STAR-FORMATION ACTIVITY IN THE NEIGHBORHOOD OF W–R 1503-160L STAR IN THE MID-INFRARED BUBBLE N46

2.1.1. Optical Spectrum

2.1.2. Near-infrared Spectra

To study the SF process, we used multi-wavelength observations from various Galactic plane surveys (e.g., the Multi-Array Galactic Plane Imaging Survey (MAGPIS; λ = 20 cm; Helfand et al. 2006), the Galactic Ring Survey (GRS; λ = 2.7 mm; Jackson et al. 2006), the Bolocam Galactic Plane Survey (BGPS; λ = 1.1 mm; Ginsburg et al. 2013), the APEX Telescope Large Area Survey of the Galaxy (ATLASGAL; λ = 870 μm; Schuller et al. 2009), the Herschel Infrared Galactic Plane Survey (Hi-GAL; λ = 70, 160, 250, 350, 500 μm; Molinari et al. 2010), the MIPS Inner Galactic Plane Survey (MIPSGAL; λ = 24 μm; Carey et al. 2005), the Wide Field Infrared Survey Explorer (WISE; λ = 12 μm; Wright et al. 2010), the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE; λ = 3.6, 4.5, 5.8, 8 μm; Benjamin et al. 2003), the UKIRT NIR Galactic Plane Survey (GPS; λ = 1.25, 1.65, 2.2 μm; Lawrence et al. 2007), the Galactic Plane Infrared Polarization Survey (GPIPS; λ = 1.6 μm; Clemens et al. 2012), and 2MASS (λ = 1.25, 1.65, 2.2 μm)). In the following sections, we give a brief description of these archival data.

2.2.1. NIR Data

2.2.2. H-band Polarimetry

We obtained photometric images and magnitudes of point sources in our selected region from the Spitzer-GLIMPSE survey at 3.6–8.0 μm (resolution ∼2''). The data were retrieved from the GLIMPSE-I spring 2007 highly reliable catalog. The image at MIPSGAL 24 μm was collected. We also used MIPSGAL 24 μm photometry (from Gutermuth & Heyer 2015) in this work.

Herschel far-infrared and sub-millimeter (sub-mm) continuum maps at 70, 160, 250, 350, and 500 μm were obtained from the Herschel data archive. The beam sizes of these images are 58, 12'', 18'', 25'', and 37'' (Griffin et al. 2010; Poglitsch et al. 2010), respectively. The processed level2₋5 products were downloaded using the Herschel Interactive Processing Environment (HIPE; Ott 2010).

The sub-mm continuum map at 870 μm (beam size ∼192) was retrieved from the ATLASGAL. A Bolocam 1.1 mm image was also obtained from the BGPS. The beam size of the 1.1 mm map is ∼33''.

The ¹³CO (J = 1−0) line data were obtained from the GRS. The GRS data have a velocity resolution of 0.21 km s⁻¹, an angular resolution of 45'' with 22'' sampling, a main beam efficiency (η_mb) of ∼0.48, a velocity coverage of −5 to 135 km s⁻¹, and a typical rms sensitivity (1σ) of ≈0.13 K (Jackson et al. 2006).

Radio continuum map at 20 cm was extracted from the VLA MAGPIS. The map has a 62 × 54 beam and a pixel scale of 2'' pixel⁻¹.

Spectroscopic observations offer a unique opportunity to characterize the individual astronomical objects. In Figure 1, we present optical (5500–9200 Å) and NIR H- and K-band spectra of the W–R 1503-160L star. The optical and NIR spectra mainly show the broad emission lines of ionized helium and nitrogen, confirming it as a W–R star. In general, due to the strong stellar wind of a W–R star, the outer hydrogen envelopes are thrown away and thus strong lines of He and other heavy elements (N, O, C) are observed in the spectra of W–R stars. In the present work, we have utilized both optical and NIR spectra to determine the spectral type of this source. Smith et al. (1996) have demonstrated a classification scheme of W–R stars using the helium lines (He ii at 4686 and 5411 Å, and He i at 5875 Å) in the optical spectrum. Figer et al. (1997) have also shown that several characteristic lines in the NIR K-band spectrum can be used to determine the spectral subtype of a W–R star. Our optical spectrum does not cover the lines in the blue part, as used by Smith et al. (1996), hence, we primarily carried out the identification of the spectral type using our K-band spectrum. The absence of carbon lines and the presence of highly ionized nitrogen lines in all the spectra preliminarily confirm the source as a nitrogen-rich W–R star (e.g., WN subtype; see Table 1). From our observed K-band spectrum, the equivalent widths of the 2.112 μm (He i + N iii), 2.166 μm (H i + He i + He ii), and 2.189 (He ii) lines were calculated by fitting Gaussian profiles (see Table 1). Good agreement of the ratio of equivalent widths (2.189/2.112 ∼ 0.85 and 2.189/2.166 ∼ 0.72) corresponding to the WN7 star, as reported in Figure 1 of Figer et al. (1997), confirms the source as a WN7-type star. Furthermore, we have compared the ratio of equivalent widths of several optical lines with those values reported in Conti et al. (1990). The ratios of the equivalent widths of He i 6678 Å and He ii 8237 Å (ratio ∼1.27), and He i 6678 Å and blended He i 7065 Å + N iv 7715 Å (ratio ∼0.37) agree well with the values reported for several WN7 stars (e.g., WR 55, WR 120, WR 148, and WR 158) in Conti et al. (1990). Hence, from our observed optical and NIR spectra, we confirm the source as a WN7 W–R star. Note that the source was previously characterized as a WN7 W–R-type star using the K-band spectrum (with a spectral resolution of ∼1200) by Shara et al. (2012). Our K-band spectrum (with a spectral resolution of ∼1000) is similar to the spectrum reported in Shara et al. (2012). Similar characteristics of broad emission lines of He and N are seen in the present study as well as previously published work. We have given equivalent widths of our observed lines in Table 1, which are consistent with the NIR lines listed in Table 5 in Shara et al. (2012). They utilized the 2MASS colors of the source and estimated an extinction to be∼0.84 mag in the K_s band. They also computed a distance to the W–R 1503-160L star of ∼3.1 kpc, assuming it is a strong-lined WN7 star.

