 of the 6 high-confidence near-IR sources was detected by 2MASS (Src 6 from Table 2, or B33-1). The source to the far right of the main- sequence tracks (Src 16 with J −H = 0.468 and H −KS = 1.247) lies on the Herbig AeBe Locus (Lada & Adams 1992) and sits near the Northern “nose” of the Horsehead.

4.3.2 Mid-IR Color-Color Diagrams The mid-IR color-color diagrams are used to confirm the youth of candidate YSOs, to include stars which are young and that have strong mid-IR disk emission but weak near-IR disk emission, and to assess the evolutionary classes of probable young stars. These diagrams were previously used to reject background galaxies (see section 4.1). They are particularly useful in distinguishing between highly reddened stars and those with true IR excesses. Various regions on the [3.6]−[4.5] vs. [5.8]−[8.0] diagram have been identified as being the loci of different young evolutionary classes. Allen et al. (2004) model the positions of Class I sources (with accreting envelopes and disks) and Class II sources (with accreting disks) for various accretion rates, as well as for blackbody sources. These regions were later tested through observations of various star forming regions and were shown to be in good agreement with models (Megeath et al. 2004; Hartmann et al. 2005; Megeath et al. 2005; Luhman et al. 2006; Hernández et al. 2007; Balog et al. 2007; Winston et al. 2007). Hartmann et al. (2005) and Hernándezet al. (2007) note that young stars with anemic disks or no disks at all (wTTSs/Class

36 III sources) have similar near- and mid-IR colors to older dwarf stars, where photospheric emission dominates over any slight disk emission for young stars in this color range. All 22 sources detected in the four IRAC channels (with errors ≤ 0.1 in [3.6] and [4.5] and errors ≤ 0.2 in [5.8] and [8.0]) after extragalactic rejection criteria were applied are plotted in the [3.6]−[4.5] vs. [5.8]−[8.0] diagram in Figure 4.5. The loci for older stars or young stars with little excess emission (open squares) and Class II sources (filled triangles) are shown as a dot- dash box and a dotted box, respectively, as inferred from Allen et al. 2004, Hartmann et al. 2005, and Luhman et al. 2006. Class I sources (filled circles) are located redward of the Class II source locus for both [3.6]−[4.5] and [5.8]−[8.0] colors. Possible Class III sources with slight IR excesses are labeled as filled squares, although older YSOs with less apparent excess may have “normal” photospheric colors in these wavelengths. Figure 4.5 also shows the [3.6]−[4.5] vs. [4.5]−[5.8] color-color diagram with labeled sources from the [3.6]−[4.5] vs. [5.8]−[8.0] diagram. Two high-confidence near-IR excess sources were detected in both diagrams (the two Class I sources). No low-confidence near-IR excess sources were detected in either the 5.8 or 8.0µm channels. See Table 2 for source numbering. Out of the 18 near-IR excess sources (6 high-confidence and 12 low- confidence), 11 were detected by Spitzer in at least one channel (all 6 high- confidence sources and 5 low-confidence sources). We would expect the low- confidence near-IR excess non-detected sources to be seen in these images if they were indeed young and were located at this distance, even down to the brown-dwarf limit (M'0.08M ). So the 7 low-confidence sources that were not detected are considered extragalactic or possibly faint background M-stars. Two sources are identified as probable Class I IR sources from these diagrams. They correspond to B33-1 (Src 6 from Table 2) and Src 9, two embedded sources in the western bright limb of the Horsehead. Each have a high-confidence near-IR excess in Figure 4.4. B33-1 had previously been identified as a young star (Reipurth & Bouchet 1984), but its IR class had not been ascertained. The three probable Class II sources are Src 24, B33-25 (Src 31 ), and B33-21 (Src 27 ). All three were detected in the IRSF images but none were identified as having an IR excess. This is somewhat unexpected, but it should be noted that the classical T Tauri star (CTTS) locus extends well within the “normal” photospheric emission region between the giant reddening envelope

37 Figure 4.5: Mid-IR color-color diagrams are useful tools to differentiate between reddened stars and young stars with thermal disk emission. It is also useful for the preliminary determination of YSO classes for suspected young stars. Top: YSOs of differing evolutionary states and hence IR classes fall into well-determined regions on the [3.6]−[4.5] vs. [5.8]−[8.0] CCD. The dot-dashed region represents the locus of normal photospheres with little or no IR excess (open squares) while the dotted region represents the locus of Class II sources (filled triangles), as is inferred from Allen et al. (2004), Megeath et al. (2004), Hartmann et al. (2005), Megeath et al. (2005), and Luhman et al. (2006). The region between photospheric colors and Class II colors often contains sources with slight IR excesses and may represent Class III or transition disk objects (filled squares). Sources redward of the Class II locus in both colors is the location of Class I sources (filled circles). Bottom: The [3.6]−[4.5] vs. [4.5]−[5.8] CCD has less distinct regions but still represents IR class trends where sources redward of the origin in both colors tend to have earlier evolutionary states. Source labeling is the same as those in the upper CCD.

38 and the late-M type reddening envelope. However, if they are indeed Class II sources, an estimate of their reddenings can be obtained from the near-IR CCD, assuming they sit on the CTTS locus (dot-dash line in Figure 4.4). Src 24 is displaced from the CTTS locus by AV ' 7. B33-25 and B33-21 are both within 1.5σ of the CTTS locus, and so would probably have very little reddening if they are indeed Class II sources. All three of these sources don’t appear to be associated with the Horsehead pillar. They are likely associated with the cloud L1630 from which the Horsehead protrudes from or, in the case of Src 24, might be a background galaxy based on its high reddening. The two candidate Class III sources are Src 5 and B33-22 (Src 28 ). Neither of these sources appeared to have a near-IR excess in Figure 4.4. If both these sources are only slightly reddened then they would fall within the range of typical photometric scatter in the [3.6]−[4.5] vs. [5.8]−[8.0] diagram in Figure 4.5 (Winston et al. 2007) (the reddening vector is nearly vertical in this plot). The classification of these two sources as Class III is therefore more tentative than for the earlier classes identified above. The 38 sources detected in the 3.6, 4.5, and 5.8µm bands (with errors ≤ 0.1 in the first two bands and ≤ 0.2 in the third band) after extragalactic criteria were applied are also plotted in Figure 4.5. It is more difficult to assess the locus of IR evolutionary classes in the [3.6]−[4.5] vs. [4.5]−[5.8] diagram, as there is less photometric scatter (Hernándezet al. 2007) and more blending of classes. There are nevertheless reddening trends among the various evolutionary classes with younger evolutionary classes being redder in both colors. The adjusted positions of some of these sources from the 4-channel CCD to the 3-channel CCD is noteworthy. The Class I sources still have a strong excess of emission in the latter diagram. Two Class II sources also have an excess, although not as strong as in the 4-channel CCD. One Class II source and the two candidate Class III sources have lost most of their excess in the 3-channel CCD and have normal photospheric colors. The combination of near- and mid-IR color color diagrams will aid in discerning the true evolutionary class of these sources.

4.3.3 IRSF and IRAC Combined Color-Color Diagrams The combination of near- and mid-IR photometry in color-color diagrams has been shown to be useful in determining bona ﬁde IR excess sources (Hartmann et al. 2005; Winston et al. 2007). The 8µm band is particularly sensitive to PAH emission, which has relevant emission features at 3.3,

39 6.2, 7.7, and 8.6µm (Wu et al. 2005). These PAH emission features are related to star formation rates in galaxies, so color color diagrams without the 8µm channel will be least affected by background extragalactic contaminants which weren’t rejected from previous selection criteria. The shorter wavelength IRAC channels are also more sensitive and will detect bona fide YSOs from the low- and high- confidence near-IR excess sources, as any excess in the near-IR should also be apparent in the near- and mid-IR combined CCDs.

