1 Draft version September 8, 2021 Typeset using LATEX manuscript style in AASTeX63
The Sejong Open Cluster Survey (SOS). VII. A Photometric Study of the Young Open Cluster IC 1590
1 2, 3 4 5 1 2 Seulgi Kim, Beomdu Lim, Michael S. Bessell, Jinyoung S. Kim, and Hwankyung Sung
1 3 Department of Astronomy and Space Science
4 Sejong University, 209 Neungdong-ro, Gwangjin-gu, Seoul 05006, Republic of Korea
2 5 School of Space Research, Kyung Hee University, 1732, Deogyeong-daero, Giheung-gu, Yongin-si, Gyeonggi-do 17104,
6 Republic of Korea
3 7 Korea Astronomy and Space Science Institute, 776 Daedeokdae-ro, Yuseong-gu, Daejeon 34055, Republic of Korea
4 8 Research School of Astronomy & Astrophysics, The Australian National University, Canberra, ACT 2611, Australia
5 9 Steward Observatory, University of Arizona, 933 N. Cherry Ave., Tucson, AZ 85721-0065, USA
10 (Received; Revised; Accepted)
11 Submitted to AJ
12 ABSTRACT
13 Young open clusters are ideal laboratories to understand star formation process. We
14 present deep UBVI and Hα photometry for the young open cluster IC 1590 in the center
15 of the H II region NGC 281. Early-type members are selected from UBV photometric
16 diagrams, and low-mass pre-main sequence (PMS) members are identified by using Hα
17 photometry. In addition, the published X-ray source list and Gaia astrometric data are
18 also used to isolate probable members. A total of 408 stars are selected as members. The
19 mean reddening obtained from early-type members is hE(B − V )i = 0.40 ± 0.06 (s.d.).
20 We confirm the abnormal extinction law for the intracluster medium. The distance
Corresponding author: Beomdu Lim [email protected] 2 Kim et al.
21 modulus to the cluster determined from the zero-age main-sequence fitting method is
+0.24 22 12.3±0.2 mag (d = 2.88±0.28 kpc), which is consistent with the distance d = 2.70−0.20
23 kpc from the recent Gaia parallaxes. We also estimate the ages and masses of individual
24 members by means of stellar evolutionary models. The mode of the age of PMS stars
25 is about 0.8 Myr. The initial mass function of IC 1590 is derived. It appears a steeper
26 shape (Γ = −1.49 ± 0.14) than that of the Salpeter/Kroupa initial mass function for
27 the high mass regime (m > 1 M ). The signature of mass segregation is detected from
0 28 the difference in the slopes of the initial mass functions for the inner (r < 2. 5) and
29 outer region of this cluster. We finally discuss the star formation history in NGC 281.
30 Keywords: open clusters and associations: individual (IC 1590) — stars: formation —
31 stars: pre-main sequence — stars: luminosity function, mass function
32 1. INTRODUCTION
33 Young open clusters are the prime targets to study star formation processes. In particular, they
34 make it possible to obtain distance, age, luminosity and mass function because their members have
35 the same origin. These fundamental parameters are basic information to understand the formation
36 and dynamical evolution of stellar clusters (Lada & Lada 2003; Sagar et al. 1986). Young open
37 clusters can also be used to verify stellar evolution theory in a wide range of stellar masses (Sung et
38 al. 2013) and trace the spiral structure of the Galaxy as they are the building blocks of the Galactic
39 disk (Carraro et al. 1998; Chen et al. 2003; Piskunov et al. 2006). In addition, young stars in open
40 clusters are now being targeted for probing the formation processes of exoplanets and investigating
41 properties of protoplanetary disks where planets could be formed (e.g., Alcal´aet al. 2017; Eisner et
42 al. 2018; Ansdell et al. 2020; Sanchis et al. 2020).
43 Stellar photometry is a basic tool to study the physical properties of young stellar systems Lim et al.
44 2015a; Abdelaziz et al. 2020. However, since most young open clusters are distributed in the Galactic
45 plane, it is essential to discriminate cluster members from field interlopers to determine the reliable
46 results. Classical T-Tauri stars (CTTSs) tend to show a prominent UV excess and Hα emission due IC 1590 3
47 to accretion activity, as well as a mid-infrared (MIR) excess from their warm circumstellar disks
48 (Sung et al. 1997, 2009; Megeath et al. 2012; Lada 1987). On the other hand, weak-line T-Tauri
49 stars (WTTSs) have no or little accretion activity, but are still very bright in X-ray (Feigelson et al.
50 2003). Therefore, a photometric study complemented by multi-wavelength data is required to select
51 genuine members.
52 The initial mass function (IMF), originally introduced by Salpeter(1955), is a fundamental tool for
53 understanding star formation processes. If the IMF is universal in all star forming regions (SFRs),
54 one can expect common factors to control the star formation process. Conversely, variety in the IMFs
55 implies that the star formation process may be dependent on the physical properties of the natal
56 cloud. The issue of the universality of the IMF is still an ongoing debate. Young open clusters are
57 ideal targets for the study of the IMF across a wide range of stellar masses because their members
58 are less affected by stellar evolution and dynamical evolution (Sung et al. 1997). The universality of
59 the IMF can therefore be examined through the study of many young clusters.
60 Star formation can be regulated by massive stars in young stellar clusters. If massive stars (earlier
61 than B3) are formed in a molecular cloud, HII regions can be created by their intense UV radia-
62 tion. Strong stellar winds during their lifetime and subsequent supernova events could push out the
63 surrounding material and may create a huge void-like structure called a superbubble or a galactic
64 chimney. As the size of the superbubble increases, the accumulated material in the peripheral regions
65 may induce subsequent star formation (Bally 2008).
66 In this context, we have dedicated a homogeneous photometric survey of young open clusters in
67 the Galaxy (Lim et al. 2011; Sung et al. 2013; Lim et al. 2014a,b, 2015a,b; Sung et al. 2019). This is
68 the seventh paper as part of this project. The target is IC 1590 in the center of the HII region NGC
h m ◦ 0 ◦ ◦ 69 281 (Sharpless 184, PacMan Nebula, α2000 = 00 52 , δ2000 = +56 4 and l = 123. 07, b = −6. 31)
70 (Guetter & Turner 1997). The most prominent feature of IC 1590 is the Trapezium-like system HD
71 5005 composed of four O-type stars (O4V((fc)), O9.7II-III, O8.5V(n), O9.5V) (Sota et al. 2011).
72 These massive stars are the main ionizing sources of this HII region. The cluster is known to be 4 Kim et al.
73 young (3.5 ∼ 4.4 Myr), and 3 kpc away from the Sun (Guetter & Turner 1997; Sharma et al. 2012;
74 Sato et al. 2008).
75 The NGC 281 region has some interesting structures that are related to star formation history.
76 The south-western region of IC 1590 is obscured by the adjoining molecular cloud NGC 281 West.
77 Several X-ray sources and mid-IR sources, as well as an H2O maser source are detected in the region.
12 13 78 Lee & Jung(2003) studied NGC 281 West in CO(J=1-0) and CO(J=1-0) and found that the
12 79 CO(J=1-0) emission peak is very close to the H2O maser source. Some indication of an outflow
80 were also found near the maser source, implying ongoing star formation in NGC 281 West (Snell et
81 al. 1990; Hodapp 1994; Wu et al. 1996). An age sequence between IC 1590 and NGC 281 West was
82 found. Class II young stellar objects (YSOs) are located close to the cluster, while Class 0/I YSOs
83 and molecular clumps are distributed away from the cluster (Megeath & Wilson 1997; Sharma et
84 al. 2012). Furthermore, VLA 20 cm continuum images reveal compact structures and edges at the
12 85 north of the CO peak (Megeath & Wilson 1997). Accordingly, NGC 281 West has been suggested
12 86 as a SFR triggered by IC 1590. Another SFR, NGC 281 East, was found from CO observation by
87 Elmegreen & Lada(1978). Three IRAS sources were embedded in this region, and a highly reddened
88 star was found in the clump at the northern edge of NGC 281 East. Megeath & Wilson(1997)
89 suggested that the star formation of NGC 281 East might have been started or strengthened by
90 shocks. However, it is still necessary to probe the physical associations between the SFRs and IC
91 1590.