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In this work, we have also utilized the widths of lines seen in our observed spectra and have estimated the terminal velocity of the emitting gas of the W–R star. In the first step, we determined the instrumental broadening of the lines using the isolated observed sky emission lines, which is found to be ∼430 km s⁻¹. Furthermore, the equivalent widths of the He ii+ Hα line at 6563 Å and the Brγ + He i + He ii line at 2.166 μm were used to determine the velocity of the expanding gas. Finally, after correcting the velocity for the instrumental broadening, we obtain the terminal velocity of the outflowing gas to be ∼1600 ± 400 km s⁻¹.

Crowther et al. (2006) extensively studied the properties of W–R stars in a super cluster, Westerlund 1, and suggested a criterion to infer the strong-lined and weak-lined WN7 W–R stars. They mentioned that the strong-lined WN7 W–R stars have FWHMs of the He ii 2.1885 μm line of ≥130 Å. They also listed the absolute magnitudes of strong-lined and weak-lined WN7 W–R stars, which are different for these two subclasses. The absolute magnitudes (in the K_s band) of strong-lined and weak-lined WN7 W–R stars were reported to be −4.77 mag and −5.92 mag, respectively (see Table A1 in Crowther et al. 2006), which will lead two different estimates of distance. Taking into account these absolute magnitudes along with 2MASS photometry (m${}_{{K}_{s}}$ = 8.51 mag) and extinction (A${}_{{K}_{s}}$ ∼0.84 mag; as mentioned above) in the K_s band, we compute the distances to strong-lined and weak-lined WN7 W–R 1503-160L stars of 3.1 ± 0.7 kpc and 5.2 ± 0.7 kpc, respectively. Using our K-band spectra, the FWHM of the He ii 2.1885 μm line of W–R 1503-160L is about 110 Å, which implies the source to be a weak-lined WN7 star. Considering this result and our calculations, we adpot the distance to the weak-lined WN7 W–R star of 5.2 ± 0.7 kpc, which is in good agreement with the kinematic distance of the bubble N46. Hence, taken together, we suggest that the bubble N46 hosts the W–R 1503-160L star.

The knowledge of the ionized, dust, and molecular emissions in a given star-forming complex allows us to infer physical conditions, and to provide details about H ii regions, possible outflows, and photon-dominated regions. Additionally, the distribution of molecular gas in a given star-forming region also helps us to trace its exact physical extension/boundary.

3.2.1. Global View

In Figure 2, we show the spatial distribution of ionized emission, warm dust emission, and molecular gas around the bubble N46. Figure 2(a) is a three-color composite image made using MAGPIS 20 cm in red, Herschel 70 μm in green, and GLIMPSE 8.0 μm in blue. The radio continuum map at MAGPIS 20 cm reveals several ionized regions. The image at 70 μm represents the warm dust emission, whereas the 8 μm band is dominated by the 7.7 and 8.6 μm polycyclic aromatic hydrocarbon (PAH) emission (including the continuum). In Figure 2(b), we show the molecular ¹³CO (J = 1–0) gas emission in the direction of the bubble N46. The ¹³CO profile depicts the bubble region in a velocity range of 87–100 km s⁻¹. The supernova remnant (SNR) G027.3+00.0 (also known as Kes 73) is seen as the most prominent ionized region in the radio map (see Figure 2(a)). The radio map also traces some compact point-like radio sources. The composite map also displays the edges of the bubble as a prominent bright feature coincident with the 20 cm emission (see Figure 2(a)). In Figure 2(a), we have marked the positions of some known astronomical objects (see Table 2 for coordinates), which are seen in the global view of the bubble N46. As mentioned before, the bubble harbors the WN7 W–R 1503-160L star that is situated at a projected linear separation of ∼867 from the SNR G027.3+00.0. Ferrand & Safi-Harb (2012) tabulated the distance and age of SNR G027.3+00.0 to be 7.5–9.1 kpc and 750–2100 years, respectively. Using photometric and spectroscopic analyses (as mentioned earlier), we find the distance of the W–R 1503-160L star to be 5.2 ± 0.7 kpc, which is consistent with the distance of the bubble (i.e., 4.8 ± 1.4 kpc). In general, W–R stars have very high mass-loss rates ((2–5) × 10⁻⁵ M_⊙ yr⁻¹) with terminal velocities of 1000–2500 km s⁻¹ (Prinja et al. 1990) and have lifetimes of roughly a few million years (Lamers & Cassinelli 1999). Considering the distance of the bubble N46, these two sources, SNR G027.3+00.0 and bubble N46, do not appear to be physically linked. Therefore, the impact of SNR G027.3+00.0 on the surroundings of the bubble N46 is unlikely.

