Neelam Panwar

M R Samal

A K Pandey

J Jose

W P Chen

D K Ojha

K Ogura

H P Singh

R K Yadav

Low-mass young stellar

Keywords: stars, formation -stars, pre-main-sequence -ISM, globules Hii regions -open cluster, initial mass function; star formation

HAL is

INTRODUCTION

Most recent studies show that stars are formed in clusters or groups and nearly half of the low-mass populations are born in massive young clusters or OB associations (Lada & Lada 2003;Allen et al. 2007). Young clusters are natural laboratories for testing the star formation processes and the stellar evolutionary theory as the cluster members have a wide mass spectrum while sharing the same chemical composition and distance.

However, the presence of very luminous (O-or early B-type) stars in such systems profoundly influences their environment by compressing and sweeping the surrounding material because of their strong ionizing radiation and powerful stellar winds (see e.g., [START_REF] Preibisch | Low-mass young stellar population of the cluster[END_REF]Zavagno et al. 2007;Povich et al. 2009;Chauhan et al. 2009Chauhan et al. , 2011;;Deharveng et al. 2010;Jose et al. 2013;Samal et al. 2014). Ultraviolet (UV) radiation from massive stars not only affects the surrounding molecular cloud and its star formation but also causes photoevaporation of disks around the young low-mass stars in their vicinity (see e.g., Bally et al. 1998;Adams et al. 2004;Balog et al. 2007;[START_REF] Gorti | [END_REF], thereby playing a key role in shaping the fundamental properties of the cluster such as stellar mass function (MF), total star formation efficiency, and evolution of circumstellar disks around the young stars.

IC 1805, also known as Melotte 15 or OCL 352, is a young open cluster associated with the Hii region W4 (the Heart Nebula) which is a part of the W3-W4-W5 complex, and contains dozens of massive OB-type stars (Massey, Johnson & Degioia-Eastwood 1995). The W4 complex evacuated by the combined energetic winds of the OB stars within IC 1805 (Goudis & White 1980;Ninkov et al. 1995), appears to be a void in the atomic hydrogen ⋆ E-mail:neelam 1110@yahoo.co.in c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 3 layer with a width of about 100 pc (Taylor et al. 1999). The W4 complex also contains several cometary-shaped bright-rimmed clouds (BRCs) or pillars, pointing towards the luminous stars of IC 1805 (Sugitani, Fukui & Ogura 1991;Lefloch, Lazareff & Castets 1997), similar to those found in Ogura, Sugitani & Pickles (2002), and Chauhan et al. (2011). Such structures are generally the result of compression and erosion of pre-existing clouds by massive stars (Bisbas et al. 2009;Miao et al. 2009), indicating that stellar feedback is playing a strong role in altering the morphology and physical conditions of the W4 complex. Therefore, the complex has been a subject of numerous investigations for understanding the role of massive stars on the star formation process in the surrounding cloud (e.g., [START_REF] Lada | [END_REF]Taylor et al. 1999;Oey et al. 2005;Koenig et al. 2012;Jose et al. 2016).

IC 1805 is located at a moderate distance of ∼ 2.0 kpc (Straižys et al. 2013), similar to the adjacent Hii regions W3 (2.0 kpc, Hachisuka et al. 2006;Xu et al. 2006) and W5 E (2.1 kpc, Chauhan et al. 2011). It is embedded in a very low extinction cloud of E(B -V ) ∼ 0.8 mag or visual extinction (A V ) ∼ 2.4 mag (Joshi & Sagar 1983;Straižys et al. 2013). Also, according to Guetter & Vrba (1989a), Sagar & Yu (1990), and Hanson & Clayton (1993) the extinction law in the direction of IC 1805 is normal. Hence, it is a good target for studies on star formation and evolution of low-mass stars; in particular, how they are influenced by OB type stars. Though the cluster has been investigated by several authors at optical bands (Joshi & Sagar 1983;Kwon & Lee 1983;Sagar 1987;Guetter & Vrba 1989b;Massey, Johnson & Degioia-Eastwood 1995;Sung & Lee 1995;Ninkov et al. 1995), most of these studies were limited to the high-mass stars only. These observations were not only too shallow to detect the low-mass stellar populations but also cover only a relatively small area of the complex. Though several attempts have been made to identify young stellar objects (YSOs) in the complex (Koenig et al. 2012;Povich et al. 2013;Straižys et al. 2013;Broos et al. 2013), no adequate attention has been paid to the characterization of low-mass young stellar population of the cluster using deep optical photometric data.

Since low-mass stars constitute the majority of the stellar population in a cluster and even in the Galaxy (Kroupa 2002), the low-mass stellar content and its MF are essential to understand the nature of the star formation process and the properties of open clusters.

Young stars can be identified using multiwavelength observations depending on their evolutionary stages and spectral types (e.g. Carpenter et al. 2006;Dahm & Hillenbrand 2007).

Overall, the Spitzer results suggest that though primordial circumstellar disks around solar and sub-solar mass stars can last for up to 10 Myr, disk lifetimes are a factor of 2 shorter for higher mass objects (see Williams & Cieza 2011, references therein).