92 The IMF of IC 1590 has been investigated in several studies. Guetter & Turner(1997) obtained a
93 shallow IMF (Γ = −1.00 ± 0.21) of IC 1590 from photoelectric and CCD photometry. From UBVI
94 photometric and Hα spectroscopic observations by 105-cm Schmidt telescope, Sharma et al.(2012)
95 also found the slope of IMF to be flatter (Γ = −1.11 ± 0.15) than those of the Salpeter/Kroupa IMF
96 (Γ = −1.35; Salpeter 1955; Kroupa 2001). This shallow IMF implies that this SFR is a favorable
97 site of massive star formation. However, these previous studies derived the IMF using bright stars
98 (m > 1.5 M by Guetter & Turner 1997; m > 2 M by Sharma et al. 2012). Their selection
99 of low-mass PMS stars may be less complete due to the limited data. With our deep photometry, IC 1590 5
100 complemented by several archival datasets, it is possible to derive the IMF down to 1 M . In addition,
101 the recent Gaia Early Data Release 3 (EDR3; Gaia Collaboration et al. 2020) allows us to reliably
102 select kinematic members. Therefore, the IMF of this cluster can be studied across a wide mass
103 range.
104 The main scientific goals of this study are to derive the fundamental parameters of IC 1590 (red-
105 dening, distance, age, and the IMF) and to study the star formation process in this SFR. To achieve
106 these aims, we conducted deep UBVI and Hα observations and investigated the properties of IC
107 1590 using new optical data and additional archival data. In Section2, we describe our observations,
108 data processing, and archival datasets. The member selection criteria are addressed in Section3. In
109 Section4, the reddening and distance of IC 1590 are determined. The IMF is derived in Section5.
110 Finally, we discuss the triggered star formation in this SFR and the properties of some individual
111 stars in Section6 and summary in Section7.
112 2. OBSERVATIONS AND ARCHIVAL DATA
113 2.1. Optical Photometry
114 2.1.1. Maidanak AZT-22 1.5m Observations
115 The UBVI CCD observations of the young open cluster IC 1590 were conducted on 2005 August
00 −1 116 9, using the AZT-22 1.5 m telescope with the SITe 2000×800 CCD (0. 266 pixel ) at the Maidanak
00 117 Astronomical Observatory in Uzbekistan. The typical seeing was ∼ 1 . We used two sets of exposure
118 times for each band to cover both bright and faint stars. To determine the atmospheric extinction
119 coefficients and photometric zero-points, we also observed several SAAO standard stars (Menzies et
120 al. 1991) in the Landolt standard fields and additional standard stars with extremely blue and red
121 colors (Kilkenny et al. 1998). The transformation relation to the standard system is addressed in
122 Lim et al.(2009). The field of view (FoV) of the Maidanak observations is shown in Figure1 (green
123 rectangle). We summarize our observation in Table1.
00 124 2.1.2. Kuiper 61 Telescope of Steward Observatory 6 Kim et al.
Table 1. Observing log
00 Filter Exposure time (s) Seeing ( ) k1λ k2λ ζλ Maidanak AZT-22 (2005. 8. 9)
U 30, 600 1.22 0.539 ± 0.022 0.023 21.560 ± 0.011
B 5, 300 1.09 0.324 ± 0.010 0.026 23.261 ± 0.006
V 3, 180 0.88 0.231 ± 0.009 - 23.396 ± 0.001
I 3, 60 0.74 0.139 ± 0.010 - 23.017 ± 0.009
Kuiper 6100 (2011. 10. 31)
U 15, 600 1.6 0.451 ± 0.007 0.002 ± 0.003 22.181 ± 0.008
B 7, 300 1.4 0.253 ± 0.005 0.025 ± 0.002 23.563 ± 0.007
V 5, 180 1.4 0.134 ± 0.005 - 23.566 ± 0.008
I 5, 120 1.2 0.043 ± 0.004 - 22.181 ± 0.008
Hα 30, 600 1.2 0.061 - 19.494
Note— k1λ is the primary atmospheric extinction coefficient, k2λ is the secondary atmospheric extinction
coefficient, ζλ is the photometric zero-point from Lim et al.(2009).
00 125 We performed additional UBVI and Hα observations with the Kuiper 61 Telescope (f/13.5) and
126 Mont4k CCD of the Steward Observatory at Mt. Bigelow in Arizona on 2011 October 31. The pixel
00 −1 0 0 127 scale is 0. 42 pixel in a 3 × 3 binning mode, and the FoV is about 9. 7 × 9. 7. The exposure times
128 and seeing for each filter are summarized in Table1. A number of standard stars from the same
129 catalogues used in the previous section were observed. Transformation to the standard system is
130 detailed in the Appendix of Lim et al.(2015a). IC 1590 7
10
5 ] n i m c r a
[ 0
. c e D
5
10 10 5 0 5 10 R.A. [arcmin]
Figure 1. Optical image of NGC 281 taken from the Digital Sky Survey 2. The green rectangular and white square represent our target regions covered by the Maidanak AZT-22 observations and the Kuiper 6100 telescope observations, respectively. The large blue box depicts the FoV of the Chandra X-ray observations. The positions of stars are relative to the coordinate R.A. = 00h52m49s, decl. = +56◦3704200 (J2000). 8 Kim et al.
0.20 0.20
0.15 0.15
0.10 0.10
0.05 0.05 ) I - V
0.00 V 0.00 ( 0.05 0.05
0.10 0.10
0.15 V = -0.009 ± 0.032 0.15 (V-I) = 0.001 ± 0.021
0.20 0.20 10 11 12 13 14 15 16 17 0.25 0.50 0.75 1.00 1.25 1.50 1.75 V V-I 0.20 0.20
0.15 0.15
0.10 0.10
0.05 0.05 ) ) B V - - B 0.00 U 0.00 ( ( 0.05 0.05
0.10 0.10
0.15 (B-V) = 0.005 ± 0.021 0.15 (U-B) = 0.025 ± 0.030
0.20 0.20 0.2 0.4 0.6 0.8 1.0 1.2 1.4 0.5 0.0 0.5 1.0 1.5 B-V U-B
Figure 2. Comparison of two sets of photometry. The difference ∆ means Kuiper 6100 data - Maidanak data. Stars brighter than V = 17 mag were used in the comparison to avoid the large photometric errors of faint stars. The thick lines and shaded areas represent mean and standard deviation, respectively. The mean difference and the standard deviation between the two sets of the photometry are displayed in the lower-left corner of each panel. IC 1590 9 2 emission α H obs N α H B − U V − B I − V V I 1 α H B − . Photometric Data for IC 1590 emission candidates are represented as h. VU α − Table 2 ) defined by Sung &). Bessell ( 2000 IB I + 2 − V − Hα m = α IVV index (H α ] [mag] [mag] [mag] [mag] [mag] [mag] [mag] [mag] [mag] [mag] [mag] [mag] ◦ 0 00 means H emission stars are represented as H and H α α H H ][ 1 2 h m s [ ID R.A.21 0 Dec. 52 20.922 +56 0 37 52 49.6 21.223 17.28 +56 0 38 52 17.3 21.524 18.46 18.21 +56 0 35 52 1.18 57.3 21.725 20.30 17.20 0 +56 52 39 2.09 0.87 21.8 3.526 19.45 +56 0 17.30 37 52 2.25 1.79 39.5 0.21 21.927 18.75 19.04 +56 0 39 52 -0.12 1.63 16.9 22.028 20.45 1.45 - 18.39 +56 0.01 0 39 52 1.41 13.2 22.129 20.71 1.21 - 16.83 0.01 -0.74 +56 0 37 52 2.32 1.17 41.6 22.130 19.23 0.80 0.01 0.01 18.74 -1.10 +56 0 36 52 2.40 52.2 22.3 21.28 0.09 0.04 - 0.01 - 0.01 15.88 +56 35 2.54 1.70 53.6 0.01 0.04 16.84 0.02 0.01 19.55 0.96 0.01 - 0.08 - 0.01 21.00 0.05 - - 0.01 5 1.45 0.76 4 0.04 4 0.01 - 3 2 -0.41 1 - 0.02 0.79 0.20 - 0.04 - 0.02 0.09 0.06 0.02 0.06 0.04 3 0.01 - 2 0.04 0.05 2 - 2 0.04 0.01 0.07 0 1 0.02 5 5 2 3 2 3 0.04 0.01 2 2 0 2 0.04 - 0.03 1 1 - 0.01 0.06 - H 0.07 - 0.01 0.07 - H 0.06 0.01 0.06 - 3 - 2 0.09 2 5 2 0.01 2 0 2 0 2 0.11 0 5 1 - 4 4 4 - 4 2 - 3 2 2 0 0 0 - - 3 2 2 0 2 0 2 0 2 1 0 0 10 Kim et al.