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Table 2.  Coordinates of Some Referred Astronomical Objects Seen in a Global View of the Bubble N46, as Marked in Figures 2 and 3

ID	Galactic	Galactic	R.A.	decl.
	Longitude (degree)	Latitude (degree)	[J2000]	[J2000]
Bubble N46	27.3100	−0.1100	18:41:33.0	−05:03:07.6
W–R 1503-160L	27.2987	−0.1207	18:41:34.1	−05:04:01.2
IRAS 18385−0512	27.1853	−0.0822	18:41:13.3	−05:09:01.1
IRAS 18391−0504	27.3689	−0.1598	18:41:50.2	−05:01:21.0
SNR G027.3+00.0	27.3866	−0.0060	18:41:19.2	−04:56:11.0
IRDC (SDC G27.356−0.152)	27.3568	−0.1530	18:41:47.4	−05:01:48.5
6.7 GHz methanol maser	27.3650	−0.1660	18:41:51.0	−05:01:42.8
N46clm	27.2940	−0.1540	18:41:40.7	−05:05:11.0
clm1	27.4653	−0.1510	18:41:58.9	−04:55:58.0
clm2	27.2418	−0.1751	18:41:39.5	−05:08:33.0
clm3	27.4653	−0.1510	18:41:23.3	−05:04:59.2

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3.2.2. Local Environment

In Figure 2(b), based on the distribution of molecular emission, we have selected the region around the bubble for our present study, where a majority of molecular gas is distributed (hereafter, N46 molecular cloud; see the dotted–dashed box in Figure 2(b)). In Figure 3, we show the multi-wavelength images of the selected region around the bubble N46 at 12, 160, 250, 350, 500, 1100 μm, integrated ¹³CO emission, and 20 cm. The radio continuum and integrated CO maps are also shown for comparison with the infrared and sub-mm images. The details of the integrated CO map as well as the kinematics of molecular gas are given in Section 3.3. We have highlighted the positions of IRAS 18385−0512, IRAS 18391−0504, 6.7 GHz methanol maser emission (MME), and the W–R star. The WISE 12 μm band traces the warm dust emission and contains the PAH emission feature at 11.3 μm. The 160–1100 μm images indicate the presence of the cold dust emission (see Section 3.4 for quantitative estimates). We identify about five molecular condensations (e.g., clm1, clm2, clm3, MME, and N46clm) based on visual inspection of the integrated molecular map, which are labeled in the map. The MME condensation is found to be associated with an infrared dark cloud (IRDC) in the 8 μm image. The IRDC appears as a bright in emission at wavelengths longer than 24 μm. In Figure 4(a), we show a color composite image made using 8 μm (blue), 24 μm (green), and 70 μm (red) images. In Figure 4(a), the inset on the bottom right shows the zoomed in view toward the MME and IRAS 18391−0504 (a color composite image: 5.8 μm (red), 4.5 μm (green), and 3.6 μm (blue) images). The observed MME is located ∼253 away from IRAS 18391−0504 (see the inset in Figure 4(a)). In Figures 4(b) and (c), we present the zoomed in view of the bubble N46, which depict the edges of the bubble. Figure 4(c) shows the absence of radio continuum peak toward the W–R star located inside the bubble. Note that the 20 cm continuum emission is lacking inside the bubble except a detection of a point-like radio source. Additionally, the 24 μm emission, which traces the warm dust emission, is not enclosed by the 8 μm emission. It indicates that the bubble structure is unlikely to be originated by the expanding H ii region.

