STAR FORMATION ACTIVITY IN THE GALACTIC H ii COMPLEX S255–S257

The imaging observations of the S255 SFR centered on α₂₀₀₀ = 06^h12^m53^s, δ₂₀₀₀ = +17°59'21'' in the J (λ = 1.25 μm), H (λ = 1.63 μm), and K_s (λ = 2.14 μm) bands were obtained on 2000 October 13 with the 2.2 m UH telescope and SIRIUS, a three-color simultaneous camera (Nagashima et al. 1999; Nagayama et al. 2003). In this setup, each pixel of the 1024 × 1024 HgCdTe array corresponds to 028, yielding a field of view (FOV) of ∼49 × 49 on the sky. A sky field 10' south of the target position (centered on α₂₀₀₀ = 06^h12^m53^s, δ₂₀₀₀ = +18°09'26'') was also observed and used for sky subtraction. We obtained 18 dithered frames, 20 s each, thus giving a total integration time of 360 s in each band. The observing conditions were photometric and the average FWHM during the observing period was ∼07–09. Photometry was done using the point spread function algorithm of the DAOPHOT package (Stetson 1987) in IRAF,⁷ following the same procedure described in Samal et al. (2010). The photometric calibration was done by observing the standard star P9107 in the faint NIR standard star catalog of Persson et al. (1998) at air masses closest to the target observations. A comparison of present photometry with that of Two Micron All Sky Survey (2MASS) photometry as a function of J mag for the common sources is shown in Figure 2, where the 2MASS sources with photometric error ⩽0.1 mag and "read-flag" values between 1 and 3 are selected to ensure good quality (see Samal et al. 2010 for description). We find that the photometric scatter in magnitude increases at the fainter end, whereas the mean photometric uncertainties are ⩽0.1 mag in all the three bands. Absolute position calibration of the detected sources was achieved using the coordinates of 20 isolated bright 2MASS sources. The astrometric accuracy is estimated to be ∼005. To plot the data in color–color (CC) and color–magnitude (CM) diagrams, the JHK_s data were transformed to the California Institute of Technology (CIT) system using the relations given by Nagashima et al. (2003).

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2.2.1. Photometric Observations

2.2.2. Spectroscopic Observations

We have also analyzed the data of S255-IR at 14.94 GHz, taken from the National Radio Astronomy Observatory (NRAO) data archive (Project ID: AM697), obtained using the Very Large Array (VLA) in the D configuration. The observation was carried out in 2001 November 28. The flux calibrator was 0542+498 and 0625+146 was used as a phase calibrator. NRAO AIPS was used for the reduction of the data. The estimated uncertainty of the flux calibration was within 5%. The image of the field was formed by the Fourier inversion and cleaning algorithm task IMAGR, with a Briggs weighting function (robust factor = 0) halfway between uniform and natural weighting. We applied two iterations of phase self-calibration to improve the map. The resulting image has a beam size of ∼44 × 42 and the rms noise in the map is 0.11 mJy beam⁻¹.