0.2 0.2 0.2 0.2
0.1 0.1 0.1 0.1 ) ) ) I B V - - - V V B 0.0 0.0 U 0.0 0.0 ( ( (
0.1 0.1 0.1 0.1
V = 0.007 ± 0.056 (B-V) = 0.000 ± 0.021 (U-B) = -0.012 ± 0.063 (V-I) = -0.026 ± 0.018 0.2 0.2 0.2 0.2 10 12 14 0.25 0.50 0.75 1.00 0.5 0.0 0.5 0.20 0.25 0.30 0.35 0.40 V B-V U-B V-I
Figure 3. Comparison of our photometric data and photoelectric photometric data by Guetter & Turner (1997). The thick line and shaded regions represent the mean and standard deviation, respectively. The mean difference and the standard deviation between the two sets of photometry are displayed in the lower-left corner of each panel.
131 2.1.3. Data Reduction and Comparison
1 132 Pre-processing was performed using the IRAF /CCDRED package to remove the instrumental
133 artifacts. The brightness of sources were measured applying a point-spread function (PSF) fitting
134 method to the target images with IRAF/DAOPHOT (Stetson 1987). A total of 731 stars and
00 135 1473 stars were detected from the Maidanak AZT-22 observations and the Kuiper 61 telescope
136 observations, respectively. These data were transformed to the SAAO standard system as described
137 in Lim et al.(2009) and Lim et al.(2015a). The photometric data of stars in common between these
138 datasets were obtained from their weighted mean values, where the inverse of squared photometric
139 errors were adopted as the weight. We present the photometric uncertainties in Figure5.
140 Figure2 shows the difference between the two datasets. These two photometric datasets show
141 a very good consistency. The mean difference is less than 0.01 mag for V , V − I, and B − V ,
142 and about 0.025 mag for U − B. Figure3 shows the difference between our optical data and the
143 photoelectric data of Guetter & Turner(1997). We excluded known variables and emission-line stars
144 in the comparison. The differences are less than 0.01 mag for V,B − V and about 0.01 mag for
145 U − B. The difference for V − I is about 0.03 mag, which is consistent with the expected error
1 IRAF is developed and distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy under a cooperative agreement with the National Science Foundation. IC 1590 11
0.2 0.2
0.1 0.1 ) I - V
0.0 V 0.0 (
0.1 0.1
V = 0.033 ± 0.075 (V-I) = -0.001 ± 0.043 0.2 0.2 14 15 16 17 0.5 1.0 1.5 V V-I
Figure 4. Comparison of our photometric data with the previous CCD data from Sharma et al.(2012). Stars brighter than V = 17 mag were used in the comparison to avoid the large photometric errors of faint stars. The thick line and shaded area represent the mean and standard deviation, respectively. The mean difference and the standard deviation between the two sets of photometry are displayed in the lower-left corner of each panel.
Figure 5. Phtometric uncertainties as a function of magnitude in U, B, V, and I bands for the observed region. The large errors of some stars at a given magnitude may result from the deblending of duplicated sources.
146 in the photoelectric photometry of Guetter & Turner(1997) which is vulnerable due to short-term
147 variations in the sky conditions. Figure4 shows the difference between the CCD photometric data
148 of Sharma et al.(2012) and ours. The mean differences are about 0.03 mag in the V band and less 12 Kim et al.
3.0
2.5
2.0 N g o l 1.5
1.0
0.5
0.0 8 10 12 14 16 18 20 I [mag]
Figure 6. I band luminosity function derived from all observed stars. We adopted a bin size of 0.5 mag. The turnover appears at about 17.5 mag, which is related to the completeness limit of our photometry.
149 than 0.01 mag in V − I. Our data are also well consistent with the previous CCD photometric data
150 of Sharma et al.(2012).
151 The brightest stars, HD 5005A and HD 5005B, were too bright to be measured reliably. They
00 152 were saturated in the Kuiper 61 telescope images. Although HD 5005B was not saturated in the
153 observation with the 2k CCD, the errors were large due to the saturated neighboring star HD 5005A.
154 We used instead the photoelectric photometric data of HD 5005A+B of Guetter & Turner(1997).
155 Figure6 shows the luminosity function of all stars in our survey region. The least-squares method
156 was used to determine the slope of the luminosity function between 13 and 17 mag at which our IC 1590 13
157 photometry is complete. If the linear relation is extrapolated beyond 17 mag, our photometry may
158 be 84% complete at the turnover magnitude.
159 We transformed the CCD coordinates of the observed stars to the equatorial coordinates tied to the
160 Two Micron All Sky Survey (2MASS; Cutri et al. 2003) using several hundreds of stars in common
161 between the two data. A total of 833 2MASS counterparts were found in our data. The near-infrared
162 (NIR) data of some stars are used to investigate the extinction law.
163 2.2. Archival Datasets
164 2.2.1. Gaia Data
165 Since stars in a cluster are considered to have formed from the same molecular cloud on a short
166 timescale, cluster members share almost the same kinematic properties. Astrometric data such
167 as proper motions and parallaxes are useful to isolate members of the cluster from the back-
168 ground/foreground field stars. The Gaia EDR3 catalog (Gaia Collaboration et al. 2020) provides
169 accurate proper motion and parallax data for sources across the whole sky. The nominal uncertainties
−1 −1 170 reach 0.02 mas and 0.03 mas yr for bright sources (G < 15 mag) and 1 mas and 1.3 mas yr for
171 G ∼ 21 mag sources. It has been known that the Gaia EDR3 parallaxes have a global zero-point
172 offset of −0.017 mas (Lindegren et al. 2020). We searched the Gaia EDR3 catalog for counterparts
00 173 of our observed stars, and all stars were matched in a radius of 0. 5. To obtain reliable results, we
174 did not use stars with a negative parallax and compensated for the global zero-point of 0.017 mas.
175 2.2.2. Chandra ACIS Data
−3 176 PMS stars are known to be bright X-ray sources (LX /Lbol ∼ 10 , Feigelson et al. 2003; Linsky et
177 al. 2007). Furthermore, their X-ray activity persists for a long time (Sung et al. 2008). For these
178 reasons, PMS stars in SFRs can be well identified through X-ray imaging observations (Sharma et
179 al. 2012). We took a list of 1185 X-ray sources in this SFR (Townsley et al. 2019). A total of 210
00 180 X-ray counterparts were found within a searching radius of 0. 7.
181 3. MEMBER SELECTION 14 Kim et al.
182 As the majority of open clusters are distributed in the Galactic plane, a large number of field
183 interlopers are observed along the line of sight in the same FoV. Member selection is therefore of
184 crucial importance in the study of open clusters.
185 3.1. Low-Mass Stars
186 Material in the circumstellar disk of YSOs may accrete onto the stellar surface, manifesting in a UV
187 excess and Hα emission (Calvet & Gullbring 1998; Hartmann 1999). Therefore, Hα emission can be
V +I 188 a useful indicator of CTTS in young open clusters (Sung et al. 1997). The Hα index (≡ Hα − 2 )
189 in Figure7(a) is a measure of H α emission introduced by Sung & Bessell(2000). Figure7(a)
190 shows the Hα indices of the observed stars with respect to the (V − I) color. The fiducial line
191 for Hα emission stars is adopted from Lim et al.(2015a). The selection criteria were set based
192 on the distribution of ∆Hα, where ∆Hα means the difference between the measured Hα indices
193 and the fiducial line at a given (V − I) color. The ∆Hα distribution appears close to a Gaussian
194 distribution with a weak skewed wing toward negative ∆Hα values. The standard deviation of
195 the Gaussian distribution is about 0.15 mag. Normal stars without Hα emission lines were found
196 within the standard deviation from ∆Hα = 0, while Hα emitting stars may be detected from the
197 skewed wing. Therefore, we selected Hα emission candidates as having ∆Hα values smaller than
198 a standard deviation (∆Hα < −0.15) and Hα emission stars as having ∆Hα values smaller than
199 twice the standard deviation (∆Hα < −0.30) from the zero value. Some stars satisfying the criteria
200 were excluded because of their large photometric errors. A total of 39 Hα emission stars and 15 Hα
201 emission candidates were identified. We visually inspected these stars in the Hα images. In the same
202 FoV, Sharma et al.(2012) found four H α emission stars, which were high mass or intermediate-mass
203 stars. In this study, we identified a larger number of low-mass stars showing Hα emission.