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Shirley et al. (2013) reported the molecular HCO⁺ (3–2) gas velocities to be ∼91.1 and ∼89.2 km s⁻¹ toward the N46clm and clm1 condensations, respectively. The peak velocity of 6.7 GHz MME is ∼99.7 km s⁻¹ (Szymczak et al. 2012). Ammonia (NH₃) line observations were also reported toward MME condensation by Wienen et al. (2012). They found the radial velocities of the NH₃ gas to be between 91.21 and 91.70 km s⁻¹ and also estimated the near kinematic distance as 5.06 kpc. However, the line of sight velocity of molecular gas toward IRAS 18385−0512 is ∼26 km s⁻¹ (Sridharan et al. 2002). With knowledge of these velocities, it appears that the bubble N46, the W–R star, and five condensations (including IRAS 18391−0504 and 6.7 GHz MME) belong to the same region (see Figure 4(b)); however, IRAS 18385−0512 is not physically associated with the N46 molecular cloud. Considering this fact, we do not discuss the results related to IRAS 18385−0512 in the present work.

All together, with the help of molecular line data and the distribution of cold, warm, and ionized emission, we find the physical association of the W–R star and molecular condensations located inside the N46 molecular cloud. Furthermore, the ionized emission appears to be well correlated with the warm dust emission at the edges of the bubble.

3.2.3. Compact Radio Sources

3.2.4. Infrared Counterpart of CRSs and MME

This section deals with a kinematic analysis of the molecular gas in the N46 molecular cloud. The kinematics of the molecular gas may help us to explore the expansion of the bubble. Figure 5 shows the integrated GRS ¹³CO (J = 1−0) velocity channel maps (at intervals of 1 km s⁻¹), revealing different molecular condensations along the line of sight. The channel maps show a clear separation of the high- and low-velocity emission toward the MME condensation. In the velocity range from 88 to 95 km s⁻¹, the highlighted molecular condensations (i.e., clm1, clm2, clm3, MME, and N46clm) are well detected. Using GRS ¹³CO (J = 1−0) and radio continuum data, Anderson et al. (2009) examined the properties of the molecular cloud associated with the Galactic H ii regions including the bubble N46, which is referred to as C27.31−0.14 (V_lsr ∼ 92.3 km s⁻¹) in their catalog. The integrated GRS ¹³CO intensity map⁵ was provided by these authors, however, the position–velocity analysis of this cloud is still lacking. In Figure 6, we present the integrated ¹³CO intensity map and the position–velocity maps. The galactic position–velocity diagrams of the ¹³CO emission reveals an almost inverted C-like structure as well as the presence of a noticeable velocity gradient toward the 6.7 GHz MME (see Figures 6(b) and (d)). In Figure 6(c), we show the bubble location traced at 8 μm along with the distribution of cold dust emission. As previously mentioned, the 6.7 GHz MME represents MSF and is associated with the densest clump in the N46 molecular cloud. Additionally, the MME condensation contains a cluster of young populations (see Section 3.5.2). The position–velocity diagrams convincingly suggest the presence of an outflow activity toward the MME condensation. In the channel maps, the orientations of the high- and low-velocity lobes appear significantly different (see the velocity range between 88 to 95 km s⁻¹ in Figure 5). These features suggest the presence of multiple CO outflows in the MME condensation. We cannot further explore this aspect in this work due to the coarse beam of the GRS CO data (beam size 45''). The existence of an inverted C-like structure in the position–velocity plot often indicates an expanding shell (Arce et al. 2011). Arce et al. (2011) examined the observed molecular gas distribution in the Perseus molecular cloud along with a modeling of expanding bubbles in a turbulent medium. They indicated that the semi-ring-like or C-like structure in the position–velocity plot represents an expanding shell. The position of the W–R star is approximately situated at the center of the inverted C-like structure. Our position–velocity analysis indicates the presence of an expanding shell with an expansion velocity of the gas to be ∼3 km s⁻¹.

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Taken together, the CO kinematics suggest the presence of at least one outflow as well as the expanding shell.

In this section, we study the distribution of column density, temperature, extinction, and clump mass in the N46 molecular cloud. We followed the procedures given in Mallick et al. (2015) to obtain the Herschel temperature and column density maps (see below). We obtained these maps from a pixel-by-pixel SED fit with a modified blackbody curve to the cold dust emission in the Herschel 160–500 μm wavelengths. The images at 250–500 μm are in the surface brightness unit of MJy sr⁻¹, while the image at 160 μm is calibrated in the units of Jy pixel⁻¹. The plate scales of the 160, 250, 350, and 500 μm images are 6.4, 6, 10, and 14 arcsec pixel⁻¹, respectively.

Here, we provide a brief step-by-step description of the procedures. All images were first converted to Jy pixel⁻¹ unit and were convolved to the angular resolution of 500 μm image (∼37'') using the convolution kernels available in the HIPE software. The images were regridded to the pixel size of a 500 μm image (∼14'') using the HIPE software. Next, the sky background flux level was measured to be 1.476, 3.705, 1.759, and 0.643 Jy pixel⁻¹ for the 160, 250, 350, and 500 μm images (the size of the selected region ∼66 × 73; central coordinates: α_J2000 = 18^h44^m45^s, δ_J2000 = −05°40'00''), respectively. The featureless dark region far from the N46 molecular cloud, to avoid diffuse emission associated with target, was carefully chosen for the background estimation.