We now discuss the nature of the associated low-mass stellar population using our NIR observations. To examine the nature of the stellar population in the gas ridge, we used NIR CM (H−K/K) and CC (H−K/J−H) diagrams of the sources detected in the HK_s and JHK_s bands, respectively, by comparing with the CM and CC diagrams of a nearby control field. Since the extinction is more sensitive to the shorter wavelength ranges and since young protostars are more prominent at longer wavelengths, we use the sources detected only in the H and K_s bands to construct H − K/K CM diagrams. Figure 6 shows the CM diagrams for the S255-IR SFR and for the nearby control field, where the nearly vertical solid lines represent ZAMS at a distance of 2.5 kpc, reddened by A_V = 0, 2.5, 20, 40, and 60 mag. The parallel slanting lines represent the reddening vectors to the corresponding spectral type. We would like to mention that the uncertainty in spectral type determination is mainly due to the error in the distance determination as well as the K_s-band excess of the individual sources. A comparison with the control field reveals that the majority of the cluster population shows an apparent separation from the foreground field stars distributed at (H−K) ∼ 0.2 mag. Figure 6 manifests that the majority of the sources in the S255-IR region are likely to be reddened MS members and/or pre-main-sequence (PMS) sources with strong infrared excess emission. Most of the sources with H−K ⩾ 2, not detected in the J band, are likely candidates for young sources with disks and envelopes. Massive stars (⩾8 M_☉) do not have a PMS phase and appear on the ZAMS while still embedded in the molecular cloud, whereas lower mass stars spend a significant amount of time in the PMS phase. Therefore, in Figure 6(a), not all of the sources may be reddened ZAMS stars. There are 22 sources in Figure 6 which appear to be massive with spectral type earlier than B3. Probably, these sources appear luminous due to excess emissions in the K_s band because one would expect radio free–free emissions in the vicinity of such sources. However, we detected only four compact radio emissions at 14.94 GHz in the vicinity of SFRs S255-IR and S255-N. The map is sensitive down to 0.11 mJy beam⁻¹ at 14.94 GHz, which corresponds to ∼2.9 × 10⁴⁶ Lyman continuum photons per second for an optically thin thermal source situated at a distance of 2.5 kpc, for an electron temperature of 10⁴ K. This corresponds to a ZAMS spectral type close to B2 (Panagia 1973) and therefore sets an upper limit of ∼10 M_☉ for the sources that do not show compact radio emissions. The sources associated with the radio emissions are marked in the H − K/K CM diagram (Figure 6), following the same nomenclature by Howard et al. (1997). Source IRS3 corresponds to the radio peak S255-2a, the strongest radio continuum source and the bluest of all the stars of the S255-IR cluster. The CM diagram indicates the star as a B2 type. The source IRS2 corresponds to the radio peak S255-2b and is situated ∼2'' southwest of S255-2a. The CM diagram suggests the source as a B2-B3 type. The infrared spectral-type estimation for both sources is consistent with the radio estimation at 14.94 GHz within a subclass. The position of star IRS1b coincides with the radio source S255-2c and the spectral type of the star appears to be earlier than O5, whereas the radio estimation suggests a spectral type of B1. Since the star shows large ∼ H−K color and is not detected in the J band, it is probably the youngest among the three radio sources. The polarization maps by Tamura et al. (1991) at the K band show the presence of a dusty disk toward this star. Therefore, we believe that the discrepancy of infrared and radio luminosities is caused by the presence of excess emission in the K_s band due to the circumstellar disk. We do not discard the possibility of a few reddened background field stars, which might have contaminated the CM diagram. However, by taking into account the distribution of the high column density of ¹³CO molecular gas (Chavarría et al. 2008), it seems unlikely that the stellar sources in the gas ridge consist of significant background contamination.

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To characterize the different populations of the region, we use the H − K/J − H CC diagram, as shown in Figure 7. The solid thin and thick dashed curves, taken from Bessell & Brett (1988), represent the loci of early MS and late giant stars, respectively. The left and middle parallel dashed lines are the reddening vectors for early MS and late giant stars (drawn from the base and tip of the two branches), respectively. The dotted line represents the locus of classical T Tauri stars (CTTS; Meyer et al. 1997) and the right parallel dashed line represents the reddening vector drawn from the tip of the locus of CTTS. The reddening vectors are plotted using A_J/A_V = 0.265, A_H/A_V = 0.155, and A_K/A_V = 0.090 for the CIT system (Cohen et al. 1981). The sources between the left and middle reddening vectors are considered to be either field stars or reddened MS stars or weak-line T Tauri stars, whereas sources with NIR-excess usually fall to the right of the middle reddening vector and are considered as young stellar objects (YSOs) with disk and envelope (Lada & Adams 1992). The NIR CC diagram can also be contaminated by foreground and background field stars and evolved stars. Therefore, to distinguish the cluster members and YSOs from these objects, we used a statistical approach by comparing the distribution of the sources in the NIR CC diagram of a nearby control region as shown in Figure 7(b). Comparison reveals that sources lying between the left and middle reddening vectors with (J − H) ⩾ 1.1 and the sources to the right of the middle reddening vectors are likely to be members and YSOs of the region. We do not rule out the possibility that a few massive and intermediate-mass members may fall below with (J − H) ⩽1.1, if they are reddened by the same amount like IRS3 and IRS2. However, considering that the initial mass function is dominated by the lower mass stars, most of the members should have (J − H) ⩾ 1.1 and are possibly reddened MS members or PMS stars with an inner circumstellar disk. The evolutionary stages of these sources are difficult to estimate until we have information about their SEDs. As can be seen in the CC diagram of the control field, the NIR-excess zone is free of field populations; therefore, 83 sources found in the NIR-excess zone in Figure 7(a) are likely to be YSO candidates. However, this number should be considered as a conservative lower limit, as there are also many sources detected only in the H and K_s bands and only in the K_s band, but not in the J band. The fraction of YSOs increases with high sensitivity observations at the L and M bands (e.g., Hoffmeister et al. 2008). Spitzer-IRAC observations by Chavarría et al. (2008) uncover several new clusters/groups as well as distributed population of YSOs in the direction of the S255–S257 complex. Giving emphasis to the distribution of IRAC classified Class II and Class I sources, they found that the ratio of the Class I/Class II sources is highest along the direction of the gas ridge between S255 and S257. In Figure 7(a), squares represent the distribution of 21 IRAC classified Class I and Class II YSOs and crosses denote the 115 IR-excess YSOs by Chavarría et al. (2008). It is important to note that many sources appear as normal stars in the NIR CC diagram, whereas they show IR-excess in the Spitzer bands. This gives an important clue that many sources that appear likely to be reddened MS stars in the NIR CC diagram could in reality be PMS sources with disks. In order to discuss the star formation activity in the gas ridge, we used the spatial distribution of the 83 YSOs identified from the NIR CC diagram and 26 sources having H−K > 2, along with the YSO candidates identified by Chavarría et al. (2008; see Section 4.4).