204 As mentioned above, identifying X-ray emission stars is the most effective way to select low-mass
205 PMS members (Flaccomio et al. 2006; Rauw & Naz´e 2016; Caramazza et al. 2012) because stars in
206 the PMS stage are much brighter in X-ray than main-sequence stars (Prisinzano et al. 2008). The
207 majority of X-ray sources as well as Hα emission stars are found in a specific locus (hereafter PMS
208 locus) above the field main-sequence stars at given (V − I) colors (see Figure7(c)). Most X-ray IC 1590 15
2.5 (a) (c) H emission 2.0 8 H candidates X-ray emission 1.5 proper motion & plx x e d n
i 1.0 10
H 0.5
0.0 12
0.5 0.0 0.5 1.0 1.5 2.0 2.5 3.0 3.5 4.0 V-I I 14 (b)
0
] 16 r y / s a m [
2 18
4 20 4 2 0 0 1 2 3 cos [mas/yr] V-I
Figure 7. Member selection criteria for IC 1590. The small dots, red star symbols, magenta star symbols, blue squares, and green triangles represent all the detected stars, Hα emission stars, Hα emission star candi- dates, X-ray emission stars, and member candidates selected from the Gaia EDR3 (Gaia Collaboration et al. 2020), respectively. (a) Hα index versus (V − I) two-color diagram. The solid fiducial line is adopted from Lim et al.(2015a). The dashed and dot-dashed lines represent the fiducial of H α emission candidates and Hα emission stars. (b) Proper motion map of IC 1590. X-ray emission stars and Hα emission stars/candidates −1 are concentrated around (µα cos δ, µδ) = (-1.3, -1.5) mas yr . (c) Color-magnitude diagram of IC 1590. The solid line represents the zero-age main-sequence relation (Sung et al. 2013). We set a PMS locus between two thick dashed lines where the majority of PMS members are distributed. The reddening vector with
AV = 1 is shown as an arrow which is nearly parallel to the PMS locus. 16 Kim et al.
20
15 N 10
5
0 0.20 0.25 0.30 0.35 0.40 0.45 0.50 0.55 [mas]
Figure 8. Distribution of parallax of X-ray emission stars (blue) and members (green) selected in Section3. The solid lines are fit to Gaussian functions using the maximum likelihood method. Stars with parallaxes larger than five times the errors are used in the histogram. We corrected for the global zero-point offset of 0.017 mas (Lindegren et al. 2020).
209 emission stars in the PMS locus do not show any appreciable Hα emission, which implies that they
210 are YSOs without an active accretion disk, at least at the time of the Hα observations. On the other
211 hand, some X-ray emission stars are found below the PMS locus. These stars may be PMS stars
212 going through disk-bearing, or field late-type stars undergoing prolonged X-ray activity due to them
213 being in low-mass close binary systems (Sung et al. 2008; Fang et al. 2021).
214 We also used proper motions of stars to select additional members. Figure7(b) shows the proper
215 motion of stars in the observed FoV. We fit a 2-D Gaussian function, with the principal axis anal- IC 1590 17
216 ysis, to the proper motion distribution of X-ray emission stars and Hα emission stars that were
−1 217 concentrated around the mean value of (µα cos δ, µδ) = (−1.3, −1.5) mas yr with the principal axis
◦ 218 inclined by 19.4 . The standard deviations of the major and minor axes were computed to be 0.2
219 and 0.1, respectively, yielding an elliptical shape that encompasses as many members as possible.
220 Stars with proper motions within three times the standard deviations from the mean values and
−1 221 associated errors smaller than 0.5 mas yr were considered as member candidates (see Figure7(b)).
222 The Gaia parallaxes of stars (Gaia Collaboration et al. 2020) provide crucial constraints on cluster
223 membership. Figure8 displays the parallax distribution of stars. The mean and standard deviation
224 of X-ray emission stars are 0.36 ± 0.03 mas. Stars with parallaxes between 0.30 mas and 0.45 were
225 selected as probable members, where we only considered parallaxes larger than three times the errors.
226 In addition, we considered 44 stars in the PMS locus as PMS star candidates. They were not
227 detected by the Chandra observation (Townsley et al. 2019) and do not show the Hα emission. Their
228 Gaia parallaxes and proper motions are similar to those of the cluster although the associated errors
229 are large due to their faintness. However, we considered these stars as member candidates in this
230 study. The later release of the Gaia data will be able to judge their membership
231 Other properties of YSOs are UV excess. The UV excess originates from hot spots in the photo-
232 sphere due to the release of gravitational energy from accreting material (Calvet & Gullbring 1998).
233 This phenomenon that occurs in actively accreting YSOs is normally accompanied by strong Hα
234 emission. Our optical data are rather limited to detect such UV excess stars because our faint YSOs
235 have large photometric errors in the U band. However, two Hα emission stars definitely show UV
236 excesses in Figure9. In addition, some faint stars, although they have large photometric errors, are
237 bluer in (U − B).
238 Stars satisfying more than two criteria among X-ray emission, Hα emission, proper motion, parallax,
239 and PMS locus are selected as the members of IC 1590. If stars fulfill only one criterion, the stars
240 were considered as member candidates. We summarize the criteria of member selection as below:
241 1. X-ray emission stars 18 Kim et al.
242 2.H α emission stars or candidates
243 3. proper motion within three times the standard deviations from the mean value (µα cos δ, µδ) =
−1 244 (−1.3, −1.5) mas yr
245 4. parallaxes between 0.03 mas and 0.45 mas
246 5. stars in the PMS locus in HRD
247 Finally, we selected 369 of the low-mass stars as members.
248 3.2. Massive and Intermediate-mass Stars
249 Massive O- and B-type stars are classified as probable cluster members because their lifetime is so
250 short that they cannot migrate far away from their birth place. Such early-type stars, being luminous
251 and blue in color, can be easily identified in photometric diagrams. But selection of late B- to F-type
252 members of young open clusters is very difficult because Hα photometry is less effective to identify
253 them (except for Herbig Ae/Be stars), and these stars are quiet in the X-ray (Sung et al. 2017).
254 However, high precision proper motion and parallax data from the Gaia mission make it possible
255 to identify these intermediate-mass members. In Figure9, stars with B − V < 0.3 and V < 17
256 mag corresponds to late-B-type stars, and therefore we considered these stars as early-type member
257 candidates. Their membership was assessed by using proper motions and parallaxes. To do this, we
258 adopted the same criteria 3 and 4 used for the selection of low-mass members. A total of 39 massive
259 and intermediate-mass stars were selected as members.
260 Figure 10 shows the CMD of IC 1590. Grey and black dots represent, respectively, field stars and
261 cluster members. In the (V,V − I) CMD, PMS star population is clearly seen. On the other hand,
262 some PMS stars have colors bluer than those of the zero-age main sequence (ZAMS) relations in the
263 (V,B − V ) and (V,U − B) diagrams, which implies a UV excess for these PMS stars, although they
264 have large photometric errors. Indeed, three of these stars were identified as Hα emission stars.
