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Wide-Field Survey of Emission-line Stars in IC 1396

We have made an extensive survey of emission-line stars in the IC 1396 HII region to investigate the low-mass population of pre-main sequence (PMS) stars. A total of 639 H-alpha emission-line stars were detected in an area of 4.2 deg^2 and their i'-photometry was measured. Their spatial distribution exhibits several aggregates near the elephant trunk globule (Rim A) and bright-rimmed clouds at the edge of the HII region (Rim B and SFO 37, 38, 39, 41), and near HD 206267, which is the main exciting star of the HII region. Based on the extinction estimated from the near-infrared (NIR) color-color diagram, we have selected pre-main sequence star candidates associated with IC 1396. The age and mass were derived from the extinction corrected color-magnitude diagram and theoretical pre-main sequence tracks. Most of our PMS candidates have ages of<3 Myr and masses of 0.2-0.6 Mo. Although it appears that only a few stars were formed in the last 1 Myr in the east region of the exciting star, the age difference among subregions in our surveyed area is not clear from the statistical test. Our results may suggest that massive stars were born after the continuous formation of low-mass stars for 10 Myr. The birth of the exciting star could be the late stage of slow but contiguous star formation in the natal molecular cloud. It may have triggered to form many low-mass stars at the dense inhomogeneity in and around the HII region by a radiation-driven implosion.


Introduction
As observational instruments have developed, systematic studies of stellar clusters have revealed many low-mass populations of pre-main sequence (PMS) stars in massive star-forming regions. It has been found that most stars are born in groups or clusters, not in isolated regions (e.g., Lada & Lada 2003;Allen et al. 2007). Stellar clusters provide a tool for investigating stellar evolution, because they are considered to constitute a natural sample of coeval stars at the same distance and with the same chemical composition. Young clusters appear to be a snapshot of stellar evolution over a wide range of masses. However, a careful study of the Orion Nebula Cluster (Hillenbrand 1997) has revealed that the younger stars ( < 0.3 Myr ) are clustered towards the center of the Trapezium Cluster (a projected distance of < 0.3 pc from θ 1 C Ori), while the older stars are widely distributed.
This indicates that star formation has recently occurred in the center of the cluster and that star formation with a low rate has taken place over a long time scale over the whole region. Muench et al. (2008) has suggested that there is no physical reason to separate the Orion Nebula Cluster and the Trapezium Cluster, and the entire region is a single contiguous star forming event. There are also examples of stellar clusters still producing massive stars. Cep A HW2 is the brightest radio source in Cep A (Hughes & Wouterloot 1984), an active star forming region within the molecular cloud associated with the cluster with an age of 5.5 Myr (Cep OB3b;Jordi et al. 1996). Jiménez-Serra et al. (2007) found a hot core internally heated by a massive protostar, and a circumstellar rotating disk around of the next generation of stars might be triggered.
The apparent age spread observed in the color-magnitude diagram of young stellar groups has been investigated by Burningham et al. (2005). They found that the spread could not be explained adequately by the combination of the binarity, photometric uncertainty and variability of the members. In the case of the σ Ori young stellar group, they suggested an age spread of 2-4 Myr. A possible explanation is that accretion from a disk causes scatter of the stellar continuum light (Tout et al. 1999). However, another possibility is that stars in a cluster are formed over a period of time. Three-dimensional numerical simulations including protostellar outflow feedback show that the cluster-forming clumps can maintain a quasi-virial equilibrium state for several periods of free-fall time (Nakamura & Li 2007), suggesting continuous star formation. Wang et al. (2010) also demonstrated outflow-regulated massive star formation. The episodic star formation causes an intrinsic spread in age and hence in the position in the color-magnitude diagram.
We have reported that young stars near the exciting star of the W5E H ii region are systematically older than those near the edge of the H ii region (Nakano et al. 2008). We further suggested that the birth of the most massive star occurs in the late stage of star formation in the cluster. Although the structure of the H ii region is simple to investigate, the distance of W5 is large, at 2.3 kpc, and therefore it is not easily accessible to the wide range of stellar masses.
IC 1396 is a ring-shaped giant H ii region in Cep OB2 with a size of about 3 • , located at the edge of the Cepheus bubble, which is a large, far-infrared shell with a diameter of ∼10 • (Kun et al. 2008). The structure of IC 1396 is rather simple and is powered by the multiple O star system HD 206267 (RA = 21 h 38 m 57. s 62, Dec = +57 • 29 ′ 20. ′′ 5), which is located near the geometrical center of IC 1396 in the young cluster Trumpler 37. The bright members of Trumpler 37 have been investigated by many authors (e.g., Garrison & Kormendy 1976;Marschall et al. 1990;de Zeeuw et al. 1999;Contreras et al. 2002). There are a number of bright-rimmed globules (Pottasch 1956) at its periphery. Its modest distance (< 1 kpc) and low foreground extinction (Av < 1 mag) make IC 1396 a good subject for our study of low-mass population associated with OB stars. The mean distance of 76 members of Cep OB2 derived from the Hipparcos parallax is 615 pc (de Zeeuw et al. 1999). Contreras et al.
(2002) made a color-magnitude diagram of the intermediate-mass members and fitted them to main-sequence stars at 870 pc. In this paper, we adopt 870 pc, which is frequently quoted by the recent works (e.g., Sicilia-Aguilar et al. 2004) .
Hα emission, a characteristic of T Tauri stars (TTS), is radiated from the solar-type magnetic activity in weak-lined T Tauri Stars (WTTS) or the chromospheric activity and magnetospheric accretion shock in classical T Tauri stars (CTTS) (Briceño et al. 2007).
Thus, a survey of Hα emission-line stars is a powerful method for searching for young low-mass populations (e.g., Nakano et al. 1995). Recently, NIR observations with array detectors and MIR observations from space (e.g. with the Spitzer Space Telescope) have advanced the protostar survey in the dense parts of giant molecular clouds. However, the Hα survey is still useful in areas with moderate or low extinction, and it is also a complementary method to search for PMS populations in a wide and rather dispersed region. Here, we report an extensive survey of emission-line stars in the IC 1396 H ii region and discuss the formation sequence of low and high mass populations in IC 1396.