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The K_s-band luminosity function (KLF) of an embedded cluster is useful in constraining the age of the cluster (Megeath et al. 1996; Ojha et al. 2004b, 2004c). In order to derive the observed KLF, one needs to apply corrections for (1) the incompleteness of the star counts as a function of the K_s magnitude and (2) the field star contribution in the line of sight of the cluster. The completeness was determined through artificial star experiments (see Ojha et al. 2004c). To correct for the foreground and background contamination, we used the stellar population synthesis Galactic model (Robin et al. 2003) using a similar procedure described in Ojha et al. (2004c). The star counts were predicted in the direction of S255-IR using the average extinction to the embedded cluster A_V = 10 mag for a distance of 2.5 kpc. After determining the fraction of contaminating stars over the total model counts, we scaled the model prediction to the star counts in the reference field and subtracted the foreground and background contamination from the KLF of the S255-IR region.

After correcting for the foreground and background star contamination and photometric completeness, the resulting KLF for the S255-IR region is presented in Figure 8. In Figure 8, a power law with slope α ($d N(m_K)/dm_K \propto 10^{\alpha m_K}$, where N(m_K) is the number of stars brighter than m_K) has been fitted to the KLF using a linear least-squares fitting routine. The KLF of the S255-IR region shows a slope of α = 0.17 ± 0.03 over the magnitude range 12.5–17 in the K_s band (solid line in Figure 8). The derived KLF slope is lower than those generally reported for the young embedded clusters (α ∼ 0.4; e.g., Lada et al. 1991, 1993; Lada & Lada 2003), indicating a younger age of ∼1 Myr. However, a break in the power law can be noticed at K_s = 15.25 mag and the KLF seems to be flat in the magnitude range of 15.25–17.75. The slope of the KLF in the magnitude range of 12.5–15.25 (dashed line in Figure 8) comes out to be 0.25 ± 0.05, which is still lower than the average value of slopes (α ∼ 0.4) for young clusters of similar ages (Lada et al. 1991, 1993; Lada & Lada 2003). A turnoff in the KLF has also been observed in a few young clusters, e.g., at K ∼ 14.5 mag, K ∼ 16.0 mag, and K ∼ 15.75 mag in the case of Tr 14 (distance ∼ 2.5 kpc; Sanchawala et al. 2007), NGC 7538 (distance ∼ 2.8 kpc; Ojha et al. 2004c), and NGC 1624 (distance ∼ 6.0 kpc; Jose et al. 2011), respectively.