265 4. REDDENING AND DISTANCE
266 4.1. The Extinction Law IC 1590 19
1.5 X-ray emission H emission 1.0 No.409 H candidates proper motion & plx 0.5
0.0 U-B No.1254 0.5
No.1210 1.0
1.5
0.5 0.0 0.5 1.0 1.5 2.0 B-V
Figure 9. (U-B, B-V) two-color diagram. The small open grey circles represent all detected stars and the
black big dots represent stars with small errors U−B < 0.07 and B−V < 0.07 where U−B and B−V mean the errors of (U − B) and (B − V ), respectively. The solid and dashed lines represent intrinsic color-color relation from Sung et al.(2013) and the relation reddened by E(B − V ) = 0.35, respectively. The other symbols are the same as in Figure7. Some notable sources among early-type stars (No.409, No.1210, and No.1254) are discussed in Section 6.4.
267 The interstellar extinction can be determined in the (U − B, B − V ) two-color diagram (TCD)
268 shown in Figure9. The reddening vector is well established as follows (Sung et al. 2013);
E(U − B)/E(B − V ) = 0.72 + 0.025 × E(B − V ) (1)
269 We obtained the reddening of individual early-type members by comparing their observed with
270 intrinsic colors (Sung et al. 2013) along this reddening vector. 20 Kim et al.
8
10
12
14 V 16
18
20
22
0 1 2 3 0 1 2 1 0 1 V-I B-V U-B
Figure 10. Color-magnitude diagrams of IC 1590. Left panel : V − I vs. V diagram. Middle panel : B − V vs. V diagram. Right panel : U − B vs. V diagram. The small grey dots represent all observed stars and the big black dots represent members of IC 1590. The solid lines represent the zero-age main sequence relations reddened by average reddening of members. A distance modulus of 12.3 mag was applied to the ZAMS relation (see the main text in Section 4.2 for detail).
271 The mean reddening is hE(B−V )i = 0.40±0.06 (s.d.). There is non-negligible differential reddening
272 given the standard deviation larger than the typical photometric errors of 0.017 mag at V < 20 mag.
273 Figure 11 shows the spatial variation of E(B−V ). The southern region is more obscured by remaining
274 material than the northern region. The presence of the dense gas toward NGC 281 West was also
275 confirmed from the CO observations (Elmegreen & Lada 1978; Megeath & Wilson 1997; Lee & Jung
276 2003). The reddening of individual PMS stars could not be obtained from the photometric diagrams.
277 Therefore, we used this map to correct the reddening of PMS stars.
278 Light from stars in very young open clusters is absorbed by foreground and intracluster media (see
279 Hur et al. 2012; Lim et al. 2015a). The properties of the line of sight extinction in given passbands
280 are, in general, described by the total-to-selective extinction ratio (RV ), which depends on the size IC 1590 21
281 distribution of dust grains (Draine 2003). The canonical value of the total-to-selective extinction
282 ratio (RV = 3.1) of the diffuse ISM is generally accepted by many investigators (Wegner 1993; Lada
283 & Lada 1995; Winkler 1997). However, a somewhat lower value of RV = 2.9 was reported toward
284 the Perseus arm (Sung & Bessell 2014). In the case of intracluster media, RV,cl can be high because
285 some SFR are so young that large dust grains may still remain without having been destroyed by UV
286 photons from massive stars. Abnormal high extinction laws have been reported in some young open
287 clusters (e.g., Hur et al. 2012; Lim et al. 2015a; Fang et al. 2021), while others have not (e.g., Sung
288 et al. 2017). The interstellar extinction due to the dust associated with IC 1590 was claimed to be
289 larger from RV (Guetter & Turner 1997; Sharma et al. 2012). We corrected for the total extinction
290 AV of each star as follows;
AV =AV,fg + AV,cl (2)
=E(B − V )fg × RV,fg + E(B − V )cl × RV,cl. (3)
291 where AV,fg, E(B − V )fg, RV,fg, AV,cl, E(B − V )cl, and RV,cl represent the foreground extinction,
292 foreground reddening, foreground RV , intracluster extinction, intracluster reddening, and intracluster
293 RV , respectively.
294 To check whether the extinction law of IC 1590 is abnormal or not, the properties of extinction were
295 investigated using color excess ratios. We estimated the intrinsic colors of early-type stars in the NIR
296 wavelength passbands interpolating their reddening-corrected B − V to the intrinsic color relations
297 between (B − V ) and (V − λ), where λ is the I, J, H, or Ks magnitude (Sung et al. 2013). The
298 color excesses E(V − λ) were calculated by comparing the observed colors with the intrinsic colors.
299 Figure 12 shows the color excess ratios of various passbands with respect to E(B − V ). RV,cl was
300 obtained from the color excess ratio of E(V − λ) to E(B − V ). In some regions, RV from the I band
301 showed a considerably different value than those from the NIR JHKs bands (e.g., Hur et al. 2015),
302 so we did not use the I band for obtaining RV . There are several stars with an IR excess emission
303 (see Figure 12). Their IR excess may originate from circumstellar material. Hence, we cautiously 22 Kim et al.
6
0.40 0.50
4 0.35
0.45 ] n
i 2 0.40 m c r a [
0.40 .
0 E(B-V) c e D
0.40 2 0.35
0.45
4 0.45 0.30 6 4 2 0 2 4 R.A. [arcmin]
Figure 11. Interstellar extinction map. The early-type stars used for deriving the extinction map are shown as large dots. Each color of dots represents the value of E(B − V ). Magenta diamond symbols represent foreground stars. E(B − V ) of a given position in the extinction map is obtained from the weighted average of all early-type stars and the weight is adopted as exp (− ri ) where r is the distance from the i-th early-type r0 i
member and the scaling length r0 is adopted as 1 arcmin.
304 excluded stars with E(V − λ) larger than 2.5 times the standard deviation from the rest. The mean
305 RV,cl obtained from the NIR color excess is 3.94 ± 0.34 (s.d.).
306 In our FoV, there are five foreground early-type stars. Their mean reddening is E(B − V )fg =
307 0.24 ± 0.02 (s.d.). From these stars, we can examine the foreground extinction law in the same way.
308 The RV,fg is 2.94 ± 0.10 (s.d.), which is normal value in the direction of the Perseus Arm (Sung &
309 Bessell 2014). IC 1590 23
1.5 2.5 3 RV, fg = 3.03 ± 0.40 RV, fg = 2.92 ± 0.17 RV, fg = 2.94 ± 0.20 RV, fg = 2.95 ± 0.18 R = 3.98 ± 0.55 3 V, cl 2.0 RV, cl = 3.77 ± 0.63 RV, cl = 4.07 ± 0.82 RV, cl = 3.99 ± 0.92 1.0 2 )
1.5 s
K 2
1.0 V E(V-I) E(V-J) ( 0.5 E(V-H)
1 E 1 0.5 0.0 0.0 0 0 0.0 0.2 0.4 0.6 0.8 0.0 0.2 0.4 0.6 0.8 0.0 0.2 0.4 0.6 0.8 0.0 0.2 0.4 0.6 0.8 E(B-V) E(B-V) E(B-V) E(B-V)
Figure 12. Color excess ratio diagrams. Black, cyan, and red colors represent, respectively, the early-type stars used in the reddening estimation, foreground stars, and NIR excess stars. The black solid and dashed
lines represent the adopted color excess ratios for RV,fg = 2.94 and RV,cl = 3.94, respectively. The blue thin solid line shows the relation between E(B − V ) and E(V − λ) for each band.
6 6 6 6
8 8 8 8
10 10 s 10 10 s I J K K V V V V Q Q Q 12 12 12 Q 12
14 14 14 14
16 16 16 16 1.00 0.75 0.50 0.25 0.00 1.00 0.75 0.50 0.25 0.00 1.00 0.75 0.50 0.25 0.00 1.00 0.75 0.50 0.25 0.00 Q' Q' Q' Q'
Figure 13. ZAMS fitting to the early-type main-sequence stars of IC 1590. The ZAMS relations of Sung et al.(2013) were fit to the lower ridge line of the members. The solid line and shaded region represent the ZAMS relation with the distance modulus of 12.3 ± 0.2 mag. The color of symbols are the same as in Figure 12.
310 4.2. Distance
311 The colors and magnitudes of young main-sequence stars are empirically well established (Sung
312 et al. 2013). We determined the distance of IC 1590 by comparing the CMD of the early-type
313 main sequence members with the ZAMS relations of Sung et al.(2013). Johnson & Morgan(1953)
314 introduced a reddening-free Q value, Q = (U − B) − RV × E(B − V ), and investigated the Q value
315 for O- and B-type stars (e.g., Johnson & Morgan 1953; van den Bergh & Hagen 1968). To minimize
316 the effect of reddening, we used reddening-free quantities and a modified Johnson Q in the distance 24 Kim et al.