H −[3.6] vs. [3.6]−[4.5] and KS−[3.6] vs. [3.6]−[4.5] There were 90 sources detected in the 3.6 and 4.5µm bands with errors ≤ 0.1. However, 15 of these sources were not detected by IRSF nor by 2MASS and so are removed from the sample. The remaining 75 sources are plotted in the H −[3.6] vs. [3.6]−[4.5] and KS−[3.6] vs. [3.6]−[4.5] color-color diagrams in Figure 4.6. Symbols follow the classes defined in the [3.6]−[4.5] vs. [5.8]−[8.0] diagram (Figure 4.5; Class I colors: filled circles, Class II colors: filled triangles, Class III colors: filled squares). Hartmann et al. (2005) de- fine regions of IR excess in these diagrams using known young stars in the Taurus star-forming region. In their work there are clean breaks between wTTSs (Class III sources) and CTTSs (Class II sources). These breaks are represented in Figure 4.6 as dotted lines. In this sample there is a contin- uous blend between regions, therefore sources with [3.6]−[4.5] colors > 1σ to the right of the dotted lines are considered to have IR excesses. Sources that satisfy this criterion are depicted as crosses in the upper and lower left diagrams. Their error bars are plotted in the upper and lower right plots for reference. The positions of the two Class I sources are consistent with this age class from Hartmann et al. (2005). One previously identified Class II source (Src 24 ) does not show an excess in either diagram and instead sits alone in an abnormal region ([3.6]−[4.5] = 0.272, H −[3.6] = 2.83, KS−[3.6] = 1.26) with respect to the Taurus sources in Hartmann et al. (2005). Another Class II source (Src 27 ) lies to the right of the dotted line in both plots. The third source with mid-IR Class II colors (Src 31 ) does not appear to have an excess at these shorter wavelengths and may therefore represent a more evolved young star. The possible Class III sources also have interesting positions in these plots. Src 28 appears to have an IR excess in both diagrams, something that would not be expected of Class III source but rather of a Class II source.

40 Figure 4.6: The combination of near- and mid-IR are useful in distinguishing among sources with an excess of IR emission from those with normal or reddened photospheric emission. Sources with [3.6]−[4.5] colors > 1σ redward of the dotted line in the H−[3.6] vs. [3.6]−[4.5] (top) and KS−[3.6] vs. [3.6]−[4.5] (bottom) diagrams have IR excesses (Hartmann et al. 2005) and are added to the sample of candidate YSOs in the region (represented as crosses in the top and bottom left plots). 1σ error bars are plotted in the upper and lower right diagrams. Filled symbols are the same as in Figure 4.5.

41 However its positions on the mid-IR diagrams, especially the [3.6]−[4.5] vs. [4.5]−[5.8] diagram (Figure 4.5) are indicative of a Class III or transition disk source. There are 11 sources with positions ≥ 1σ redward of the dotted line in the [3.6]−[4.5] color and which were either not detected or did not have an excess in the longer-wavelength mid-IR diagrams. Four sources (Src 2, 4, 14, and 15 ) have an excess in both diagrams. Src 4 and Src 15 were determined to have a high-conﬁdence near-IR excess (Figure 4.4). Src 8 has an excess in the KS−[3.6]vs.[3.6]−[4.5] diagram and has a high-conﬁdence near-IR excess.

J −H vs. H −[4.5] and H −KS vs. KS−[4.5] The combination of the near- and mid-IR color color diagrams as J −H vs. H −[4.5] and H −KS vs. KS−[4.5] have been shown to be useful in identifying the IR excesses of young stars with circumstellar disks (Gutermuth et al. 2004; Winston et al. 2007). Any sources which show an excess in the near-IR color color diagram should also display an excess in these combined color- color diagrams if they are indeed young stars with disks. All 80 sources detected in the 4.5µm band with errors ≤ 0.1 and which were also detected by IRSF are plotted in Figure 4.7. Symbols are the same as for Figure 4.5, with IR excess sources (≥ 1σ redward of the dotted lines in H −[4.5]) represented as crosses (top and bottom left diagrams). The dotted lines mark the separation of sources with and without IR excesses, and also represent the slope of the reddening vector (Winston et al. 2007). In the top and bottom right diagrams the 1σ errors for IR excess sources are plotted. One source (Src 32 )) has an excess in both diagrams. Src 4 and Src 15 have an excess in the J −H vs. H −[4.5] plot. Both of these have a high- conﬁdence near-IR excess and were also determined to have an excess in both the H −[3.6] vs. [3.6]−[4.5] and KS−[3.6] vs. [3.6]−[4.5] plots.

4.3.4 Summary of Color-Color Diagrams A summary of all IR excess detections is provided in Table 4.1. Source coordinates, running numbers, and other identiﬁers are listed along with color-color plots for each source that has an apparent IR excess. Sources with detections in at least three color-color plots are listed as young stars (“Y”), while those with excesses in two diagrams have colors that are either inconsistent between diagrams (e.g., Source 8 has a near- and mid-IR excess,

42 Figure 4.7: Labeling is similar to that in Figure 4.6. Sources with H−[4.5] colors (top) and KS−[4.5] colors (bottom) ≥ 1σ redward of the dotted lines have an IR excess (crosses; Winston et al. 2007). Filled symbols are from Figure 4.5. The upper and lower right panels show errors for IR excess sources in these CCDs.

43 –2–

Table 2. Summary of Color-Color Diagrams for Candidate Young Stars

αJ2000.0 δJ2000.0 Source Plots with YSO (h m s) (◦ ’ ”) Number 2MASS J B33-a IR Excesses Colors?

05 40 47.45 −02 26 57.9 1 ······ 1(l) N 05 40 48.20 −02 25 37.7 2 05404820−0225377 ··· 4,5 U 05 40 50.20 −02 26 55.9 3 ······ 1(l) N 05 40 51.20 −02 29 46.1 4 ······ 1(h), 2, 4, 5 Y 05 40 51.58 −02 24 54.8 5 05405158−0224548 ··· 6 N 05 40 51.72 −02 26 48.9 6 05405172−0226489 1 1(h), 2, 3, 4, 5, 6 Y 05 40 51.78 −02 28 54.2 7 ······ 1(l) N 05 40 52.04 −02 25 56.5 8 ······ 1(h),5 U 05 40 52.39 −02 27 12.7 9 05405239−0227127 ··· 1(h), 2, 3, 4, 5 6 Y 05 40 53.04 −02 26 07.3 10 ······ 1(l) N 05 40 53.24 −02 23 41.6 11 ······ 1(l) N 05 40 53.27 −02 26 03.5 12 ······ 1(l) N 05 40 55.42 −02 23 33.3 13 05405542−0223333 ··· 5 N 05 40 56.40 −02 23 43.9 14 05405640−0223439 ··· 4,5 U 05 40 56.77 −02 26 53.2 15 ······ 1(h), 2, 4, 5 Y 05 40 56.80 −02 25 06.2 16 ······ 1(h) N 05 40 56.92 −02 23 05.3 17 ······ 1(l) N 05 40 57.11 −02 24 33.0 18 ······ 5 N 05 40 58.68 −02 25 26.7 19 05405868−0225267 10 3 N 05 41 01.14 −02 28 12.0 20 05410114−0228120 13 5 N 05 41 02.10 −02 23 44.1 21 05410210−0223441 ··· 5 N 05 41 03.18 −02 23 33.9 22 ······ 1(l) N 05 41 04.62 −02 24 06.4 23 05410462−0224064 ··· 1(l) N 05 41 05.52 −02 27 17.8 24 05410552−0227178 ··· 6 N 05 41 08.45 −02 26 20.9 25 ······ 1(l) N 05 41 08.87 −02 25 25.3 26 ······ 1(l) N 05 41 10.32 −02 29 03.9 27 05411032−0229039 21 4, 5, 6 Y 05 41 12.22 −02 27 34.4 28 05411222−0227344 22 4, 5, 6 Y 05 41 13.05 −02 29 12.2 29 05411305−0229122 23 5 N 05 41 13.68 −02 27 20.4 30 05411368−0227204 24 5 N 05 41 14.06 −02 25 05.1 31 05411406−0225051 25 6 N 05 41 14.15 −02 29 02.5 32 04411415−0229025 ··· 2,3 U 05 41 14.52 −02 29 16.4 33 ······ 1(l) N