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As discussed in Section 4.1, the excess emission in the K_s band may make the YSOs luminous, hence the mass estimation on the basis of H − K/K CM diagram will lead to an overestimation of the mass. To minimize the effect of NIR excess emission on mass estimation, we preferred J-band luminosity rather than H or K band. Figure 9 represents J − H/J CM diagram for the sources detected in our observed area, where asterisks represent YSOs selected from the NIR H−K/J−H CC diagram. We assumed an age of 1 Myr to estimate the mass, which would be appropriate because the mean age of 16 YSOs in the direction of the embedded cluster is found to be ∼1.2 Myr (see Section 6). The solid and dashed curves in the figure denote the loci of 1 Myr PMS isochrones by Siess et al. (2000) for 1.2 M_☉ ⩽ M ⩽ 7 M_☉ and Baraffe et al. (1998) for 0.05 M_☉ ⩽ M ⩽ 1.4 M_☉, respectively. The dotted slanting lines are the reddening vectors for 4, 0.5, 0.2, and 0.1 M_☉ stars for 1 Myr isochrones. For an assumed age of 1 Myr, the lowest mass detected in our J-band image is about 0.1 M_☉ for A_V = 5 mag and is about 0.3 M_☉ for A_V = 11.5 mag. In Figure 9, squares represent the Spitzer-IRAC identified Class I and Class II YSOs, which indicates that the IRAC observations in the gas ridge are sensitive only down to 0.5 M_☉, whereas the IR-excess YSOs (crosses) identified using NIR plus IRAC bands are sensitive down to 0.2 M_☉ for A_V = 5 mag. The lower sensitivity of IRAC observations is most likely due to the combined effect of saturation, bright polycyclic aromatic hydrocarbon (PAH) emission toward S255-IR region, and lower resolution of the IRAC data. Indeed, Chavarría et al. (2008) found that the area with PAH emissions in the complex decreases the sensitivity to detect faint sources at 5.8 and 8.0 μm compared to the area free of PAH emissions. These are the possible causes that might have limited the sensitivity of the IRAC bands to detect YSOs up to 0.5 M_☉, which is less sensitive than our NIR detection (∼0.1 M_☉). The reddening vectors for the assumed age suggest that most of the YSOs have masses in the range of 0.1–4 M_☉. With the assumed age, the distribution of YSOs with wide colors (Figure 9) probably indicates the combined effect of variable extinction, weak contribution of excess emission in the J and H bands, and/or sources in different evolutionary stages (see also Wang et al. 2011). We note that the estimated stellar masses from the infrared CM diagrams rely on uncertain age and distance. The ambiguity is more severe among the early B-type stars because the inferred stellar mass for a 4 M_☉ star from the 1 Myr PMS isochrone can be as high as 18 M_☉ when estimated from the MS isochrones of a younger age. For this purpose, we have also drawn the reddening vectors (slanting dash-dotted lines) from the tip of the B0 and B5 ZAMS (dash-dotted vertical curve) locus. The concept of PMS is not applicable to stars above 8 M_☉ since the birth line and the ZAMS unify at those mass levels (Palla & Stahler 1993). Therefore, the ionizing sources IRS3 and IRS2 are likely to be reddened ZAMS massive stars. It is worth pointing out that Howard et al. (1997), based on the stellar fit of the observed NIR data points, indicated that the region is rich in B-type stars, containing 9 B0-type (∼18 M_☉) and 19 B5-type (∼6 M_☉) MS stars within an area of 1.69 arcmin² centered on the S255-IR cluster (α₂₀₀₀ = 06^h12^m53^s, δ₂₀₀₀ = +17°59'23''). The J − H/J CM diagram does not reveal such massive members (at least earlier than B0) for an assumed age of 1 Myr, as most of the stars have masses less than 4 M_☉. Moreover, one would expect radio emission from such sources. The circumstellar disks are rare among the intermediate mass to massive PMS stars with ages ∼1 Myr, therefore the mass uncertainty due to the circumstellar disk should not be the major cause for massive members. The discrepancy could be due to the contribution from the bright nebulosity to the photometric magnitudes.