Observations
Slitless grism spectroscopic searches for Hα emission objects in IC 1396 were conducted in 2004 and between 2006 and 2009 using the Wide Field Grism Spectrograph 2 (WFGS2; Uehara et al. 2004) at the f/10 Cassegrain focus of the University of Hawaii (UH) 2.2-m telescope on Mauna Kea. The log of the observations is given in Table 1. A 300-line mm −1 grism was used in conjunction with a wide Hα filter to isolate the first-order spectra between 6265-6765Å. A Tektronix 2048×2048 CCD was used as the detector, yielding a dispersion of 3.8Å pixel −1 , a pixel scale of 0. ′′ 34 and an instrumental field of view of 11. ′ 5 × 11. ′ 5 The whole IC 1396 region was covered by 157 field positions including overlapping areas. About 4.2 deg 2 were observed, both in the grism and the direct imaging mode of WFGS2. Figure 1 shows the boundary of the observed fields and indicates the names of bright-rims, dark clouds, and two O stars in the IC 1396 H ii region. For each field, we have taken a single 30-s exposure in the direct mode with a wide Hα filter. Then, we inserted a grism and obtained three dithered frames of 300-s exposures for slitless spectroscopy.
Direct images with an i ′ -band filter were also taken with WFGS2 in 2008 August and 2009 August. Three 30-s exposures were made for each field. The i ′ -band flux is less affected by excess emission of the circumstellar disk, Hα emission, and continuum veiling of CTTS.
Thus, i ′ -band photometry is useful as a measure of the underlying photosphere of CTTS (e.g., Cieza et al. 2005). 3-5 standards for the Sloan Digital Sky Survey (SDSS) system (Smith et al. 2002) were observed each night at nearly the same airmass (difference of < 0.1) as the targets. We used a dome flat for flat fielding.
The data were reduced using the IRAF software package. We applied the standard procedure for bias subtraction and flat fielding. For each field, three dithered grism frames are combined into one frame by referring the position of the Hα emission or absorption line of the star(s) in the field. We picked all stars in each wide Hα image using the DAOFIND package and extracted their one-dimensional spectra from the grism image. We visually investigated the spectra for the presence of the Hα emission line and found 639 Hα emission-line objects in the whole surveyed area. We used the SPLOT package to measure the equivalent width of the Hα line in the spectra. Their coordinates were determined by the astrometric calibration tool of the Starlink program GAIA using USNO-B catalog stars as the reference in the field. All of the photometric measurements were made with the APPHOT package for aperture photometry. The typical seeing of the image was 2.7-3.0 pixels. We used a 4-pixel diameter aperture for the aperture photometry. We then calibrated the photometric values to the SDSS system via standards. In general, the run was of high quality, but thin cloud affected a few fields. We checked the stability of the magnitude for three direct images in each field and examined the quality of the measurement. Excluding stars with saturated images or images with low photometric quality, we measured 586 stars in 13.6 < i ′ < 19.3.