Note. — The IR Excess plot numbering system is as follows: 1l)low-conﬁdence excess in J −H vs. H −KS, 1h)high-conﬁdence excess in J −H vs. H −KS, 2)J −H vs. H −[4.5], 3)H −KS vs KS−[4.5], 4)H −[3.6] vs. [3.6]−[4.5], 5)KS −[3.6] vs. [3.6]−[4.5], 6)[3.6]−[4.5] vs. [5.8]−[8.0]. Sources labeled as “Y” have IR excesses in at least three color-color diagrams. Those with “U” have excesses in two plots and have either inconsistent color excesses for typical YSOs or have excesses in only the longer- wavelength plots. Sources labeled with “N” have IR excesses in only one color-color diagram and are not considered young stars. 2MASS coordinates are quoted for those sources with 2MASS detections in the Point Source Catalogue. aReipurth 1984

44 but no excess in those plots which combine the two wavelbands) or have an excess in the longer wavelength color-color diagrams but not in the near- IR. These latter sources may correspond to more evolved young stars with anemic circumstellar disks. Those with an excess detection in only one plot are rejected as YSO candidates and are probably due to photometric scatter or contamination from extragalactic objects. Color-magnitude diagrams and spectral energy distributions will aid in the determination of bona fide YSOs from the sources with uncertain classifications. Source 6 and Source 9 have IR excesses in all six color-color plots. They are considered Class I sources from their locations on the [3.6]−[4.5] vs [5.8]−[8.0] plot. Source 4 and Source 15 have excesses in four of the five color-color diagrams they were plotted on (neither was detected in the 8.0µm IRAC channel). Two sources (Source 27 and Source 28 ) have excesses in the three longer-wavelength plots. None of the low-confidence near-IR excess sources appeared to have an excess in any other plots, while five of the six high-confidence near-IR excess sources also had an excess in at least one other color-color diagram.

4.4 Color-Magnitude Diagrams

4.4.1 Near-IR Color-Magnitude Diagrams Color magnitude diagrams are common tools used in observational astronomy to estimate a star’s mass and age based on its magnitude and color from at least two ﬁlters. Conversions from color to temperature and magnitude to luminosity allow for the placement of stellar sources on a Hertzsprung- Russell diagram. Several color-magnitude diagrams based on the theoretical modeling of pre-main sequence evolution exist in the literature, however for low-mass young stars the tracks by Baraﬀe et al. (1998) have been shown to be the most accurate (White et al. 1999). Therefore for the near-IR color- magnitude diagrams I have used the evolutionary tracks from Barrafe et al. (1998), hereafter referenced as BCAH98. The BCAH98 models needed some slight adjustments in order to correctly overplot the near-IR data. A range of iso-age and iso-mass tracks are given for masses between 0.02M –1.4M and for ages between 1 Myrs–10 Gyrs. However, the zero-age main sequence iso-age track occurs at about 500 Myrs, so I’ve truncated the iso-age tracks to this limit. For each track is a list of

45 apparent magnitudes in the optical and near-IR. The near-IR magnitudes are given in the CIT system, which must be transformed to the 2MASS system. Then the tracks must be shifted to the appropriate distance, or, equivalently, the IRSF data must be shifted to a distance of 10 pc. I have adopted the former BCAH98 shift to the correct distance. Transformations between the 2MASS photometric system and a variety of other systems are provided in Carpenter (2001). The transformations between the 2MASS and CIT system are reproduced below:

(KS)2MASS = KCIT + (0.000 ± 0.005)(J − K)CIT + (−0.024 ± 0.003)

(J − H)2MASS = (1.076 ± 0.010)(J − H)CIT + (−0.043 ± 0.006)

(J − KS)2MASS = (1.056 ± 0.006)(J − K)CIT + (−0.013 ± 0.005)

(H − KS)2MASS = (1.026 ± 0.020)(H − K)CIT + (0.028 ± 0.005) These transformations are applied to the BCAH98 CIT near-IR colors. To ﬁnd the J2MASS magnitude, the (KS)2MASS and (J −KS)2MASS equations are added to produce the following:

J2MASS = 1.056JCIT − 0.56KCIT − 0.037

A star’s color is independent of its distance (neglecting reddening) and so only the magnitude in these color-magnitude diagrams must be corrected for distance. The distance modulus (µ) is simply the diﬀerence between a source’s apparent magnitude (which I will denote m0) and its absolute magnitude (M ; the apparent magnitude at d = 10 pc). It is deﬁned in the following way:

m0 − M = 5 log(d) − 5,

where d is the distance to the source (in parsecs). Therefore, the in- trinsic magnitude to a source is just the distance modulus plus the absolute magnitude:

m0 = 5 log(d) − 5 + M.

The complete transformed and shifted J -band BCAH98 magnitude is below:

46 J02MASS = 1.056JCIT − 0.56KCIT − 0.037 + 5 log(d) − 5

Two near-IR color-magnitude diagrams are plotted for all IRSF sources. As with the color-color diagrams, certain color-magnitude diagrams (CMDs) happen to be more useful than others. The J vs J − H and the KS vs. H − KS diagrams are commonly used in the literature (e.g. Wang et al. 2007). The J vs. J − KS and J vs. H − KS ones are also used, albeit more sparingly. Most near-IR CMDs have a relatively narrow color range in the evolutionary tracks and so it is often difficult to accurately determine the mass or age of stars from these diagrams alone. The J vs. J −KS and KS vs. H − KS CMDs were chosen for their relatively wide theoretical color spreads for this analysis (Figs. 4.8 and 4.9). Reddening also needs to be taken into account as it affects a source’s magnitude and color in well-determined ways. Reddening vectors and iso-age/mass track labels have been included in both diagrams. The reddening vector is computed relative to reddening in some wave- band, usually the V -filter bandpass. Tokunaga gives the following ratios relating reddening in the V band to that in the near-infrared in Allen’s As- trophysical Quantities (2001):

A V = 5.82 E(J − K)

A V = 1.13 E(V − K) A V = 15.3 E(H − K)