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The infrared excess in the case of YSOs can be due to circumstellar disk/envelope or weaker contribution from the reflected stellar radiation of the dust emission. In any case, infrared excess represents the association of young sources in an SFR and possibly traces the star formation history of a cloud. Figure 10 shows the spatial distribution of YSOs from an NIR CC diagram (asterisks), possible YSOs detected only in the H and K_s bands with H−K ⩾ 2 (filled circles), and sources detected only in the K_s band (crosses). The outer boundaries of the 870 μm dust continuum emission (Klein et al. 2005) observed with a resolution of ∼26'' are shown with white dashed lines in the figure. This has the shape of a bar extending in the north–south direction and separates the two evolved H ii regions S255 and S257. The outer boundaries of S255 and S257 are also marked with the red thick dashed lines, which are extracted from the ionized gas traced by the radio continuum emission at 610 MHz observed with the Giant Metrewave Radio Telescope (GMRT) at a resolution of ∼55 × 65, taken from I. Zinchenko et al. (2011, in preparation). The yellow circles denote the positions of the three prominent 1.2 mm dust continuum clumps MM1, MM2, and MM3, revealed by Minier et al. (2005) at a resolution of 24''. The large crosses represent the positions of the radio continuum sources, three of which are found in the vicinity of the S255-IR cluster and one associated with the S255N SFR. Most of the NIR YSOs are distributed in close proximity to the S255-IR cluster and are also projected in the direction of the ionized emissions of the evolved H ii regions. The majority of the YSOs with H−K ⩾ 2 and those detected only in the K_s band are distributed along the bar of dust emission. The sources detected only in the K_s band show a scattered distribution along the gas ridge, with an enhanced concentration toward the western side of the S255-IR cluster and in the close vicinity of the molecular clump MM2. Toward S255N, MM1, and MM3, we do not see an enhanced concentration of YSOs as seen toward S255-IR; possibly, they are in a more embedded phase of star formation. In our observed NIR field, we identified 109 YSO candidates, which include 83 from the NIR CC diagram and 26 sources with H − K > 2. Out of 109 sources, we find that only 40 in our observed region are common to the sources classified as YSOs by Chavarría et al. (2008). Therefore, in the present work we identified 69 new YSO candidates in the region. To see the spatial distribution of all the YSO candidates, we overlaid the stellar surface number density (SSND) contour map on the K_s-band image in Figure 10 using the kernel method (Gomez et al. 1993; Silverman 1986). The SSND contour map was made using a total of 213 YSO candidates, which include the 69 new YSOs we detected and the 144 YSOs found by Chavarría et al. (2008) using NIR and IRAC bands in the region we observed. The stellar density distribution is smoothened by a 0.1 pc × 0.1 pc sampling box. The contour levels are drawn from 20 stars pc⁻² with an increment of 30 stars pc⁻². The clustering toward the S255-IR region is obvious in the SSND map along with sub-structures between S255-IR and MM1 and toward MM2. The main peak of the projected stellar density distribution was found to be α₂₀₀₀ = 06^h12^m53^s, δ₂₀₀₀ = 17°59'25'', with a central density of 180 stars pc⁻²; one can also see a second peak at α₂₀₀₀ = 06^h12^m54^s, δ₂₀₀₀ = +18°00'06'', with a central density of 87 stars pc⁻². This distribution confirms the presence of an enhanced concentration of YSOs in the gas ridge compared to that toward S255 and S257 and is therefore a relatively more active site of star formation.