Spatial Distribution
We list the 639 Hα stars in Table 2. The table shows the (1) identification number, (2, 3) J2000 coordinates, (4) i ′ -band magnitudes and (5) equivalent widths of Hα emission (Å). The Hα equivalent widths were measured for 608 stars. We compared our list of emission stars with the Two Micron All Sky Survey (2MASS) Point Source Catalog with 12 stars supplemented by the 2MASS 6x Point Source Catalog. We found matches for 617 stars with 2MASS stars within 2 ′′ and determined that 548 of these have good quality JHKs photometry (quality flag AAA). Due to small separation angles or positional uncertainties, two pairs of Hα stars are identified with single 2MASS sources (No. 250/251 with J21381703+5739265, and No. 600/601 with J21451578+5717392). We have included these stars in the above number. We list the 2MASS designation of the NIR counterpart, photometric quality, J magnitude, and J-H, H-Ks colors in columns (6)-(10) of Table   2. Av estimated from the NIR color-color diagram (see §4.1) is also listed in column (11). About 70% of our Hα stars cannot be identified with known sources in the SIMBAD database. The identification in the literature or comments are listed in column (12). Figure 2a shows the spatial distribution of emission-line objects while Figure 2b represents the contour map of the number density of the detected emission-line objects with the kernel method (Gomez et al. 1993). We adopted a Gaussian shape for the density distribution of the kernel with a smoothing parameter of h=2. ′ 4. The isolines are drawn at intervals of 150 stars per deg 2 , which is about three times the average number density in the field region (see Table 4), starting from 100 stars per deg 2 . It is clear from Figure   2 that the distribution of emission-line stars is far from uniform. To demonstrate the different properties in the different areas, we divided the surveyed field into the following subregions according to the number density of Hα stars. The primary peak of the most rich aggregate (the West region) is located between the main exciting star HD 206267 and the western elephant trunk globule Rim A (IC 1396A). The secondary peak is at the opposite side, to the east of HD 206267 (the East region), and the third peak is at the south of The West and the East regions appear to be surrounded by a halo-like structure (the Halo region). Some Hα stars appear to be associated with other small bright rimmed clouds such as Rim C and SFO 34 (Rim D). We defined the rest of the region as the field area (the Field region).
An objective prism survey using the Schmidt telescope Kun (1986) Table 2, our grism spectrum appears saturated but indicates a P Cyg-type Hα profile. Kun 320 (GL Cep) and Kun 337 (HD 239745) are also too bright and saturated in our grism spectra, and the positions of Nos. 77, 230, 12, and 266 are offset from Kun 53, 73, 170, and 322 by around 10 ′′ . Even if we include these seven stars, only 12% of stars can be identified with Kun stars. The rest of the Kun stars show obvious Hα absorption or no emission, and 10 of them show the spectra of M-type stars in our grism spectra. Medium resolution spectroscopy by Balázs et al. (1996) show that Kun 193 (No. 340;

Comparison with IPHAS Sources
In the recent paper by Barentsen et al. (2011) Barentsen et al. (2011). We also estimated the completeness of our survey with the same manner. For the PMS candidates including smaller EW ( > 10Å ), the fraction of recovered sources by Barentsen et al. (2011) falls off to 50%. On the other hand, the fraction of recovered sources with EW > 10Å maintains 86% (51 out of 59) in our sample. We confirmed that none of their rejected candidates ( 30 stars ) by Barentsen et al. (2011) were included in our list. Thus, our observations are more sensitive than their survey for the emission-line stars with small Hα equivalent width.