Here AX is the extinction by dust (in magnitudes) in band X, and E(X − Y ) = AX − AY . Note that the optical parameter RV = AV /E(B − V ) = 3.1 has been used, where RV = 3.1 is a typical value for the interstellar medium (ISM) (Tokunaga 2001; Mathis 2001). From the above relations it is simply a matter of algebra to derive reddening vectors for a chosen value of AV . The J vs. J − KS color-magnitude diagram is plotted in Figure 4.8. The BCAH98 tracks are shifted to a distance of 400 pc (µ = 8.01) which is the distance to the dark cloud L1630 and the Horsehead (Anthony-Twarog 1982). Sources with color excesses indicative of youth (labeled “Y” in Table

47 4.2) are plotted as filled circles, while sources with an uncertain youth status (labeled “U” in Table 4.2) are plotted as filled squares. The same symbols are used in the KS vs. H − KS color-magnitude diagram in Figure 4.9. The dashed lines demarcate iso-mass loci across the isochrones. Five of the six sources with YSO colors lie well to the right of the 1 Myr isochrone track in both CMDs. It is difficult to get a quantitative estimate of the mass of these probable young stars even if the reddening to the sources was known, as an IR excess is not accounted for in these plots. However a qualitative estimate is possible by simply assuming youth, then dereddening in the opposite direction of the reddening vector. Two sources (Src 4 and Src 15 ) appear to lie in the very-low mass territory if they are indeed young, while one source (Src 8 ) has an uncertain YSO status but sits at a similar location. The hydrogen-burning (brown dwarf) limit is at roughly M = 0.08M , or J ' 14 and KS ' 13 for sources that are 1 Myrs old. These qualitative mass estimates should be analyzed with some warranted skepticism though. Tokunaga (2001) outlines the IR properties of galaxies and includes the expected number of galaxies per square degree per magnitude: dN = 4000 × 10α(K−17). dK Here α = 0.67 for 10 < K < 17. For K = 17, 4000 galaxies per square degree are expected, corresponding to roughly 67 galaxy detections at this magnitude in the IRSF 7.8’×7.8’ fov. At K = 16, roughly 14 would be expected and for K = 15, 3 are expected. For K ≤ 14, less than one galaxy detection is expected. The IRSF KS-band limiting magnitude is roughly 17.5. Although the high expectation at K = 17 will be dramatically reduced by the obscuring gas and dust from this star-forming region, galaxy contamination may still be a problem and the assumption that these three candidate YSOs with high-KS (KS ' 16) magnitudes aren’t galaxies may be invalid. The spectral energy distributions of these sources will aid in the final discrimination between the two options. The other three of the five sources (Src 6, Src 28, and Src 9 ) have consistent de-reddened (to the 1 Myr iso-age track) masses of > 1.4M , 0.4M , and 0.08M , respectively. These values have large and unknown errors caused by an unknown amount of excess in the KS band and are quoted merely to provide a qualitative physical understanding of these sources. Source 2 and Source 14 are both located consistently blueward of the ZAMS and are probable background sources. The rejection of these two

48 Figure 4.8: Sources with YSO colors (labeled “Y” in Table 2) and sources with uncertain status (labeled “U”) are plotted as filled circles and filled squares, respectively. Sources to the left of the ZAMS (at 500 Myrs) are rejected as field stars. Tracks are from Barrafe et al. (1998). They are transformed from the CIT system to the 2MASS system and are shifted to a distance of 400 pc, the distance to the Horsehead Nebula. Sources are not dereddened.

49 Figure 4.9: Similar to Figure 4.8

50 sources is consistent with the rejection of likely non-members of the σOri cluster in Hern´andezet al. (2007). Source 27 appears in the middle of the BCAH98 tracks in both CMDs; it is therefore not rejected as a candidate YSO because near-IR color-magnitude theoretical tracks are notorious for inaccuracies in exact interpretations of stellar masses and ages. The ﬁnal source with uncertain status, Source 32, is located redward of the 1 Myr iso- age track in Figure 4.8, but near the ZAMS in Figure 4.9. Its status therefore remains uncertain.

4.4.2 Mid-IR Color-Magnitude Diagram Spitzer mid-IR CMDs have been shown to be useful in the further assessment and confirmation of youth in suspected young stars (Winston et al. 2007; Allen et al. 2007). In Figure 4.10 the [3.6] vs. [3.6]−[4.5] CMD is plotted for all 75 Horsehead sources with errors ≤ 0.1 in the 3.6 and 4.5µm channels. Symbols are the same as in Figure 4.8 (sources with YSO colors: filled circles, sources with inconsistent YSO colors or uncertain status: filled squares, all other sources: open squares). The dashed line marks the empirical limit of AGN contamination from Winston et al. (2007) in the Serpens star-forming region. Most of the faint sources in their sample with [3.6] >15 mag are AGN, which often have colors similar to YSOs but have intrinsically fainter magnitudes. Four sources with YSO colors (open circles) and one with uncertain status lie in the expected region for young stars in Figure 4.10 (Sources 2, 6, 9, 27, and 28 ). The two others with YSO colors (Source 4 and Source 15 ) have [3.6] magnitudes greater than 1σ above 15, as do two sources with uncertain status Source 8 and Source 14. The remaining source with uncertain status (Source 32 ) was only detected in the 4.5µm band.

4.4.3 Summary of Color-Magnitude Diagrams Table 3 summarizes the positions of sources on the near- and mid-IR CMDs. Source numbers (from Table 4.2) are listed along with the CMDs for which each source is in the expected region for young stars. Only sources with a “Y” (YSO colors) or “U” (uncertain) status from table 4.1 are listed. Sources with expected positions for young stars in all three CMDs retain their status from Table 4.2. Those which don’t satisfy this in all three CMDs are rejected

51 Figure 4.10: Mid-IR CMD for all 75 sources with errors ≤ 0.1 in the 3.6 and 4.5µm channels. Most sources with [3.6] > 15 mag are AGN (Winston et al. 2007). Open circles have YSO colors from CCDs (Table 2) and open squares have an uncertain status. Sources having YSO colors with [3.6] > 15 have been changed to an uncertain membership status in Table 3, while those previously having an uncertain status and located in the same AGN region of this CMD are now rejected.

52 – 2 –

if they were labeled as “U” in Table 4.2 and are now labeled “U” if they were previously labeled “Y.”

Table 3. Summary of CMDs for Candidate Young Stars

Source Number YSO Region of CMDs YSO?

2 3 N 4 1, 2 U 6 1, 2, 3 Y 8 1, 2 N 9 1, 2, 3 Y 14 · · · N 15 1, 2 U 27 1, 2, 3 Y 28 1, 2, 3 Y 32 1, 2, 3 U

Note. — CMDs are labeled in the following manner: 1)J vs. J−KS , 2)KS vs. H−KS, 3)[3.6] vs. [3.6]−[4.5]

4.5 Spectral Energy Distributions

Spectral energy distributions (SEDs) are useful tools for the characterization of young stars with disks. An SED is a plot of flux vs. wavelength typically for a point source. A normal star without a warm disk would appear ap- proximately as a blackbody in such a plot. With a warm disk, however, the SED will be roughly equal to the superposition of two blackbodies: one for the stellar photosphere and one for the thermal emission from warm dust in a circumstellar disk. The shape of the SED can reveal detailed information about the physical properties of such a disk, including disk inclination, temperature, flaring properties, inner disk holes or disk gaps, and, with spectral information, grain size properties (Allen et al. 2004; D’Alessio et al. 2005). Moreover, the overall shape defines the IR class, and hence evolutionary stage, of the star plus disk/envelope (Lada 1987; Lada et al. 2006). Winston