Finally, we fitted a modified blackbody to the observed fluxes on a pixel-by-pixel basis to obtain the maps (see Equations (8) and (9) given in Mallick et al. 2015). We performed the fitting using the four data points for each pixel, retaining the dust temperature (T_d) and the column density (N(H₂)) as free parameters. In the estimations, we used a mean molecular weight per hydrogen molecule (μ_H2=) 2.8 (Kauffmann et al. 2008) and an absorption coefficient (κ_ν=) 0.1 (ν/1000 GHz)^β cm² g⁻¹, including a gas-to-dust ratio (R_t=) of 100, with a dust spectral index of β = 2 (see Hildebrand 1983). In Figure 7, we show the final temperature and column density maps (resolution ∼37'') of our selected region around the bubble N46. In Figure 7(a), we find warmer gas (T_d ∼ 23–27 K) toward the edges of the bubble. The Herschel temperature map also exhibits a temperature range of about 18–24 K toward the five molecular condensations as marked in Figure 3. We find the peak column densities of about 9.2 × 10²¹ (A_V ∼ 10 mag), 6.31 × 10²¹ (A_V ∼ 7 mag), 1.18 × 10²² (A_V ∼ 12 mag), 9.2 × 10²² (A_V ∼ 98 mag), and 9.17 × 10²¹ (A_V ∼ 10 mag) cm⁻² toward the molecular condensations clm1, clm2, clm3, MME, and N46clm, respectively (see Figure 7(b)). Here, we used the relation between optical extinction and hydrogen column density from Bohlin et al. (1978; i.e., A_V = 1.07 × 10⁻²¹ N(H₂)). Note that the clump associated with 6.7 GHz MME is relatively warmer (T_d ∼ 24 K) and is the densest one compared to other condensations. Using NH₃ line observations, Wienen et al. (2012) reported physical parameters of dense gas (i.e., the gas kinetic temperature and rotational temperature) toward the MME clump. The kinetic temperature of the MME clump was estimated to be ∼23.81 K, which is consistent with our estimated temperature value.

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We estimate the masses of four prominent clumps (MME, N46clm, clm1, and clm2) seen in the column density map (see Figure 7(b)). The mass of a single clump can be obtained using the formula:

$$\begin{eqnarray}&&{M}_{\mathrm{clump}}={\mu }_{{{\rm{H}}}_{2}}{m}_{{\rm{H}}}{\mathrm{Area}}_{\mathrm{pix}}{\rm{\Sigma }}N({{\rm{H}}}_{2}),\end{eqnarray} \tag{ 1 }$$

where ${\mu }_{{{\rm{H}}}_{2}}$ is assumed to be 2.8, Area_pix is the area subtended by one pixel, and ${\rm{\Sigma }}N({{\rm{H}}}_{2})$ is the total column density. We employed the "CLUMPFIND" program to estimate the total column densities of the clumps. Using Equation (1), we compute the masses of the clumps MME, N46clm, clm1, and clm2 to be ∼7273, ∼954, ∼412, and ∼741 M_⊙, respectively.

3.5.1. Selection of Infrared Excess Sources

The study of YSOs is considered a primary tracer of the star-formation activity in a given star-forming region. In the following, we use different photometric surveys to identify and classify YSOs.

• 1.  
  We selected sources having detections in the MIPSGAL 24 μm and GLIMPSE 3.6 μm bands. The color–magnitude plot ([3.6]–[24]/[3.6]) has been utilized to trace the youngest population in a given star-forming region (Guieu et al. 2010; Rebull et al. 2011; Dewangan et al. 2015a). In our selected region, we find 135 sources that are common in the 3.6 and 24 μm bands. In Figure 8(a), we present the [3.6]–[24]/[3.6] plot of 135 sources. We select 32 YSOs (5 Class I; 2 Flat-spectrum; 25 Class II) and 103 Class III sources. The boundary of different stages of YSOs is also marked in Figure 8(a) (e.g., Guieu et al. 2010), which depicts the criteria to infer the youngest population. The figure also shows the boundary of possible contaminants (i.e., galaxies and disk-less stars; see Figure 10 in Rebull et al. 2011). We do not identify any contaminants (i.e., galaxies) in our selected YSOs.
• 2.  
  We obtained sources having detections in all four GLIMPSE bands. Following the Gutermuth et al. (2009) schemes, we identified YSOs and various possible contaminants (e.g., broad-line active galactic nuclei (AGNs), PAH-emitting galaxies, shocked emission blobs/knots, and PAH-emission-contaminated apertures) in our selected region around the bubble. The possible contaminants are also removed from the selected YSOs. We further utilized the slopes of the Spitzer-GLIMPSE SED (${\alpha }_{\mathrm{IRAC}}$) measured from 3.6 to 8.0 μm (e.g., Lada et al. 2006) and classified the selected YSOs into different evolutionary stages (i.e., Class I (α_IRAC > −0.3), Class II (−0.3 > α_IRAC > −1.6), and Class III (−1.6 > α_IRAC > −2.56)). One can also find more details of YSO classifications in Dewangan & Anandarao (2011, and references therein). In Figure 8(b), we show the GLIMPSE color–color diagram ([3.6]−[4.5] versus[5.8]−[8.0]) for all of the identified sources. In the end, we find 19 YSOs (3 Class I; 16 Class II), 2545 photospheres, 3 PAH-emitting galaxies, and 235 contaminants.
• 3.  
  Infrared excess traced in the first three GLIMPSE bands (except 8.0 μm band) can also offer to identify the additional protostars. The color–color space ([4.5]–[5.8] versus [3.6]–[4.5]) is utilized to trace the infrared excess sources. We use color conditions, [4.5]–[5.8] ≥ 0.7 and [3.6]–[4.5] ≥ 0.7, to select protostars, as given in Hartmann et al. (2005) and Getman et al. (2007). We identify four protostars in our selected region around the bubble (see Figure 8(c)).
• 4.  
  The NIR H–K color excess also allows us to identify additional YSOs. In order to choose a color criterion, we examined the color–magnitude (H−K/K) plot of the nearby control field (size ∼6' × 6'; central coordinates: α_J2000 = 18^h41^m536, δ_J2000 = −05°14'555) and obtained a color H−K value (i.e., ∼2.0) that isolates large H−K excess sources from the rest of the population. Adopting this color H−K cut-off condition, we obtained 315 embedded YSOs in our selected region around the bubble (see Figure 8(d)).