317 determination. These quantities are defined as follows (Sung et al. 2013).
QVI =V − 2.45 × (V − I) (4)
QVJ =V − 1.33 × (V − J) (5)
QVH =V − 1.17 × (V − H) (6)
QVKs =V − 1.10 × (V − Ks) (7)
Q0 =(U − B) − 0.72 × (B − V ) − 0.025 × E(B − V )2 (8)
318 Figure 13 shows the reddening-free CMDs of early-type stars. Because a large fraction of high mass
319 stars constitute binary systems (Sana et al. 2012) as well as the effect of rapid evolution, high mass
320 stars may be brighter than the ZAMS at a given effective temperature (Lim et al. 2015a). Therefore,
321 we fit the ZAMS to the lower ridgeline of the mid- to late-B type main-sequence stars within 0.2
322 mag. The estimated distance modulus is 12.3 ± 0.2 mag, which is equivalent to the distance of
323 d = 2.88 ± 0.28 kpc.
324 We also determined a distance of IC 1590 from the inversion of the parallaxes of members. Stars
325 with a parallax larger than five times the error were used to determine the distance. We fit a Gaussian
326 distribution to the parallax distribution of members (Figure8) and obtained $ = 0.37 ± 0.03 mas
+0.24 327 (s.d.). As a result, the distance is computed to be 2.70−0.20 kpc. The distance from the parallaxes
328 is well consistent with that from the ZAMS fitting. Our result is also in a good agreement with the
329 distance d = 2.82 ± 0.24 kpc from radio VLBA astrometry of an H2O maser in NGC 281 West (Sato
330 et al. 2008).
331 5. AGE AND IMF
332 5.1. Hertzsprung-Russell Diagram
333 To construct the Hertzsprung-Russell Diagram (HRD), we corrected the reddening and total ex-
334 tinction in visual band, according to Equation3. While the reddening values of individual early-type
335 members were applied to their observed colors, the values inferred from the reddening map were
336 adopted for the reddening correction of lower-mass stars. We also applied the distance modulus of IC 1590 25
337 12.3 mag to the extinction-corrected visual magnitude. The effective temperature (Teff ) of O-type
338 stars was inferred from the spectral type-Teff relation, while that of the other stars was estimated by
339 interpolating their intrinsic colors to the color-Teff relations. The bolometric correction values were
340 obtained from the BC-Teff relation. These relations are obtained from Sung et al.(2013).
341 We present the HRD of IC 1590 in Figure 14 with several stellar evolutionary tracks (Ekstr¨omet
342 al. 2012) and PMS evolution models (Siess et al. 2000). Stars with masses larger than 7 M are in
343 the main-sequence stage. The age of this cluster inferred from the main-sequence turn-off is about
344 1.9 Myr. However, the stars in the HD 5005AB system were not spatially resolved, its luminosity
345 may therefore be overestimated. Hence, the age of these stars may be younger than 1.9 Myr. Most
346 of the PMS members are evolving along Hayashi tracks or in the Kelvin-Helmholz contraction phase.
347 The upper mass range of these PMS members appears to be 5 M and the majority of PMS stars
348 are younger than 5 Myr.
349 The masses and ages of individual members were derived by interpolating the Teff and Mbol of each
350 star to those from the evolutionary tracks for the main-sequence (Ekstr¨omet al. 2012) and PMS (Siess
351 et al. 2000). We present the age distribution of PMS stars with m < 2 M in Figure 15. The age
352 distribution shows a large spread because the ages of PMS stars are almost certainly overestimated
353 due to imperfect correction for differential reddening as well as internal extinction of disk-bearing
354 stars, variability, and photometric errors for faint stars (Sung et al. 1997, 2004). We fit a skewed
355 Gaussian function to the age distribution. The representative age of this cluster was estimated to be
356 0.8 Myr from the mode value of the age distribution. The ages of PMS stars ranges from 0.7 Myr to
357 8.4 Myr at 10% and 90% of the cumulative distribution.
358 5.2. Initial Mass Function
359 The mass function can be derived from the distribution of masses of individual stars. The IMF
360 is expressed as the number of stars in a given logarithmic mass bin in a projected area S, i.e.
N 361 ξ ≡ ∆logm·S , where N and ∆log m represent the number of stars within a given mass bin and the
362 size of the logarithmic mass bin, respectively. The IMF is approximated as a linear relation of stellar 26 Kim et al.
1.9 Myr 60 M 10.0 40 M
25 M
7.5 15 M
9 M
5.0 7 M
2.5 l
o 5.0 M b
M 3.0 M 2.0 M 0.0 1.0 M
2.5 0.5 M
5.0 1 Myr
5 Myr 7.5 10 Myr 20 Myr
4.75 4.50 4.25 4.00 3.75 3.50 log Teff
Figure 14. Hertzsprung-Russell diagram of IC 1590. The solid line in the upper part of the HRD is the ZAMS from Ekstr¨omet al.(2012). Some evolutionary tracks for post-main sequence (Ekstr¨omet al. 2012) and PMS stars (Siess et al. 2000) are plotted by green dot-dashed lines and red dahsed lines, respectively. Isochrones for four different ages interpolated from these two evolutionary models are also superimposed on the HRD (blue solid lines). IC 1590 27
50
40
30 N
20
10
0 0 2 4 6 8 10 12 14 age [Myr]
Figure 15. Age distribution of PMS stars with m < 2 M . Black solid line represents the fit skewed Gaussian function. The mode age is 0.8 Myr.
363 mass in a logarithmic scale, i.e. log ξ ∝ Γ × log m. Salpeter(1955) obtained the slope Γ = −1.35 for
364 the stars in the Solar neighbourhood.
365 We derived the IMF in the range from 0.2 M to 60 M . Different bin sizes were used to include
366 at least one star in a given mass bin; ∆ log m = 0.1 for −0.7 < log m < 0.4, ∆ log m = 0.2 for
367 0.4 < log m < 0.8, and ∆ log m = 0.4 for log m > 0.8. We did not consider the unresolved binary
368 systems. Figure 16 shows the IMF of IC 1590. Before calculating the slope of the IMF, we need to
369 limit the stars within the completeness of our photometry and member selection. Our photometry
370 has a completeness limit of 17.5 mag in the I band, which corresponds to about 1 M . Since our 28 Kim et al.
7.0 ] 2 c p k
/ 6.5 M g
o 6.0 l / # [
)
m 5.5 g o l
( 5.0 g o l 4.5
2.0 1.5 1.0 0.5 0.0 0.5 log m [M ]
√ N Figure 16. Mass function of IC 1590. The error of ξ is set as ∆ log m·S . The bin sizes are ∆ log m = 0.1 for log m < 0.4, ∆ log m = 0.2 for 0.4 < log m < 0.8, and ∆ log m = 0.4 for log m > 0.8. The dashed line shows
the best fit, using the maximum likelihood method, to the mass range m > 1M .
371 member selection depends on X-ray data, the completeness mass of the X-ray dataset should be also
372 considered. Townsley et al.(2019) detected X-ray sources complete down to 0 .4 M , and therefore
373 our member selection is likely complete for stars with masses larger than 1 M .
374 We calculated the slope of the IMF for stars with masses larger than 1 M . Using the method of
375 maximum likelihood, the slope of the IMF, Γ = −1.49 ± 0.14, was determined. It is slightly steeper
376 than the slope of −1.35 derived by Salpeter(1955) and Kroupa(2001) for the Solar neighborhood.
377 This implies that low-mass star formation in this SFR is more dominant than that in the Solar IC 1590 29
378 neighborhood. On the other hand, Guetter & Turner(1997) and Sharma et al.(2012) obtained
379 shallow IMFs. They identified a low number of low-mass members and computed the slope of the
380 IMF for stars with masses larger than 2 M . Therefore, the discrepancy between their results and
381 ours presumably originates from the completeness of the low-mass member selection.