53 et al. (2007) classify sources as Class I, Flat Spectrum, Class II, Transition Disk, and Class III based on criteria including positions on color-color diagrams, reddened and de-reddened SED slope, and the visual shape of the SED. Young stars are often classified based on the value of their mid-IR ([3.6]−[8.0]) spectral indices, defined as α = d log νFν/d log ν. In this study I use the following classification scheme: IRAC Class I sources have α > 0.3, IRAC“Flat” Spectrum Sources have −0.3 ≤ α < 0.3, IRAC Class II sources have −1.8 ≤ α < −0.3, sources with transition (“anemic”) disks have −2.56 ≤α < −1.8, while IRAC Class III sources have α < −2.56 (Hernández et al. 2007; Lada et al. 2006). The conversions between magnitudes and flux densities has been ac- complished through the absolute calibration of the 2MASS system (Cohen, Wheaton, and Megeath 2003) and Spitzer/IRAC (Reach et al. 2005). The magnitude/flux relation is just

m = c − 2.5 log Fν,

where m is the magnitude of a source, c is the zero-point magnitude, and Fν is the flux density. The constant c can be redefined as c = 2.5 log Fz, where Fz is called the zero-point flux. Then F m = 2.5 log z . Fν Rearranging this equation leads to the following, final relation.

− m Fν = Fz10 2.5 .

Zero-point flux values are from Cohen, Wheaton, and Megeath (2004) for 2MASS filters and from Reach et al. (2005) for IRAC channels (Table −23 −1 −2 −1 4). Values of Fz are given in Janskys, where 1Jy = 10 erg s cm Hz . This flux density times the frequency gives the flux at that frequency. The SEDs of YSOs identified in this study as well as for uncertain members in this work (Table 3) and previous work (Table 1) are plotted in Figures 4.11 and 4.12, respectively. Filled squares are data points for that source. For comparison I’ve overplotted the 1.25–8.0µm SED of IRAS 05384−0229, which shows no excess in any color-color diagrams and is therefore useful for a photospheric comparison to other sources (see Getman et al. 2007). The 1.65µm (H-band) IRAS 05384 flux is shifted to match that of each source as

54 – 1 –

Table 1. Absolute Flux Calibration

λ (µm) Fz (Jy)

1.25 1594 1.65 1024 2.15 666.8 3.6 280.9 4.5 179.7 5.8 115.0 8.0 64.13

Table 2. Absolute Flux Calibration

λ (µm) Fz (Jy)

1.25 1594 1.65 1024 2.15 666.8 3.6 280.9 4.5 179.7 5.8 115.0 8.0 64.13

Table 3. Absolute Flux Calibration

λ (µm) Fz (Jy)

1.25 1594 1.65 1024 2.15 666.8 3.6 280.9 4.5 179.7 5.8 115.0 8.0 64.13

Table 4. Absolute Flux Calibration

λ (µm) Fz (Jy)

1.25 1594 1.65 1024 2.15 666.8 3.6 280.9 4.5 179.7 5.8 115.0 8.0 64.13

this wavelength is not greatly affected by thermal disk emission. However, these SEDs are not dereddened and their true shapes, especially at the longer wavelengths which are more effected by nonlinear reddening, are difficult to ascertain. In this regard the overplot of the IRAS 05384 SED is merely a qualitative aid, but a shift to match the H-band fluxes was chosen over J-band because of the relatively high reddening that the latter wavelength experiences. Without knowledge of the spectral types of these sources it is difficult to quantitatively judge the reddening. One approach would be to assume these were all Class II sources, then to estimate the reddening to the CTTS locus in the J − H vs. H − KS diagram. When available the IRSF and IRAC photometry is combined with IRAS photometry quoted in Pound, Reipurth, and Bally (2003). In Figure 4.11 Sources 6 and 9 both show clear evidence of a strong disk or disk/envelope emission and are strong protostellar candidates. Source 27 is difficult to compare to IRAS 05384 as it is probably highly reddened and its H-band flux skews the overplot. Source 28 has a similar SED to the overplot and may have an 8.0µm excess, as the color-color diagrams indicated. In Figure 4.12 no sources show apparent strong emission in the 1.25–8.0µm range. However, IRAS 05384−0229, which was detected in the IRAS 100µm band, shows unusually high flux at that wavelength compared to its photospheric emission. However this source has a flag of “7” out of “9” for IR cirrus contamination, a value of “9” being the worst contamination. So contamination by IR cirrus probably accounts for the high flux at the 100 µm band. IRAS 05384 and ASCA C-20 had in previous work been suspected to

55 be young, but based on their placements on color-color diagrams and their SED shapes, they are rejected in this work as being young. Src 19 displayed an excess only in the H −KS vs. KS − [3.6] diagram, and was seen in X-rays by ASCA (Nakano et al. 1999). Because it has a spectral index of αIRAC = −2.52 it was not rejected and retains its “uncertain” classiﬁcation. If it has an inner disk it is optically thin and not substantial and represents a Class III source, as a spectral index of −2.66 is that of an M dwarf photosphere (Lada et al. 2006). IRAS 05386 was a previously suspected young star but was not detected by IRSF, 2MASS, or IRAC. Figure 4.13 shows the H-band inverted IRSF image with sources from Ta- ble 5 (photometry of bona ﬁde and uncertain YSOs) labeled. Circles are the YSOs in B33 while the uncertain members are indicated with perpendicular lines.

56 Figure 4.11: The spectral energy distributions of probable young stars (Table 3) in the 8’×8’ field of view of centered on the Horsehead. When available IRAS fluxes are supplemented to the IRSF J,H, and KS near-IR fluxes as well as the Spitzer/IRAC 4-channel mid-IR fluxes. Source 6 and Source 9 are flat-spectrum protostars, while Source 27 and Source 28 are probably transition disk/Class III field or embedded sources and are not likely to be directly associated with the Horsehead (see text). The overplotted SED is from IRAS 05384 (Figure 4.12), which did not have an excess in any CCDs. Overplots are shifted to match the H-band fluxes, as longer wavelengths are strongly affected by reddening and the shorter wavelengths are strongly affected by disk emission. 57 Figure 4.12: Similar to Figure 4.11, except for uncertain YSOs from this study and for previously suspected YSOs (Table 1 and 3).