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An embedded cluster along with three prominent clumps, MM1, MM2, and MM3 (Minier et al. 2005), appears to be sandwiched between the two evolved H ii regions S255 and S257. The clumps MM1 and MM2 contain a similar amount of dense gas of mass > 200 M_☉ and show a higher gas temperature of 40 K and are in agreement with those of Zinchenko et al. (2009). This is in contrast to MM3, which has a mass of ∼100 M_☉; the gas temperature is expected to be around 20 K. These observations led us to suggest that the temperatures of MM1 and MM2 are significantly higher than the temperatures of the starless clumps, showing that the clumps are possibly internally heated by the protostars, whereas the clump MM3 represents an earlier stage of star formation as compared to MM1 and MM2. It is to be noted that the clump MM2 lies close to the infrared cluster S255-IR, where the cluster harbors two optically visible massive stars, along with several intermediate- to low-mass stars of different properties detected in the JHK_s bands, whereas the sources in the vicinity of the clump MM2 are detected in the K_s band only and possibly represent a more embedded phase of star formation. The clump MM2 is also associated with CH₃OH, H₂O, and OH masers (Minier et al. 2005; Goddi et al. 2007). The CH₃OH maser traces the earliest stages of massive star formation (Walsh et al. 1998) and turns off shortly after the formation of the UC H ii region. The detection of maser emission toward the center of MM1, which has a characteristic lifetime of 2.5 × 10⁴ to 4.5 × 10⁴ yr (van der Walt 2005), indicates that this is not a dynamically evolved object. It seems that the sources associated with the clump MM2 are most probably younger than the sources associated with S255-IR. The sources associated with the H ii regions S255-2a and S255-2b of the cluster S255-IR are not associated with millimeter continuum emission and hence, presumably, are in a later stage of evolution. There was also evidence of bipolar outflows from the sources associated with S255-2c (Wang et al. 2011; I. Zinchenko et al. 2011, in preparation), along with the presence of a dusty disk (Tamura et al. 1991), which supports the interpretation that the B-type stars may form via an accretion outflow process similar to that of low-mass stars. Both these phenomena indicate an early evolutionary stage of the source (IRS1b) associated with S255-2c, which is supported by the dynamical age and the age derived from the SED fitting models (i.e., age ⩽0.1 Myr). An interesting question is whether the sources in the vicinity of MM2 and the massive stars of S255-IR have formed in the same region of space. Although presently we do not have the necessary data to address this issue and reach a firm conclusion, it is quite possible that both populations are spatially unrelated. We also do not discard the possibility that they may represent different evolutionary stages of the same cluster-forming environment. For example, in the case of the high-mass SFR, IRAS 20293+3952, Palau et al. (2007) found that the region consists of starless cores, millimeter sources, UC H ii region, and more evolved NIR sources. They concluded that these sources are not forming simultaneously in this cluster environment. Rather, there may be different generations of star formation and this could possibly be the case for S255-IR. In several cases, for example, W3 (Feigelson & Townsley 2008; Ojha et al. 2009) and IRAS 19343+2026 (Ojha et al. 2010), the massive embedded sources are found to be associated with a cluster of more evolved lower mass Class II and Class III sources, detectable in the NIR bands. This suggests that, in these regions, the intermediate to massive protostars have formed after the first generation of low-mass stars. The clump MM1 is associated with a UC H ii region identified in the 14.94 GHz map and is possibly ionized by a B1-type star. Kurtz et al. (2004) detected a Class I CH₃OH maser at 44 GHz and an H₂O maser at 22 GHz in the close vicinity of the UC H ii region. The non-detection of ionizing source(s) of the associated UC H ii region in the near- to mid-infrared band and the presence of H₂ knots tracing an outflow (Miralles et al. 1997) indicate a very young age. Since the initial phase of massive star formation is intimately linked to cluster formation and the fragmentation of the molecular clouds, this indicates that S255N is a site of cluster formation in its early stage. Indeed, Cyganowski et al. (2007) found evidence of a massive protocluster with the help of high-resolution (∼47 × 27) interferometric observations at 1.3 mm continuum emission and detected three compact cores with no infrared counterparts. The masses of the cores range from 6 to 35 M_☉. The identification of centimeter emission and multiple millimeter continuum sources toward the S255N region indicates an ongoing (proto)cluster formation. The features discussed above support the youthfulness of the MM2 and MM1 regions and indicate that the upper limit for star formation activity around MM2 and MM1 should not be more than ∼10⁵ yr. The clump MM3 shows no ionized gas emission in the proximity of the dust peak of mass ∼100 M_☉. For a typical star formation efficiency of 0.3 (cf. Lada & Lada 2003), the clump would likely form a star greater than 8 M_☉. If this happens, then it is a site for the earliest stages of star formation in which the massive star has not yet reached the ZAMS. This indicates that it is likely to be less than ∼10⁵ yr old because a single early B star (M ∼ 10 M_☉) reaches the ZAMS in about 10⁵ yr. It is worth mentioning that Wang et al. (2011) detected low velocity outflows in MM3 (S255S) as compared to MM2 and MM1 and suggested that MM3 is likely in the collapse phase. The lower temperature of the clump MM3 and non-detection of stellar sources in the K_s band possibly indicate that MM3 is in the earliest evolutionary stage. The above discussion reveals that the star formation is non-coeval along the gas ridge, where the stars of ages from 0.1 to 2 Myr are spatially distributed along the different locations of the cloud.

The exciting stars of S255 and S257 are isolated sources, without any cospatial clustering of low-mass stars around them (Zinnecker et al. 1993). However, Chavarría et al. (2008) hypothesize that a cluster (G192.63-0.00) is associated with S255. Future spectroscopic observations would confirm the association of G192.63-0.00 with S255. Figure 14 shows the surroundings of the evolved H ii regions S255 and S257, where the white background image represents the IRAC 8.0 μm emissions which generally trace PAHs coming from the photodissociation regions (PDRs) at the surface of the molecular cloud. Generally, these emissions are due to the absorption of far-UV radiation at the surface of the molecular cloud, which escape from the H ii region, and the PAHs within PDRs are excited and remit their energy at MIR wavelength. The green radio contours show the ionized gas from radio continuum observations at 610 MHz (I. Zinchenko et al. 2011, in preparation). In radio, both S255 and S257 appear almost circular in size. The figure also shows that 8 μm emission are prominent at the peripheries of S255 and S257, in the direction of the dark lane, and are weaker toward the interiors of S255 and S257 from the lane, particularly for S257. These emission features correspond to the dark patches seen in the optical band and also roughly coincide with CO emissions found in these directions by Chavarría et al. (2008). The ¹³CO map by Chavarría et al. (2008) reveals that the distribution of the gas is rather non-uniform and clumpy, with its surface density reaching maximum along the gas ridge, where new embedded star formation is expected to occur.