382 6. DISCUSSION
383 6.1. The Effect of binarity in IMF
384 The effect of binaries has not been considered in the derivation of the stellar IMF (Kroupa 2001).
385 However, it is a common phenomenon that most stars form or end up in binary or multiple systems
386 during the star formation process. A number of previous studies have investigated stellar multiplicity
387 and mass ratio (Duquennoy & Mayor 1991; Fischer & Marcy 1992; Delfosse et al. 2004; Raghavan et
388 al. 2010; Sana et al. 2012; Peter et al. 2012). Duchˆene& Kraus(2013) reviewed the overall properties
389 of stellar multiplicity. Here, we examined the effects of binarity on the IMF using a Monte-Carlo
390 technique.
391 We constructed a synthetic cluster containing 10000 stars in the mass range of 0.1 M to 100 M
392 using a Monte-Carlo method. The Kroupa IMF was adopted as the underlying IMF. The binary
393 fraction was determined as a function of primary mass. The mass ratio distribution in a given mass
γ 394 bin is expressed as a power law relation f(q) ∝ q where q is the mass ratio. The functions for binary
395 fraction and mass ratio were adopted from those of Duchˆene& Kraus(2013). The typical binary
396 fraction of low-mass stars and γ are about 20% and 4.2, respectively. More than 80% of massive
397 stars occur in binary systems (Sana et al. 2012), and their mass ratio distribution appears to be flat.
398 A total of 2730 binary systems and 7270 single stars were generated. The effective temperature and
399 bolometric magnitude of stars in the synthetic cluster were obtained by interpolating their masses
400 to the evolutionary models for MS and PMS stars (Ekstr¨omet al. 2012; Siess et al. 2000). All the
401 binary systems were assumed to be unresolved binaries. The temperatures of the binary systems
402 may be close to those of the primary stars, while their total luminosities can be obtained from the
403 sum of luminosities of the primary and secondary stars. In turn, we derived the mass of stars from 30 Kim et al.
total 5 observation
4 ) m
g 3 o l ( g o
l 2
1
0 2.0 1.5 1.0 0.5 0.0 0.5 log m [M ]
Figure 17. Initial mass functions of the synthetic cluster. Black dots represent the underlying IMF using the resolved binary stars. Red dots represent the IMF for the synthetic cluster when binary systems are
considered as single stars. Lines show the slope of the IMF, determined in the mass range m > 1 M .
404 the HRD using the same method described in Section 5.2. The IMF of the synthetic cluster was
405 finally derived as shown in Figure 17 (red dots). We also present the underlying IMF taking into
406 account the resolved binary stars (black dots) for comparison. The slopes of the underlying IMF
407 and simulated one are −1.25 ± 0.04 and −1.30 ± 0.04, respectively. There is no significant difference
408 between two IMFs. Therefore, the IMF of IC 1590 may not be significantly affected by the binary
409 population in the cluster. IC 1590 31
410 We can guess some reasons that unresolved binaries do not affect the estimation of the IMF very
411 much. The primary stars essentially dominate the total luminosities of binary systems. In the high-
412 mass regime, the size of the mass bins is sufficiently large to cover a number of stars even if stellar
413 masses were misderived due to the contribution of the secondary stars. The low-mass PMS stars are
414 evolving along the Hayashi tracks, and therefore their masses may not significantly misderived even
415 if their luminosities are increased by the binary effects.
416 6.2. Mass Segregation
417 Some open clusters exhibit the signature of mass segregation (e.g., the concentration of massive
418 stars; Sung et al. 2013, 2017; Dib et al. 2018; Dib & Henning 2019; Hetem & Gregorio-Hetem 2019).
419 Mass segregation is the result of energy equipartition among cluster members (Binney & Tremaine
420 1987). This dynamical process takes place on a relaxation time scale which for a typical cluster with
421 about 1000 members is longer than ten crossing times (Bonnell & Davies 1998). However, because
422 young open clusters are too young to be relaxed, the cause of mass segregation among young open
423 clusters is still debated. One possibility is that the mass segregation found in several young open
424 clusters could have a primordial origin related to star formation in inhomogeneous molecular clouds
425 (Bonnell & Davies 1998). On the other hand, some assert that the origin of mass segregation in
426 young open clusters is due to rapid dynamical evolution in their early stages (McMillan et al. 2007;
427 Allison et al. 2009, 2010). Such rapid dynamical evolution occurs when substructures in molecular
428 clouds merge together into a star cluster. In this situation, mass segregation can dynamically occur
429 on a shorter timescale than the relaxation time.
430 To investigate the signature of mass segregation in IC 1590, we derived the mass functions (MFs)
0 431 of stars in the inner and outer regions from the half-number radius of 2. 5, respectively. We derived
432 IMFs of stars in the same manner described in Section 5.2. The number of stars is 206 for the inner
433 region and 202 for the outer region. Six out of seven massive stars including HD 5005AB reside in
434 the inner region. Because there is only one star with m > 10 M in the outer region, we determined
435 the slope of the MFs using stars with 1 M < m < 10 M for a fair comparison. The slopes of the 32 Kim et al.
7.5
] 7.0 2 c p k / 6.5 M g o l / 6.0 # [
)
m 5.5 g o l ( 5.0 g o l 4.5
2.0 1.5 1.0 0.5 0.0 0.5 log m [M ]
Figure 18. Mass functions for the inner (blue) and outer (red) regions from the half-number radius of 2.05.
The slopes are determined in the mass range 1 < m/M < 10 and −1.55 ± 0.27 for the inner region and −2.31 ± 0.21 for the outer region.
436 MFs are somewhat different; Γ = −1.55 ± 0.27 for the inner region and Γ = −2.31 ± 0.21 for the
437 outer region, respectively (Figure 18). This result indicates that mass segregation occurs in IC 1590.
438 Were the mass segregation caused by rapid dynamical evolution at an early stage, there would be a
439 signature in the stellar mass and velocity correlation. Adopting a distance of 2.8 kpc, we computed
p 2 2 440 two-dimensional tangential velocities (Vt = VR.A. + VDec.) of members using the proper motions
441 from the Gaia EDR3 (Gaia Collaboration et al. 2020). Figure 19 displays the Vt distribution with
442 respect to the radial distances from the cluster center and stellar masses. We could not find any
443 mass-dependent correlation. The dispersion in Vt increases with stellar mass, but the large errors IC 1590 33
0.6
) 0.4 1 s
m 0.2 k / t V (
g 0.0 o l
0.2
0.4 1.0 0.5 0.0 0.5 log(r/0)
0.6
) 0.4 1 s
m 0.2 k / t
V Vt = 0.09 0.09 0.10 0.18 0.23 (
g 0.0 o l
0.2
0.4 1.5 1.0 0.5 0.0 0.5 log(m/M )
Figure 19. Tangential velocity Vt distribution of cluster members with respect to the radial distance (upper) and stellar mass (lower). In the upper panel, the size of dots is proportional to the mass of individual stars.
The dashed lines represent the median Vt. In the lower panel, the average and standard deviation of Vt in several mass bins are represented by red dots and error bars, respectively. It is worth noting that the typical
uncertainties of tangential velocity Vt in a given mass bin are large in low-mass part.
444 in the proper motions of faint stars may be responsible for the large scatter in Vt. Hence, these
445 results suggest that the mass segregation in IC 1590 more likely has a primordial origin, rather than
446 resulting from the violent dynamical evolution of the cluster.
447 6.3. Triggered or Sequential Star Formation 34 Kim et al.
Figure 20. Optical image of IC 1590 and two SFRs taken from the Digital Sky Survey 2. The green rectangular and red ellipse show the FoV of Kuiper 6100 telescope observation and the northern subcluster of NGC 281 West.
448 Figure 20 displays two SFRs in NGC 281 region. Based on CO observations and the age sequence of
449 YSOs, Megeath & Wilson(1997) and Sharma et al.(2012) suggested that the star formation of NGC
450 281 West and NGC 281 East were affected by IC 1590. To investigate the possibility of triggered
451 star formation of NGC 281 West, we estimated the expanding velocity of the HII region NGC 281.