58 – 2 – C 2 7 2 7 0 · · · A 0 6 5 9 3 · · · . . . 2 1 R · · · . . 2 1 2 I - - - 0 0 α 6 0 2 6 9 2 3 2 7 8 7 5 0 9 2 0 1 0 0 0 . . . . . ] 0 0 0 0 0 · · · 0 · · · . ± ± ± ± ± · · · 8 1 7 1 8 9 [ 8 5 7 1 3 9 1 1 9 9 8 6 7 7 2 . . . . . 1 0 5 8 9 1 1 8 3 6 1 3 0 6 1 6 5 3 3 0 5 0 0 0 0 0 0 . . . . . ] 0 0 0 0 0 · · · 8 · · · . ± ± ± ± ± · · · 5 4 0 1 1 2 [ 5 9 4 7 8 8 9 2 4 5 3 7 8 9 3 . . . . . 2 0 6 9 9 1 1 8 4 9 4 9 7 1 5 4 7 4 5 0 3 2 2 4 0 3 0 4 0 0 0 0 0 0 0 0 0 0 0 ...... ] 0 0 0 0 0 0 0 0 5 . ± ± ± ± ± ± ± ± 4 8 9 6 0 9 4 2 4 [ 7 3 8 1 8 2 6 4 2 4 3 8 9 5 8 8 3 5 5 9 7 4 8 5 4 ...... 3 4 5 4 0 2 0 7 9 1 1 1 1 1 1 B n i 0 5 6 3 2 4 6 1 7 4 5 5 s 2 1 4 1 2 0 0 0 0 0 0 0 0 0 O 0 0 ...... s ] 0 0 0 0 0 3 S 0 0 r · 6 3 e · . ± ± ± ± ± ± ± Y · B 3 2 1 5 3 9 b 4 8 [ 8 7 9 9 5 f 2 2 m n 8 0 4 5 9 i 6 8 e o 8 4 2 6 0 3 5 . . . . . s . . M 5 1 5 2 1 r y 8 9 1 1 1 1 1 a r n t i t S a e a t t g r 4 7 3 5 3 3 6 2 a m n e S 0 0 0 0 0 0 0 0 ...... d c u o 0 0 0 0 0 0 0 0 K n o t F ± ± ± ± ± ± ± ± Y U o − S 3 2 8 8 8 9 5 8 h 0 8 0 8 3 2 3 6 R ...... H I P 1 0 1 0 0 0 0 0 n i . n 5 8 5 7 5 4 6 6 o 5 0 0 0 0 0 0 0 0 i ...... t H 0 0 0 0 0 0 0 0 e a l r ± ± ± ± ± ± ± ± − b u 2 3 0 7 2 8 7 9 t J 5 9 1 1 4 5 9 5 a a ...... s 1 0 1 1 1 0 0 1 T d n a 4 6 4 6 4 2 5 4 b 0 0 0 0 0 0 0 0 ...... - 0 0 0 0 0 0 0 0 S ± ± ± ± ± ± ± ± J K 5 4 9 0 0 6 3 6 : 2 0 7 8 3 8 0 6 ...... S 2 8 5 7 2 3 7 3 S 1 1 1 1 1 1 1 1 A M r 2 e e c b m r 5 9 7 2 8 6 4 9 o u m 1 1 2 3 2 r o u f S N y r t e a 9 1 7 2 7 9 5 4 ...... m 8 6 2 3 6 3 2 4 o 0 ) . t 4 4 1 5 2 0 0 3 0 ” o 0 6 9 7 6 5 9 9 7 h ’ 0 2 2 2 2 2 2 2 2 p 2 ◦ J 2 2 2 2 2 2 2 2 ( d δ 0 0 0 0 0 0 0 0 n a − − − − − − − − s e t 2 0 9 7 8 2 5 2 a 7 2 3 7 6 3 1 2 0 n ) ...... i s 1 1 2 6 8 0 4 2 0 d 0 5 5 5 5 5 1 1 1 r 0 m o 2 0 0 0 0 0 1 1 1 o 4 4 4 4 4 4 4 4 J h ( C α 5 5 5 5 5 5 5 5 a 0 0 0 0 0 0 0 0

59 Figure 4.13: H-band IRSF image of B33 with YSOs (circles) and uncertain sources (perpendicular lines) labeled. Source 6 and Source 9 are ﬂat-spectrum protostars located at the western bright rim of the Horsehead.

60 Chapter 5

Discussion and Conclusion

In this paper I present the results of IRSF JHKS and Spitzer/IRAC imagery of the Horsehead Nebula in an effort to characterize the state and extent of star formation in the region. Candidate YSOs are chosen based on positions on color-color diagrams indicating an excess of IR emission, interpreted as thermal emission from a warm circumstellar disk or envelope. Color-magnitude diagrams were used to further select against field stars and faint extragalactic objects. Spectral energy distributions are created for four probable young stars, four candidate young stars, and two previously identified candidate young stars. Spectral indices are calculated for bona fide and candidate YSOs that were detected in the 3.6µm and 8.0µm bands.

5.1 The Horsehead as a Site of Triggered Star Formation

The two flat spectrum sources are located at the western bright limb of the Horsehead. A flat spectrum source represents an intermediate evolutionary phase between a Class I source (disk plus envelope) and a Class II source (optically thick disk). Source 9 is a previously unknown young star and is probably located behind the western ridge as it is not visible in optical images of the region. The evolutionary phases of these two sources are significant in the understanding of the formation and evolution of this pillar as well as for analogous pillars in other star-forming regions. Although the IR class identification is not a direct indicator of stellar age, it is a useful tool to understand the relative ages of young stars. That these two sources are

61 separated by a projected distance of only ∼26” (0.05 pc at a distance of 400 pc) and are in the same evolutionary phase indicates a common formation time, possibly from the same collapsing clump. Neither source is located in the denser sub-mm regions near this ridge identified by Johnstone, Matthews, and Mitchell (2006). Gritschneder et al. (2006) modeled the effect of high-energy radiation on turbulent gas and demonstrated that pillars the size of the B33 can form in timescales of roughly 3×105 years from the photoevaporation of molecular material by the UV flux of massive stars. White et al. (2007) review the evolutionary and absolute ages of protostars and, for the Taurus-Auriga star-forming region, infer Class I ages of ∼2×105 yr. This value is close to the estimated formation timescale of B33 (for a 0.7 pc pillar, τ ∼2×105 yr). Given the approximate evolutionary ages of these sources relative to the predicted formation timescale of the Horsehead, it is therefore likely that these stars were formed during or immediately after the formation of the pillar rather than being pre-existing young stars now being exposed at the western edge of this pillar. Furthermore, the RDI model of triggered star formation (as opposed to the “collect and collapse” model) accounts for several of the observational properties of this region (Chen, Lee, and Sanchawala 2006), including a dense sub-mm region at the ionizing edge of this pillar (Johnstone, Matthews, and Mitchell 2006; Ward-Thompson et al. 2006), possibly more evolved sources extending in the direction of the ionizing source (Hernández et al. 2007; Megeath et al. 2005), and newly-formed stars at the bright rim of the pillar which do not predate the existence of the pillar. Dobashi et al. (2001) studied the relationship between the most luminous protostar in a molecular cloud and the mass of that cloud. They find that there is a simple power law relationship between these parameters. More- over, there is a clear distinction in protostellar luminosity between protostars embedded in bright rimmed clouds near HII regions and those embedded in clouds far from HII regions. Those near HII regions have luminosities up to two orders of magnitude higher than those far from HII regions. Motoyama et al. (2007) model the collapse of stars from radiation driven implosion (RDI) caused by the UV flux of nearby massive young stars. They find that RDI can increase accretion rates in protostars by 1–2 orders of magnitude, and hence the luminosity by the same amount (for protostellar luminosities, L ∝ M,˙ where M˙ is the accretion rate). The luminosity of the most luminous protostar in a molecular cloud can therefore be used as a diagnostic tool for the formation mechanism of young stars in that region. The IR luminosities