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Finally, using all the four schemes mentioned above, we obtain a total of 370 YSOs in our selected region around the bubble (as shown in Figure 4(a)). The positions of all selected YSOs are shown in Figure 9(a).

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3.5.2. Study of Distribution of YSOs

3.5.3. The SED Modeling

In this section, we performed the SED modeling of IRcs of CRS7, CRS9, CRS10, and N46clm (see Section 3.2) using an online SED fitting tool (Robitaille et al. 2006, 2007). The SED fitting tool needs a minimum of three data points with good quality and fits the data allowing the distance to the source and a range of visual extinction values as free parameters. The accretion scenario is incorporated into the SED models, which assume a central source associated with rotationally flattened infalling envelope, bipolar cavities, and a flared accretion disk. The model grid encompasses a total of 200,000 SED models (Robitaille et al. 2006), covering a range of stellar masses from 0.1 to 50 M_⊙. The SED fitting tool fits only those models which satisfy the criterion χ²–${\chi }_{\mathrm{best}}^{2}$ < 3, where χ² is taken per data point. The resultant fitted SED models of these IRcs are presented in Figure 10. Reliable photometric magnitudes of IRcs have been used for modeling, which can be seen as data points in the fitted SED plots. We provided the visual extinction in the range 0–100 mag, as an input parameter for the SED modeling. The weighted mean values of the stellar mass (extinction) of IRcs of CRS7, CRS9, CRS10, and N46clm are 20 ± 0.1 M_⊙ (6.8 ± 0.1 mag), 14 ± 0.6 M_⊙ (9.4 ± 0.1 mag), 13 ± 1 M_⊙ (13.5 ± 1.5 mag), and 9 ± 3 M_⊙ (54 ± 19 mag), respectively. The weighted mean values of the stellar age of IRcs of CRS7, CRS9, CRS10, and N46clm are estimated to be 1.2, 1.3, 1.8, and 0.67 Myr, respectively. Note that the IRcs of CRS10 and N46clm are identified as YSOs. Our SED results are consistent with inferred radio spectral types of the CRSs and further suggest the presence of massive stars on the edges of the bubble.

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In the following, we study the large-scale morphology of the plane-of-the-sky projection of the magnetic field toward the bubble. The polarization of background starlight is a useful tool to trace the projected plane-of-the-sky magnetic field morphology. The polarization vectors of background stars give the field direction in the plane of the sky parallel to the direction of polarization (Davis & Greenstein 1951).

In Figure 11(a), we display the H-band polarization vectors overlaid on the Herschel temperature map. The polarization data toward our selected region around the bubble were retrieved from the GPIPS (see Clemens et al. 2012 for more details) and were covered in five fields i.e., GP0768, GP0769, GP0782, GP0783, GP0796, and GP0797. In order to obtain the reliable polarimetric information, we selected sources with Usage Flag (UF) = 1 and P/σ_p ≥ 5. Following these conditions, we obtain a total of 249 stars. We also find the H-band polarization of CRS7, CRS10, and the W–R stars. These sources have suffered extinction, as listed in previous sections. Therefore, it appears that the H-band polarization could be related to the sources and may not be associated with the diffuse interstellar medium. Hence, we have shown the polarization vectors of these sources with an addition of 90° in their respective position angles (see the red vectors in Figure 11(a)). To examine the distribution of H-band polarization, we show mean polarization vectors in Figure 11(b). In order to infer the mean polarization, our selected spatial area is divided into 13 × 7 equal divisions and a mean polarization value is computed using the Q and U Stokes parameters of the H-band sources located inside each specific division. In Figure 11, the degree of polarization is traced by the length of a vector, whereas the angle of a vector indicates the polarization galactic position angle. Note that the H-band polarization data cannot allow us to infer the morphology of the plane-of-the-sky projection of the magnetic field toward the dense clumps (clm1, clm2, clm3, MME, and N46clm).