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S255–S257 is a part of the Gemini OB1 association. The star formation in OB associations is often quite complex. One common property of OB associations is that they are sub-structured and often consist of sub-groups of different ages. Carpenter et al. (1995a, 1995b) detected several lower luminosity sources and presumably low-mass cores scattered through the Gemini OB1 complex, whereas most of the massive dense cores are found adjacent to the optical H ii regions and often show arc-shaped structure. They interpreted that the induced star formation in the massive dense cores is prominent in the complex due to sweeping and compression of the molecular material. The complex S255–S257 shows a clear congregation of highly obscured YSOs, Herbig–Haro objects, astronomical masers, and compact and UC H ii regions. All these sources are concentrated preferentially along the vertical dense bar of gas and dust at the interface of the two H ii regions seen by the ionized gas distribution at 610 MHz. The ionized gas emissions are delineated with PDR, and therefore indicate a definite evidence of physical interaction of the H ii regions with the surrounding molecular material and are responsible for the shaping of the ring-like PDR structure. The ionizing sources of the two adjacent evolved H ii regions are MS O9–B0 stars. If they have evolved in the same ambient density, then they might have similar ages. The MS lifetime of an O9 star is of the order of ∼6.5 Myr. It is difficult to constrain the precise age of the MS star associated with S255–S257 with sufficient accuracy to compare its relative age with that of the age of YSOs detected in the gas ridge. However, the dynamical ages of both evolved H ii regions is in the range of ∼0.8–1.6 Myr (Chavarría et al. 2008; Bieging et al. 2009), whereas the low-mass sources scattered in the complex show a mean age of ∼2.5 Myr. Low-mass stars do not dramatically alter the environment in their vicinity, whereas the stellar radiations and winds from massive star(s) play a crucial role in the evolution of the nearby cloud and thus can induce the next generation of star formation (Elmegreen & Lada 1977). The minimum criterion to ascertain such processes is that the second generation star(s) should have younger age in comparison to nearby massive star(s). The physical properties of the compact and UC H ii regions in the gas ridge suggest that they are younger than adjacent H ii regions S255 and S257 (see Section 3). The age of the sources inside the dense ridge is of the order of ∼1.2 Myr or less and considering 1.6 Myr as a typical value for the age of evolved H ii regions, it appears that the sources in the gas ridge are younger than the adjacent massive stars as well as lower mass stars of the complex. These signatures indicate that the sources found in the gas bar might have formed due to the interaction of the evolved H ii regions.

A number of mechanisms by which massive stars can affect the subsequent star formation in the region have been proposed. One such process is the compression of pre-existing molecular clumps due to the impinging of an ionization/shock front, leading to density enhancement, which when it exceeds the local critical mass, collapse to form new stars (Bertoldi 1989; Lefloch & Lazareff 1994). The observational signature of such a process would be the limb-brighten rim, having a dense head and a tail, extending away from the ionizing sources (e.g., Miao et al. 2009). The morphology of the radio emissions and/or the Hα emission around the sources S255-2a and S255-2b show symmetrical distribution, not like the cometary morphology with a bright rim as found in the case of bright-rimmed clouds (e.g., Urquhart et al. 2006; Ogura & Sugitani 1998). The millimeter contours also do not show any asymmetric distribution of molecular matter, as expected in the case of cometary morphology (e.g., Morgan et al. 2004). The overall distribution of molecular material and the distribution of radio continuum do not give strong evidence in favor of star formation due to the compression of a pre-existing molecular clumps. However, the compression of pre-existing filamentary cloud cannot be ignored. Another proposed mechanism that can induce star formation around the H ii region is a "collect and collapse" process. In this mechanism, the neutral matter piles up between the ionization and the shock fronts of the expanding H ii region. This compressed material may be dynamically unstable and may fragment into dense clumps which can then form new stars (Deharveng et al. 2005). We observed recent star formation activity at the interface of S255 and S257. If the H ii regions evolved in the same ambient medium, and S255 is being excited by a relatively early-type source, then we expect to have a larger impact on the star formation process at the interface. Therefore, we compare the observed properties with that of the analytical model of star formation at the periphery of the expanding H ii region through the "collect and collapse" process by Whitworth et al. (1994). The model predicts the time when the fragmentation starts (t_frag) and the radius at the time of fragmentation (R_frag), which weakly depend upon isothermal sound speed in the compressed layer (a_s), and varies from 0.2 to 0.6 km s⁻¹, the number of Lyman continuum photons (Q_Ly) emitted by the massive star and the density of the ambient medium (n₀) in which the expansion of the H ii region occurs and are given by