Identification with AKARI Sources
We identified 27 sources among our Hα stars with the sources of the point source catalog of the MIR All-Sky Survey obtained with the infrared camera (IRC) on board the AKARI satellite (Ishihara et al. 2010), with position differences of less than 5 ′′ . Table 3 gives the AKARI 9-µm (S9W ) and 18-µm (L18W ) Figure   4 shows the contour map of the number of emission-line stars overlaid on a false-color map using the IRC on AKARI. Blue is for 9 µm and red is for 18 µm. Probable members of 66 B-F stars from Contreras et al. (2002), supplemented by 8 OB stars from Hipparcos results (de Zeeuw et al. 1999), are also shown as filled blue circles. As for Hα stars in the West and East regions, high or intermediate-mass stars are concentrated near the center of the large dust cavity, which is presumably swept up by the expansion of the H ii region. We note that their distribution is slightly shifted to the east of HD 206267.

Probable Members
Our list of Hα stars in Table 2 may contain foreground or background, magnetically active field stars. We estimated the number of unrelated main sequence stars with enhanced chromospheric and coronal activity. Within the 4.2 deg 2 field towards the IC 1396, the Besancon model of the Galactic stellar populations (Robin et al. 2003) predicts ∼4000 and ∼ 1700 dM stars down to V = 20 mag in the foreground and the background (the distance range between 870 and 1500 pc) field, respectively. Although 44% of the foreground dM stars are in the range of V < 19 mag, 3% of the background dM stars are in the same magnitude range. Following the same approach as Dahm & Simon (2005), we used the the stellar luminosity function derived by Hipparcos (Jahreiß & Wielen 1997) and the dMe (with an Hα EW > 1Å) incidence from the nearby star survey of Hawley et al. (1996). We adopted a completeness limit for the present observations of V = 20 mag. We assumed the scale height of 325 pc for an exponential falloff in the space density, and 0.7 mag kpc −1 for the diffuse absorption. The estimated number of dMe stars in the foreground and the background (up to 1500 pc) is 193 and 33 in our observed area, respectively. The total number of the field dMe stars, 54 per square degree, is consistent with the average number density of Hα stars in the field area, 58 per square degree.
It is also possible that there is contamination by pre-main sequence stars associated with the foreground cloud (Patel et al. 1995;Weikard et al. 1996). As the AKARI 9 µm image, which traces the photodissociation region in the periphery of the H ii region ( Figure   4), is likely to be dominated by polycyclic aromatic hydrocarbons (PAHs) emission features, Hα stars associated with 9-µm emission are not likely to be foreground objects. The dark clouds B 163 and B 163SW do not show any sign of bright rims and their radial velocities of the molecular emission line are 6 km s −1 (Patel et al. 1995). The radial velocities of the other associated globules are between −9 km s −1 and +2 km s −1 . From the radial velocity, Leung et al. (1982) Bessell & Brett (1988), while the dash-dotted straight line indicates the locus of the unreddened CTTS following Meyer et al. (1997). We used the color transformation equations of Cutri et al. (2008) to convert their loci into the 2MASS system. The reddening vector adopted from Cohen et al. (1981) is also shown. We selected nearly 40% of the Hα stars whose NIR colors are consistent with the colors of the reddened CTTS in  Table 2 have Av > 1.0 mag, but 12 T Tauri stars have low extinction, and more than half of them are located near HD 206267. Figure 7 shows a histogram of the measured equivalent width. The fraction of weak (EW < 10Å) emission-line stars among the IC 1396 PMS candidates is low (27%) relative to that over the entire sample (44%). It appears that the fraction of CTTS candidates among the PMS candidates in the IC 1396 region is higher than that among the candidates proposed as foreground (or background) objects. We could not find any evidence of the difference between the spatial distribution of CTTS and WTTS.
To identify the HAeBe stars, which have larger infrared excess than CTTS, we adopted the extended criteria in the NIR color-color diagram proposed by Lee & Lim (2008). They   Figure 2b, disappears in Figure 8 and therefore could be a foreground object, other aggregates should be associated with IC 1396. Table 4 summarizes the number of emission-line stars in each region. The Hα stars, whose Av values we could not obtain, appear to be spread over the observed field rather uniformly. Thus, we suggest a large number of them are field stars.   (1989). Here, we adopted the pre-main sequence model of Siess et al. (2000). Their effective temperature and luminosity are converted into our system using the table of Kenyon & Hartmann (1995) and the equation of Jordi et al. (2006). In Figure 9, the 0.1, 1, 3, 5, 10, and 100 Myr isochrones are overlapped as well as the evolutionary tracks for masses from 0.1 to 2.0 M ⊙ , at an assumed distance of 870 pc. As the reddening vector runs almost parallel to the isochrones, the extinction value would not severely affect our age estimate.