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The mean distribution of the H-band vectors appears uniform, however, we notice a little change in the mean polarization pattern between the MME condensation and the W–R star.

The histograms of the degree of polarization and the polarization galactic position angles of the 249 stars are displayed in Figures 12(a) and (b), respectively. The degree of polarization value for the majority of the population is found to be ∼3%. Additionally, we infer an ordered plane-of-the-sky component of the magnetic field with a peak galactic position angle of ∼55°. In Figure 12(c), the GPS NIR color–color diagram of sources confirms that the majority of stars are located behind the N46 molecular cloud. In the NIR color–color diagram, the reddened background stars and/or embedded stars can be identified with (J−H) ≥ 1.0; however the foreground sources can be traced with (J−H) < 1.0.

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The presence of extended nebulous features around the O and/or W–R stars is often explained by the interaction between the massive stars and their surroundings (Lamers & Cassinelli 1999; Deharveng et al. 2010; Dewangan et al. 2015a, 2015b). Many similar studies, in particular those surrounding W–R stars, have been carried out in recent years (e.g., WR 130, Cichowolski et al. 2015; WR153ab/HD 211853, Vasquez et al. 2010; Liu et al. 2012; WR 23, Cappa et al. 2005; and WR 55, Cappa et al. 2009). The knowledge of various pressure components (pressure of an H ii region (P_{H ii}), radiation pressure (P_rad), and stellar wind ram pressure (P_wind)) driven by a massive star can be useful to infer the prominent physical feedback process. A total pressure driven by a massive star on its surroundings can be written as the sum of these three pressure components, i.e., P_total = ${P}_{{\rm{H}}{\rm{II}}}$ + P_rad + P_wind. In recent years, it is often found that the 8 μm emission surrounds the warm dust emission as well as the ionized emission (e.g., Deharveng et al. 2010; Paladini et al. 2012). In some such sites, it has been reported that the photoionized gas associated with H ii regions appears as the major contributor for the feedback process (Dewangan et al. 2015a, 2015b). However, in the case of bubble N46, the lack of extended ionized emission inside the bubble strongly indicates that the bubble is unlikely to have originated by the expanding H ii region. The spectroscopically characterized source WN7 W–R star appears to be located inside the bubble N46. Here, we estimate only two pressure components P_wind (=${\dot{M}}_{w}{V}_{w}/4\pi {D}_{s}^{2}$) and P_rad (=${L}_{\mathrm{bol}}/4\pi {{cD}}_{s}^{2}$), where ${\dot{M}}_{w}$ is the mass-loss rate, V_w is the wind velocity, L_bol is the bolometric luminosity, and D_s is the distance from the location of the W–R star where the pressure components are computed. Adopting the typical values of ${\dot{M}}_{w}$ (≈5.0 × 10⁻⁵ M_⊙ yr⁻¹) and L_bol (∼10⁶ L_⊙) for the WN7 W–R star given in Lamers & Cassinelli (1999) and Prinja et al. (1990), along with V_w (=1600 km s⁻¹) estimated using our observed spectra (see Section 3.1), we obtain the values of P_wind ≈ 1.1 × 10⁻⁹ dynes cm⁻² and P_rad ≈ 2.88 × 10⁻¹⁰ dynes cm⁻² at a distance of 1.93 pc (i.e., an average radius of the bubble). We find P_wind > P_rad. Additionally, P_wind is much higher than the pressure associated with a typical cool molecular cloud (P_MC ∼10⁻¹¹–10⁻¹² dynes cm⁻² for a temperature ∼20 K and the particle density ∼10³–10⁴ cm⁻³; see Table 7.3 of Dyson & Williams 1980, p. 204). These calculations indicate that the bubble appears to be produced by the stellar wind of the WN7 W–R star. Hence, the bubble N46 can be considered as an example of a stellar wind bubble or a wind-driven bubble (e.g., Castor et al. 1975). From the position–velocity analysis of molecular gas, we infer the presence of an expanding shell with an expansion velocity, V_exp = 3 km s⁻¹.

We estimate the expected mechanical luminosity of the stellar wind (L_w = $0.5\;{\dot{M}}_{w}$ ${V}_{w}^{2}$ erg s⁻¹) for the WN7 W–R star using the values of ${\dot{M}}_{w}$ and V_w as mentioned above. We obtain L_w ≈ 4 × 10³⁷ erg s⁻¹ for the WN7 W–R star. Using the value of L_w, we infer that the W–R star has injected a large amount of mechanical energy (E_w ≈ 1.3 × 10⁵¹ erg) in 1 Myr. During the entire life of a typical WN7 star (4 Myr; see Figure 13.3 given in Lamers & Cassinelli 1999), the star can dump E_w ≈ 5 × 10⁵¹ erg in its vicinity. In the environment surrounding WR153ab, Vasquez et al. (2010) also found that the mechanical energy released by the W–R star greatly influenced its vicinity, shaping the ring-like structure of the molecular distribution.