62 of these protostars are calculated using the method described in Myers et al. (1987), where the flux density is integrated over all measured frequencies, with a linear interpolation between data points. I compute the IR luminosity of the most luminous protostar (Src 6 ) in the Horsehead using a similar method. But rather than a linear interpolation between flux densities, I find that a power law is a good fit to the data. −7 −1.2528 The best-fit power law is Fν(ν) = 3 × 10 ν between 1.25µm and 25µm. Integrating Fν(ν)dν between the J-band frequency and the IRAS 25µm frequency, I obtain a flux of 3.112×10−10 erg s−1 cm−2. At a distance of 400 pc, this corresponds to a luminosity of LSrc6 = 1.557L . This is the near- to mid-IR luminosity and is therefore a lower limit to the bolometric luminosity of this source. The mass of the Horsehead pillar is roughly 27M (Pound, Reipurth, & Bally 2003). Taking the logarithm of the cloud mass and the protostellar luminosity and comparing this data point to those in Figure 1 of Dobashi et al. (2001), the lower-limit luminosity of Src 6 sits near the boundary between protostars in BRCs and those in other clouds. Now Src 6 is a flat-spectrum protostar and is not accreting as highly as it would have been in the main part of it’s Class I evolutionary phase. This probable history of higher accretion combined with the lower luminosity for current accretion would likely have placed this source well within the other BRC luminosity data points. The position of sources on this figure is indicative of the source’s formation mechanism, with those sources located in BRCs near HII regions likely forming through RDI. The lower-limit in the luminosity of Src 6 for a cloud mass of 27M therefore bolsters the claim that it formed through RDI. Megeath & Wilson (1997) discuss the relevant size and time scales in- volved in the RDI and collect and collapse models. They note that RDI acts on smaller scales (∼1 pc and 0.5 Myr) than the collect and collapse model (∼10 pc and 1 Myr). At roughly 0.7 pc and 0.2 Myr protostellar ages, star formation in B33 is likely triggered by RDI. However, from these observations alone it is not possible to determine whether the photoevaporative pressure at this edge induced the collapse of a pre-existing dense core or whether this same pressure created the dense core that then gravitationally collapsed to form these young stars. These results add to the limited number of studies of RDI-triggered star formation in other pillars in star-forming regions.

63 5.2 Notes on Individual Sources

Sources 27 and 28 Sources 27 and 28 both satisfy criteria for being transition-disk sources based on their positions on color-color diagrams and color-magnitude diagrams as well as the values of their spectral indices. This evolutionary class corresponds to young stars with either optically thin disks or with inner-disk holes, transitioning from optically thick disk (Class II) sources to ones with little or no disk (Class III). Both are located near the base of the pillar. Neither source had previously been identiﬁed as being young. Source 27 sits on the mid-M main-sequence locus in the near-IR CCD (Figure 4.4). This means it is either a foreground M-type young star or an embedded/background F/G- type young star. Source 28 has an extinciton value of roughly AV ∼ 6 from the same ﬁgure. It is therefore probably either embedded within L1630 or behind the cloud itself and is also not associated with the Horsehead.

Source 4 The uncertain members also have interesting properties. Source 4, which had an excess in four color-color diagrams (including a high-conﬁdence near-IR excess), is located to the south-west of the western limb of the pillar. If it is indeed young it may represent a star forming from this pillar in the past. However, in the IRSF images it appears slightly elongated in the east- west direction. This elongation was conﬁrmed in ESO/VLT images of the Horsehead. Furthermore, this source was not detected in IRAC’s 5.8µm and 8.0µm channels, detections which would usually be expected for young stars in this region.

Source15 The location of Source 15 is signiﬁcant. It is near the location that Pound, Reipurth, and Bally (2003) predicted a young, possibly Class 0 object that may indeed have formed the jaw-region of the Horsehead. It is clear that this source is not in fact a Class 0 source (Class 0 sources are not detected in λ < 10µm). This source showed an excess of IR emission in four color-color diagrams, but was below the 3.6µm limit and thus in the probable AGN region of Figure 4.10. It is therefore likely to be either a very reddened Class

64 II source or a background AGN. It too was not detected in channel 3 and 4 of IRAC.

Source 19 Source 19 was detected in X-rays by ASCA (Nakano et al. 1999) and ROSAT (unpublished archival pointed observations) and displayed an excess in the H − KS vs. KS− [3.6] diagram. It has a spectral index of αIRAC = -2.52 which is just below the limit of -2.56 for photospheric emission. So If this source is indeed young then it is nearly at a Class III stage and has little or no inner disk. However, it is in an abnormal region to the left of the left-most reddening envelope on the near-IR color-color diagram.

Source 32 This source sits in the middle of an IR-bright cloud edge that is not detected in a 0.9m KPNO Hα image (Pound, Reipurth, and Bally 2003), a detection that would be expected if the ridge was being photoionized by σ Ori and were not obscured by intervening cloud material. This ridge may therefore be an optically-thick portion of L1630 and the IR images may simply be peering through material that was not observed in the Hα. Source 32 was detected in the JHKS images and in IRAC’s 4.5µ band. On the color-color plots combining these wavelengths this source exhibited an excess.

65 5.3 Future Work

There remains much work to be done regarding the formation and evolution of the Horsehead Nebula, including present and past star formation in the region. Spectra of the uncertain members from this work would greatly aid in the age classiﬁcation of these sources. Spectral types would also enable accurate estimates of the reddening to these sources, providing insight into their physical location in our line of sight with the Horsehead. Source 6 and Source 9 are relatively nearby examples of ﬂat-spectrum protostars, of which follow-up studies could undoubtedly be useful in uncovering the properties of this class of young stars. X-ray studies of the region would also be useful for studying the high-energy properties of protostars and for searching for more evolved young stars. Now that the star-forming properties of this interesting region have been revealed, the Horsehead is bound to get the attention it so deserves from researchers in the astronomical community.

Acknowledgments I’d like to thank the committee members who kindly agreed to scrutinize this thesis and my defense. Without their help this project would not have been possible. I would like to extend my warmest thanks to S. Thomas Megeath and Motohide Tamura for providing me with the data this thesis was based on. Their comments over the last year have been invaluable and I am grateful for the time they’ve taken to answer my questions. Rob Willson was kind enough to be an unoﬃcial committee member and to provide useful feedback throughout this entire project. I’d like to thank Kevin Luhman, Eric Feigelson, and Jon Swift for their helpful comments on my work and for the references they provided me with. This work makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the Na- tional Science Foundation. This research made use of the SIMBAD database, operated at CDS, Strasbourg, France. Funding came from the Massachusetts Space Grant Consortium and the Tufts Summer Scholars Program. Finally, I’d like to thank my family and friends for their help, support, and patience with me over the last year. Thanks for putting up with me!

66 Bibliography

[1] Abergel, A., et al. 1999, in The Universe as Seen by ISO, Eds. P. Cox & M.R. Kessler, ESA-SP 427, 615

[2] Abergel, A., et al. 2003, A&A, 410, 577

[3] Allen, P.R., Luhman, K.L., Myers, P.C., Megeath, S.T., Allen, L.E., Hartmann, L., and Fazio, G.G. 2007, ApJ, 655, 1095

[4] Allen et al. 2004, ApJS, 154, 363

[5] Anthony-Twarog, B.J. 1982, AJ, 87, 1213

[6] Bally, J., and Reipurth, B. 2003, AJ, 126, 893

[7] Balog et al. 2007, astro-ph/0701741

[8] Baraﬀe, I., Chabrier, G., Allard, F., and Hauschildt, P.H. 1998, A&A , 337, 403

[9] Barnard, E.E. 1913, ApJ, 38, 496

[10] Barnard, E.E. 1919, ApJ, 49, 1

[11] Bertoldi, F. 1989, ApJ, 346, 735

[12] Bessell, M.S., and Brett, J.M. 1988, PASP, 100, 1134

[13] Bok, B. 1977, PASP, 89, 597

[14] Bonnell, I.A., Larson, R.B., and Zinnecker, H. 2007, PPV, ed. B. Reipurth, D. Jewitt, & K. Keil (Tucson: The University of Arizona Press), 149

67 [15] Carpenter, J.M. 2001, AJ, 121, 2851

[16] Chen, W.P., Lee, H.T., and Sanchawala, K., IAUS, 237, 48

[17] Clemens, D.P., Yun, J.L., and Heyer, M.H. 1991, ApJS, 75, 877

[18] Clarke, C.J., Bonnell, I.A., and Hillenbrand, L.A. 2000, PPIV, ed. V. Mannings, A.P. Boss, S.S. Russell (Tucson: The University of Arizona Press), 151

[19] Cohen, M., Wheaton, Wm.A., and Megeath, S.T. 2003, AJ, 126, 1090

[20] Dale, J.E., Clark, P.C., and Bonnell, I.A. 2007, astro-ph/0701497

[21] D’Alessio, P., Mer´ın,B., Calvet, N., Hartmann, L., and Montesinos, B.