Ages and Masses of Hα stars
The uncertainty of the distance modulus (8.94-9.70) and the choice of PMS models affects the estimation of ages and masses of Hα stars from Figure 9. We used the pre-main sequence models of Palla & Stahler (1999), which are converted to our photometric system by a table of Kenyon & Hartmann (1995) Figure 10. They indicate that the numbers of young stars ( < 1 Myr ) predicted from the Siess models (Fig 10a) are generally smaller than those predicted from the Palla & Stahler models (Fig 10b). On the other hand, the number of low mass stars ( < 0.4 M ⊙ ) derived from the Siess models (Fig 10c) are slightly large compared to those derived from the Palla & Stahler models (Fig 10d). We should notice the incompleteness of our PMS candidate sample along with possible effects on the distribution of masses and ages. Da Rio et al. (2010) carefully investigated the stellar population of the Orion Nebula Cluster by using these two PMS models. The completeness function in the mass-age plane, derived from their statistical simulation for two evolutionary models, shows that a large incompleteness correction is needed in the low mass (< 0.2 M ⊙ ) range.
They attributed the difference of results to the different shapes of tracks and isochrones in this mass range. Although we adopted the Siess models, the uncertainties would be larger for stars with ages < 1 Myr and masses < 0.2 M ⊙ . Furthermore, we could not detect emission-line stars near the region of higher obscuration such as the globules, and within 1 ′ of the bright O star HD 206267, due to saturation of the detectors, light contamination from this star, and diffraction spikes from this star. Table 5 shows the number of Hα stars in each age and mass bin in Figure 9. The age and mass histograms of Hα stars in each subregion are shown in Figure 11. We combined the SFO37/38/39/41/Rim B regions in the BRC (bright-rimmed cloud) region. Most of the Hα stars have ages of < 3 Myr and masses of 0.2-0.6 M ⊙ . The number of stars with an age less than 1 Myr and with an age of 1-3 Myr are comparable in the region associated with bright rims, i.e., the BRC region. We confirmed that the above result is also true for the age assigned from the Palla and Stahler model, and is not dependent on the choice of distance. The histogram in the West region is similar to that of the BRC region, probably because most of the Hα stars in the West region are associated with Rim A. ii regions ionized by O-B2 stars, many are formed to be surrounded by cold dust shells (good candidates for C&C processes), and many have bright-rimmed dust condensations protruding inside the H ii region (good candidates for RDI process). In the IC 1396 region, the presence of aggregates associated with the bright rims immediately inside the IC 1396 H ii region or the near side of the ionizing star HD 206267, such as the West (Rim A), Rim B, and SFO38 regions, suggest that the star formation in these regions has advanced away from HD 206267 in < 3 Myr. Such a configuration is consistent with squeezing of the pre-existing condensations compressed by the pressure of the ionized gas. The kinematical structure of SFO 37 (Sugitani et al. 1997;Ikeda et al. 2008) and an X-ray study of SFO 38 (Getman et al. 2007) also favored the RDI model.