452 If we assume that the Lyman continuum photons from O-type stars are used only for expanding the
453 HII region, the expansion velocity is given by
NLyC v = 2 (9) 4πr nH IC 1590 35
454 where NLyC and nH are the number of the Lyman continuum photons per second and the number
455 density of hydrogen atoms. Lim et al.(2018) derived an expansion velocity for the NGC 1893 HII
−1 456 bubble of 8.6 km s . NGC 1893 contains the same number of O-type stars and has a HII bubble
457 with the similar diameter to NGC 281. We compared the NLyC from O-type stars and the cluster size
458 of IC 1590 with those of NGC 1893 where NLyC was obtained from Draine(2011). The expanding
459 velocity of NGC 281 is comparable to that of NGC 1893 when the hydrogen number densities nH of
460 the two regions are the same.
461 The projected distance between the O-type Trapezium system of IC 1590 and the northern sub-
462 cluster of NGC 281 West is about 4.1 pc at d = 2.8 kpc (Figure 20). The travel time of the ionization
−1 463 front to the northern subcluster of NGC 281 West is 0.5 Myr at the expanding velocity of 8.6 km s .
464 If IC 1590 has indeed triggered the star formation of NGC 281 West, the age of NGC 281 West
465 would be 0.3 Myr. Then most of YSOs in NGC 281 West would be at Class 0/I stage if we consider
466 that the lifetime of Class 0/I YSOs is about 0.55 Myr (Evans et al. 2009). However, Sharma et al.
467 (2012) detect 13 Class II YSOs in the northern subcluster of NGC 281 West. Hence, it is not clear
468 whether the HII region driven by massive stars in IC 1590 has triggered the formation of the northern
469 subcluster in NGC 281 West.
470 Megeath et al.(2002, 2003) also suggest that the NGC 281 complex was formed in a fragmenting
471 superbubble. Based on VLBA parallax and proper motion data of the IRAS sources, Sato et al.
472 (2008) suggested that the 3-D structure of the superbubble is a slightly elongated shape along the
473 Galactic disk. They proposed that the first-generation of OB stars in IC 1590, which were formed
474 by the superbubble shock, ionized the surrounding materials and could have influenced the next-
475 generation of star formation in the neighboring clouds NGC 281 West. If IC 1590 did not trigger
476 the star formation of NGC 281 West, then the star formation of NGC 281 West might have been
477 sequentially affected by the external stimulus such as the shock of superbubble.
478 6.4. Individual Notable Stars
479 6.4.1. No. 409 36 Kim et al.
480 No. 409 is an early-type star (see the TCD in Figure9). This star has almost the same proper
481 motion and parallax of IC 1590. Also, No. 409 shows strong Hα emission and a NIR excess, which
482 implies the existence of an actively accreting hot disk. Therefore, No. 409 seems to be a Herbig Be
483 star. We excluded No. 409 in constructing the extinction map.
484 6.4.2. No. 1210 and No. 1254
485 There is no definitive evidence for the membership of No. 1254 because its proper motion
486 (µα cos δ, µδ) = (0.197 ± 0.012, −3.221 ± 0.014) largely deviates from the cluster’s value and the
487 star is not an X-ray source nor an Hα emitter. No. 1210 has a similar proper motion and parallax as
488 members but X-ray emission or Hα emission were not detected. These two sources follow a normal
489 interstellar extinction law, but have similar E(B − V ) with other members of IC 1590 in Figure 12.
490 Furthermore, they follow the cluster’s sequence in Figure9 so they are probable members of IC 1590.
491 The normal extinction law for No. 1210 and No. 1254 implies that the outer region of NGC 281
492 West follows a normal interstellar extinction law.
493 6.4.3. HD 5005AB
494 The most massive stars HD 5005AB (=No. 380) in IC 1590 is too bright to be de-convolved. The
495 estimated mass of the system is 47 ± 7 M . Sota et al.(2011) reported the spectral types of HD
496 5005A and HD 5005B as O4V((fc)) and O9.7II-III, respectively. However, it is not consistent with
497 the binary formation scenario and the stellar evolution models for the less massive star HD 5005B
498 to have already evolved away from the main-sequence while the more massive star HD 5005A is still
499 on the main-sequence. Sota et al.(2011) also reported that the luminosity class of HD 5005B from
500 HeII λ4686/HeI λ4713 conflicts with a very weak Si IV line. Therefore, a luminosity class V may be
501 more appropriate for HD 5005B.
502 7. SUMMARY
0 0 503 In this paper, we presented deep UBVI and Hα photometry for the central 9. 7 × 9. 7 area of the
504 young open cluster IC 1590. We identified 39 Hα emission stars and 15 Hα emission candidates IC 1590 37
505 from Hα photometry. We selected 408 members from optical TCDs, CMDs, Hα photometry, X-ray
506 emission, Gaia EDR3 parallax, and proper motion. Our photometry is complete down to 1 M .
507 The mena reddening was estimated to be < E(B −V ) >= 0.40±0.06 (s.d.), and there is differential
508 reddening across the cluster. The extinction law toward IC 1590 was ionvestigated by using the
509 color excess ratios. As a result, we confirmed an abnormal extinction for the intracluster medium
510 (RV,cl = 0.39 ± 0.34), while a normal extinction law was found for the foreground medium toward
511 the Perseus arm. The distance modulus of IC 1590 determined from the reddening-free CMDs is
512 12.3 ± 0.2 mag corresponding to d = 2.88 ± 0.28 kpc. This result is in good agreement with the
513 distance derived from the Gaia EDR3.
514 The ages and masses of members were obtained from stellar evolution models and PMS evolutionary
515 tracks. The representative age obtained from the distribution of PMS stars is 0.8 Myr and the ages
516 of PMS stars distributes from 0.7 Myr (10% percentile) to 8.4 Myr (90% percentile). Our estimation
517 is smaller than the upper limit of 4.4 Myr estimated by Sharma et al.(2012). The IMF of IC 1590
518 was derived, and its slope is Γ = −1.49 ± 0.14. It is steeper than Salpeter/Kroupa IMF (-1.35),
519 implying that low-mass star formation is dominant in this SFR.
520 To find a signature of mass segregation in the IC 1590 region, we compared the IMF for the inner
0 0 521 region (r < 2. 5) with that for the outer region (r > 2. 5). The slope of the IMF for the outer region
522 appears steeper than that for the inner region. In addition, the distribution of massive stars with
523 m > 8 M is concentrated in the inner region. These observations indicate the signature of mass
524 segregation in IC 1590. But, we could not find any evidence for the energy equipartition among
525 members in the correlation between stellar masses and tangential velocities.
526 To investigate the possibility of triggered star formation of NGC 281 West, we estimated an ex-
−1 527 pansion velocity of the HII region NGC 281 of 8.6 km s . If the O-type stars of IC 1590 have
528 triggered the star formation of the northern sub-cluster of NGC 281 West, the age of the northern
529 sub-cluster would be 0.3 Myr. However, since there are a number of Class II YSOs probably older
530 than 0.55 Myr in NGC 281 West, it is not evident that the massive stars of IC 1590 triggered the
531 star formation of NGC 281 West. Also, to study further star formation history of NGC 281 region, 38 Kim et al.
532 another star forming region NGC 281 East (Elmegreen & Lada 1978; Megeath & Wilson 1997) and
533 several molecular clouds nearby NGC 281 (Lee & Jung 2003) should be investigated thoroughly.
ACKNOWLEDGMENTS
The authors thank the anonymous referee for many constructive comments and suggestions. We dedicate this work to our colleague Hwankyung Sung, who passed away during the final stages of
this investigation. This work was supported by the National Research Foundation of Korea (NRF) grant funded by the Korean government (MSIT) (Grant No: NRF-2019R1A2C1009475 and NRF- 2019R1C1C1005224). This work has made use of data from the European Space Agency (ESA)
mission Gaia (http://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Anal- ysis Consortium (DPAC, http://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the
Gaia Multilateral Agreement. The Digitized Sky Surveys were produced at the Space Telescope Sci- ence Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the
UK Schmidt Telescope. The plates were processed into the present compressed digital from with the permission of these institutions. IRAF was distributed by the National Optical Astronomy Obser- vatory, which was managed by the Association of Universities for Research in Astronomy (AURA)
under a cooperative agreement with the National Science Foundation.
534 Facilities: Maidanak:1.5m, SO:Kuiper
535 Software: IRAF
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