[22] Dobashi, K., Yonekura, Y., Matsumoto, T., Momose, M., Sato, F., Bernard, J.P., Ogawa, H. 2001, PASP, 53, 85

[23] Elmegreen, B.G., and Lada, C.J 1977, ApJ, 214, 725

[24] Fazio, G.G. et al. 2004a, ApJS, 154, 10

[25] Fazio, G.G. et al. 2004b, ApJS, 154, 39

[26] Fukuda, N., Hanawa, T, and Sugitani, K. 2002, ApJ, 568, L127

[27] Geller, M.J., Kenyon, S.J., Barton, E.J., Jarrett, T.H., and Kewley, L.J. 2006, AJ, 132, 2243

[28] Getman, K.V., Feigelson, E.D., Garmire, G., Broos, P., Wang, J. 2007, ApJ, 654, 316

[29] Gritschneder, M., Naad, T., Heitsch, F., and Burkert, A. 2006, IAUS, 237, 43

[30] Gutermuth, R.A., Megeath, S.T., Muzerolle, J., Allen, L.E., Pipher, J.L., Myers, P.C., and Fazio, G.G. 2004, ApJS, 154, 374

[31] Hily-Blant, P., Teyssier, D., Philipp, S., and G¨usten,A&A, 440, 909

[32] Kramer, C., Stutzki, J., and Winnewisser, G. 1996, A&A, 307, 915

[33] Hartmann, L., M et al. 2005, ApJ, 629, 881.

68 [34] Hern´andez et al. 2007, astro-ph/0701476v1

[35] Hoskin, M.A. 1963, William Herschel and the Construction of the Heav- ens, (New York: Norton & Company, Inc.), 98

[36] Hosokawa & Inutsuka 2006, ApJ, 646, 240

[37] Jeﬀries, R.D., Maxted, P.F.L., Oliveira, J.M., and Naylor, Tim 2006, MNRAS, 371, L6

[38] Lada, C.J. 1987, IAUS, 115, 1

[39] Lada, C.J. and Adams, F.C. 1992, ApJ, 393, 278

[40] Lada, E.A., Bally, J., Stark, A.A. 1991, ApJ, 250, 341

[41] Lada, C.J., and Lada, E.A. 2003, Annu.Rev.Astron.Astrophys., 41, 57

[42] Lee, H.-T. and Chen, W.P. 2007, ApJ, 657, 884

[43] Lee, H.-T., Chen, W.P., Zhang, Z.-W., and Hu, J.-H. 2005, ApJ, 624, 808

[44] Mathis, J.S. 2001, Circumstellar and Interstellar Material in Allen’s Astrophysical Quantities, 4th edition, ed. A.N. Cox (Springer), 523

[45] Megeath, S.T., Wilson, T.L. 1997, AJ, 114, 1106

[46] Megeath, S.T., et al. 2004, ApJS, 154, 367

[47] Megeath, S.T., Hartmann, L., Luhman, K.L., Fazio, G.G. 2005, ApJ, 634, L113

[48] Meyer, M.R., Calvet, N., and Hillenbrand, L.A. 1997, AJ, 114, 228

[49] Motoyama, K., Umemoto, T., and Shang, H. 2007, astro-ph/0702332

[Myers et al. 1987] Myers, P.C., et al. 1987, ApJ, 319, 340

[50] Nagashima, C., et al. 1999, Proc. Star Formation 1999, ed. T. Nakamoto (Nagano: Nobeyama Radio Obs.), 387

[51] Neckel, T., and Sarcander, M. 1985, A&A, 147, L1

69 [52] Oliveira, J.M., Jeﬀries, R.D., van Loon, J.T., and Rushton, M.T. 2006, MNRAS

[53] Ogura, K., and Sugitani, K. 1998, PASP, 15, 91

[54] Pound, M.W., Reipurth, B., and Bally, J. 2003, 125, 2108

[55] Reach, W.T. et al. 2004, ApJS, 154, 385

[56] Reach, W.T. et al. 2005, PASP, 117, 978

[57] Reipurth, B. 1983, A&A, 117, 183

[58] Reipurth, B., and Bouchet, P. 1984, A&A, 137, L1

[59] Romanishin, W. 2002, An Introduction to Astronomical Photometry Us- ing CCDs, http://observatory.ou.edu

[60] Sanchawala, K., et al. 2007, ApJ, 656, 462

[61] Skrutskie, M.F. et al. 2006, AJ, 131, 1163

[62] Sugitani, K., Fukui, Y., and Ogura, K. 1991, ApJS, 77, 59

[63] Sugitani, K., and Ogura, K. 1994, ApJS, 92, 163

[64] Tokunaga, A.T. 2001, Infrared Astronomy in Allen’s Astrophysical Quantities, 4th edition, ed. A.N. Cox (Springer), 143

[65] Urquhart, J.S., Thompson, M.A., Morgan, L.K., Pestalozzi, M.R., White, G.J., and Muna, D.N. 2007, astro-ph/0703252v1

[66] Wang, J., Townsley, L.K., Feigelson, E.D., Getman, K.V., Broos, P.S., Garmire, G.P., and Tsujimoto, M. 2007, ApJS, 168, 100

[67] Warren-Smith, R.F., Gledhill, T.M., and Scarrott, S.M. 1985, MNRAS, 215, 75

[68] White, R.J., Ghez, A.M., Reid, I.N., and Schultz, G. ApJ, 520, 811

[69] White, R.J., Greene, T.P., Doppmann, G.W., Covey, K.R., and Hillen- brand, L.A. 2007, PPV, ed Bo Reipurth, David Jewitt, and Klaus Keil (Tuscon: The University of Arizona Press)

70 [70] Whitworth, A.P., and Zinnecker, H. 2004, A&A, 427, 299

[71] Winston et al. 2007, in preparation

[72] Wu, H. et al. 2005, ApJ, 632, L79

[73] Zapatero Osorio, M.R., B´ejar,V.J.S., Mart´ın,E.L., Rebolo, R., Barrado y Navascu´es,D., Bailer-Jones, C.A.L., and Mundt, R. 2000, Science, 290, 103

[74] Zavagno, A., Deharveng, L., Comer´on,F., Brand, J., Massi, F., Caplan, J., and Russeil, D. 2006, A&A, 446, 171

[75] Zealey, W.J., Ninkov, Z., Rice, E., Hartley, M., and Tritton, S.B. 1983, ApJL, 23, 119

[76] Zhou, S. et al. 1993, ApJ, 419, 190

[77] Zissell, R.E. 1998, JAAVSO, 26, 151