The starlight H-band polarization data allow us to examine the large-scale magnetic field structure projected in sky plane. The magnetic field lines are almost uniformly distributed within the N46 molecular cloud. However, we notice a variation in the distribution of the mean polarization position angles between the MME condensation and the W–R star. We have also found that P_wind is the dominated term compared to P_rad. It is known that the aligned dust can be used as a tracer for magnetic fields. In the case of wind bubbles, it is often found that only some amount of the injected mechanical energy will be transferred into kinetic energy of the expanding gas (Weaver et al. 1977; Arthur 2007). Hence, the remaining energy propagates in the surrounding interstellar medium. Considering this fact, the impact of the stellar wind of the W–R star seems to be responsible for the noticeable change in the polarization angles. Note that in this work we are unable to trace the small-scale magnetic field structure toward the dense regions.

In the N46 molecular cloud, SF activities are traced by the presence of compact H ii regions, clusters of YSOs, 6.7 GHz MME, and molecular outflow. In our clustering analysis, we found that the clusters of YSOs are spatially seen toward the molecular condensations, which are associated with a range of temperature and density of about 18–24 K and (0.6–9.2) × 10²² cm⁻² (A_V ∼ 7–98 mag). Some YSOs are also seen on the edges of the bubble without any clustering. Several compact radio peaks are identified on the edges of the bubble and are consistent with radio spectral types B0V–B0.5V. Infrared images revealed the individual IRcs of three CRSs that have masses (13–20 M_⊙) and ages (1–2 Myr). It is evident that SF continues into the N46 molecular cloud. Interestingly, our observational investigation suggests that the early phase of MSF (<0.1 Myr) is occurring in the MME condensation, as traced by the presence of the 6.7 GHz MME. This condensation is the most massive clump and is associated with significant clustering of YSOs and at least one molecular outflow. An embedded IRc of the 6.7 GHz MME is identified and is located within the YSOs cluster. Due to coarse beam size, the GRS ¹³CO data are limited and cannot allow us to study the gas kinematics within MME condensation. Hence, high-resolution molecular line observations will be helpful to further explore the MME. We also identify an embedded young massive protostar (Mass ∼9 M_⊙, A_V ∼ 54 mag, and age ∼0.7 Myr) associated with the N46clm clump, which is located near the edge of the bubble. The average ages of Class I and Class II YSOs are generally reported to be ∼0.44 Myr and ∼1–3 Myr (Evans et al. 2009), respectively. In the N46 molecular cloud, we find the presence of an evolved and powerful WN7 W–R star (with typical age ∼4 Myr). It appears that the strong stellar wind from the W–R star has affected its local vicinity, which has been inferred from the presence of a wind-driven bubble/shell. It has been suggested that the stellar winds from O and B stars can induce the formation of massive stars in molecular clouds (Lucy & Solomon 1967, 1970; Morton 1967). We notice the existence of an apparent age gradient between the W–R star and young sources (i.e., YSOs, IRcs of CRSs, and 6.7 GHz MME). Hence, it is most likely that the stellar wind from the WN7 W–R star plays a significant role in the propagation of SF in the N46 molecular cloud. Previously, several authors also presented results in favor of SF induced by the strong wind of the WR stars (e.g., WR 130, Cichowolski et al. 2015; WR153ab/HD 211853, Vasquez et al. 2010; Liu et al. 2012; WR 23, Cappa et al. 2005; and WR 55, Cappa et al. 2009).

To assess the possibility of triggered SF, we computed P_wind values at different distances with respect to the W–R star. The W–R star is situated at a projected linear separation of ∼3.4, ∼4.4, ∼6.5, ∼6.7, and ∼14.1 pc from the N46clm, clm3, clm2, MME, and clm1 molecular condensations. In previous section, we estimated P_wind near the edge, which is equal to be ∼1.1 × 10⁻⁹ dynes cm⁻² at D_s = 1.93 pc and is much higher than P_MC. We also estimate P_wind ∼0.94 × 10⁻¹⁰ and ∼2.1 × 10⁻¹¹ dynes cm⁻² at D_s = 6.7 and 14.1 pc, respectively. The clm1 condensation is located farthest from the W–R star, indicating the impact of wind may not be very strong. Considering the tremendous energy output of the W–R star, the SF activity could be triggered on the edges of the bubble. Furthermore, SF in the N46clm, clm2, clm3, and MME condensations might also be influenced by the feedback of the WN7 W–R star. However, the triggered SF in the clm1 condensation seems unlikely.