$$\begin{equation} R_{{\rm frag}}\simeq 5.8 a_{.2}^{4/11} Q_{49}^{1/11} n_3^{-6/11} \,{\rm pc} \end{equation} \tag{ 2 }$$

$$\begin{equation} t_{{\rm frag}}\simeq 1.56 a_{.2}^{7/11} Q_{49}^{-1/11} n_3^{-5/11} \,{\rm Myr}, \end{equation} \tag{ 3 }$$

where a_.2 = [a_s/0.2 km s⁻¹], Q₄₉ = [Q_Ly/10⁴⁹ s⁻¹], and n₃ = [n₀/10³ cm⁻³] are dimensionless variables. Adopting 0.2 km s⁻¹ for a_s, 2.4 × 10⁴⁸ photons s⁻¹ for Q_Ly for the ionizing star (O9.5V) of S255 (Vacca et al. 1996), and 3 × 10³ for n₀ (Chavarría et al. 2008 and references therein), we obtained t_frag ≈ 1.1 Myr and R_frag ≈ 2.8 pc. This indicates that the fragmentation will start at a radius of 2.8 pc, whereas the young cluster and associated stellar sources are situated at a radius of ∼1.7 pc from the ionizing star of S255. If one considers the present age of the YSOs (∼1.2 Myr) plus the fragmentation time (∼1.1 Myr), one would expect an age difference of >2 Myr between the ionizing source of S255 (age ∼1.6 Myr) and the YSOs (age ∼1.2 Myr) in the core. At the same time, to reach from the fragmentation time of the shell to the present age of the YSOs (∼1.2 Myr), one would expect a shell radius much bigger than 2.8 pc, though the above estimations are based on the assumptions on the sound speed in the compressed layer. Increasing the value of the sound speed increases the fragmentation radius and time further. It appears that the H ii region is not old enough to produce the S255-IR cluster through the "collect and collapse" process. Morphologically, it also seems that the "collect and collapse" process is not viable here, as the cluster S255-IR is not associated with the ring of PAH emissions that surround the H ii regions S255 or S257, which generally represent the compressed neutral matter piled up between the ionization and shock fronts. Rather, the cluster is situated at the tip of the ring of PAH emissions. For example, in the case of a sample of H ii regions by Deharveng et al. (2005), where star formation at their peripheries is expected due to the "collect and collapse" process, one can see young sources/condensations, which are associated with the ring of PAH emissions that delineate H ii regions. Also in an ideal case, one would expect regularly spaced clumps along the periphery of H ii regions and/or young protostars of roughly similar ages, distributed preferentially along the swept-up material, which is not the case here, whereas we are seeing sources at different evolutionary stages across the gas ridge and dense clumps. The radio emission from S255 and S257 at 610 MHz is rather spherical and does not show morphology similar to the Champagne flow or blister type H ii regions, where one would expect the evolution of the H ii region in a non-uniform medium. As a result, clump formation due to the "collect and collapse" process may not be uniformly distributed. This discussion indicates that the "collect and collapse" process is also not at work here.

In the gas ridge, which scenario should work better is difficult to conclude, but the existence of molecular clumps is not unique to the "collect and collapse" model. Moreover, in this process the clumps should be distributed in a ring-like structure in the immediate H ii regions, not like the vertical bar structure which we are witnessing here. The morphological similarity of the matter sandwiched between S255 and S257 resembles on a large scale the process of star formation that has been observed in the dense bar-like structure in the Large Magellanic Cloud at the interfaces of two expanding supergiant shells LMC4 and LMC5 (Cohen et al. 2003), where it is believed that the collisions between the dense swept-up neutral materials lead to star formation (Yamaguchi et al. 2001a, 2001b). A similar kind of star formation activity might have happened here. The above discussion reveals that it is difficult to constrain a particular scenario that is at work with the present analysis, though we favor the last scenario. However, to strengthen the notion that the radiation-driven implosion and "collect and collapse" scenarios are not at work, it would be very useful to obtain the radial velocity information and the age of the stellar members using spectroscopic observations to confirm their association and relative evolutionary status, in order to obtain a better picture of triggered star formation.