Star Formation in the Central Region
In our previous results for the W5E H ii region (Nakano et al. 2008), we have found three aggregates of Hα emission-line stars in the W5E H ii region. Two of them are located at the edge of the H ii region, and one is near the central exciting O7 star. The former two aggregates are systematically younger than the latter.
Based on a comparison of the stellar ages in the BRC region with those in the East region of IC 1396 (Figure 11), it might appear that the low-mass population in the East region is older. We only find two stars younger than 1 Myr in the East region, in contrast to the West region. It may suggest that the star formation activity in the East region has ceased recently. As there remains no molecular material associated with the population in the East region, the exciting star HD 206267 has already dissipated the natal molecular cloud. However, a Kolmogorov-Smirnov (K-S) test gives a probability of > 15 % that the two samples have the same distribution in age, even if we choose either PMS model. The difference between the ages of Hα stars in the BRCs region and the East region is not statistically significant.

Sicilia-Aguilar et al. (2006a) estimated ages for six members near Rim A or in the
globule. They found that the ages of five stars are younger than 1 Myr, which is significantly lower than the average age of Trumpler 37. Figure 12 plots the spatial distribution of Hα stars with derived ages. At a glance, younger stars in Figure 12a are preferentially located at the West region compared to those at the East region. Again we used K-S test to compare the distribution in right ascension and declination derived from different choice of PMS models. A probability that two samples of different age (Fig 12a and 12b) have the same distribution in two coordinate axis is > 25 % and > 54 %. Thus, the age difference of young stars in the West and the East region is not statistically clear from the spatial distribution of our sample.
In addition to the West region, Figure 12b  Sicilia-Aguilar et al. (2005) suggested that the interstellar material, lying between powerful O6 stars and young stars with disks, shields the disk from UV radiation. However, there is no evidence of such an asymmetry of interstellar material from molecular line (Patel et al. 1995) and millimeter dust continuum observations (Rosolowsky et al. 2010). Actually, significant destruction of disks occurs for stars located out to distances of < 0.5 pc, which corresponds to only a few arcmin from HD 206267, as found in NGC 2244 (Balog et al. 2007). Sicilia-Aguilar et al. (2005) additionally suggested probable age differences for the explanation for the asymmetry in the CTTS distribution. Instead of disk photo-evaporation, we prefer that this spatial asymmetry represents the star formation sequence. If most of the young low-mass population in the West region was formed as a result of a triggered event, those in the East region could be formed along with massive stars.  Barentsen et al. (2011) reported the age gradient of the PMS candidates along the distance from this hot star. Their results also favors the triggered formation of stars in IC 1396. HD 206267 is known as the Trapezium system of OB stars with a maximum separation of 15000 AU, and the main component A is a triple (Stickland 1995). Burkholder et al. (1997) confirmed the observed masses and evolutionary model masses and suggested that the two bright components of the triple are O6.5 V((f)) (A 1 ) with 38-61 M ⊙ and O9.5: V (A 2 ) with 17-28 M ⊙ . Therefore, the age of HD 206267 is estimated to be < 3 Myr from the main sequence lifetime (Schaerer & de Koter 1997).

Star Formation in the BRC Region
Numerical simulation considering protostellar outflow-driven turbulence has been discussed by Li & Nakamura (2006) and Nakamura & Li (2007). Following their study, we suggested that the massive exciting star of the W5 E H ii region was possibly formed after the continuous production of stars (Nakano et al. 2008). It is reasonable to suppose that the formation of massive stars and the associated low-mass population in Trumpler 37, continued for a few Myr. Although it is difficult to separate the low-mass stars formed as the same generation with HD 206267 and formed by the trigger mechanism of HD 206267, we believe that many of H α stars at the rim were triggered to form by a radiation driven implosion mechanism after the formation of HD 206267.