The infrared (IR) observations are particularly useful to identify pre-main-sequence (PMS) stars exhibiting IR excess emission. Also, with the help of X-ray data, young stars can be identified from their excess X-ray emission due to magnetic reconnection flares similar to those seen on the solar surface, but with X-ray fluxes 2 to 3 orders of magnitude larger in comparison to solar-type stars (Feigelson & Montmerle 1999;Preibisch & Feigelson 2005;Güdel et al. 2007). While near infrared (NIR), mid-infrared (MIR) and X-ray observations are suitable to identify PMS stars, they have limitations in characterization of these stars themselves (e.g., age and mass determinations). This can be done best by optical observations as PMS stars possess little or no circumstellar emission at optical wavelengths.

In the present work, we attempt to investigate the low-mass stellar content of the cluster IC 1805 based on high spatial resolution, deep optical photometric observations taken with the Canada-France-Hawaii Telescope (CHFT) in combination with Spitzer, Two Micron All Sky Survey (2MASS), and X-ray data available for the region. We identify and characterize a sample of low-mass young cluster members with circumstellar disks and evolved diskless candidate members. We examine the physical properties of the cluster (e.g., cluster extent, age, disk evolution and mass-function) and try to understand the star formation process in this region.

OBSERVATIONS AND DATA REDUCTION

Fig. 1 shows the DSS2-R image of the IC 1805 region. The regions marked in the blue, green and red are the areas covered by X-ray, optical and Spitzer-Infrared Array Camera (IRAC) observations, respectively. Multi-band Imaging photometer for Spitzer (MIPS) observations are available for the entire area covered by the IRAC observations. The individual observations and photometric catalogs are described below.

Optical Observations

The pre-processed deep V and R band images of the IC 1805 region were obtained from the CFHT archive. The data were obtained with CFH12K, a mosaic CCD camera on CFHT, on 2002, January 07 (P. I. -Sung H.). With a plate scale of 0.21 arcsec per pixel, CFH12K provides a 42 ′ × 28 ′ field of view (FOV). Three exposures of 150s were taken in V and R bands each and the typical seeing during the observations was ∼ 0 ′′ .9. We obtained the photometry of the stars by using the standard IRAF tasks and DAOP HOT -II software package (Stetson 1987). We used the DAOFIND task in IRAF to extract the point sources, and selected only those sources having S/N 5σ above the background. We Observational Sciences (ARIES), Nainital, India. The CCD camera with a pixel size 13.5 µm and a plate scale 0.54 arcsec per pixel, covers a ∼ 18 ′ × 18 ′ FOV. During the observations the seeing was ∼ 1 ′′ .5. To standardize the observations, the SA 98 field of Landolt (1992) was also observed on the same night at various airmasses. A number of bias and twilight flat frames were also taken during the observations.

The DOT images were reduced using various tasks available in the IRAF software package. Instrumental magnitudes were obtained by using the IRAF/DAOP HOT -II package via PSF fitting. The PSF was obtained for each frame by using several isolated stars in the frames. The atmospheric extinction coefficients and zero-point magnitudes obtained from 1. The instrumental magnitudes were converted to the standard values by using a least-square linear regression procedure outlined by Stetson (1992).

The following transformation equations were used to calibrate the observations:

(V -R) = m 1 (v -r) + c 1 , V = v + m 2 (v -r) + c 2 ,
where v, r are the instrumental magnitudes corrected for the atmospheric extinctions, and V, R are the standard magnitudes; c 1 , c 2 and m 1 , m 2 are zero-point constants and colour coefficients, respectively. The final photometry from the DOT observations is used to calibrate the CFHT photometry. We obtained 27,427 sources with a magnitude uncertainty 0.2 mag in both V and R bands.

SP IT ZER IRAC Observations

The NIR/MIR data from the Spitzer space telescope (Werner et al. 2004) using IRAC (Fazio et al. 2004) 8.0 µm, all the sources were visually examined and ambiguous sources were removed. In addition, we also included point sources manually which were not automatically detected by AP EX and supplied the list of these sources to Apexuserlist pipeline to perform the PRF photometry. We have adopted the zero-points for the conversion between flux densities and magnitudes to be 280.9, 179.7, 115.0 and 64.1 Jy in the 3.6, 4.5, 5.8 and 8.0 µm bands, respectively, as given in the IRAC Data Handbook (Reach et al. 2006). The sources with photometric uncertainties 0.2 mag in each band were considered as good detection. To obtain the final catalogue of the sources detected in all IRAC bands, we made a catalog for each channel from the short and long exposures separately and then looked for the closest match within 1 ′′ .2. The final catalog contains 32,014 sources which are detected at least in two IRAC bands, of which 5288 sources have photometry in all IRAC bands.

SP IT ZER MIPS Observations

The region was observed at MIPS 24 µm in 2005 September (P.I. -J. S. Greaves, Program ID -3234). We downloaded the MIPS post-BCD images from the archive, which were created at the image scale of 2 ′′ .45 per pixel and at a spatial resolution of ∼ 6 ′′ . We supplied a list of coordinates of visually identified sources to the Apexuserlist module and performed the PRF fitting to extract the fluxes of the sources. The zero point value 7.14 Jy (adopted from MIPS Data Handbook2 ) is used to convert the flux densities to magnitudes. We cross matched the MIPS catalogue to the IRAC source catalogue using a matching radius of 2 ′′ .5 (Megeath et al. 2012;Jose et al. 2016) and found that 164 sources have counterparts in one or more IRAC bands.

Near-infrared photometry from 2MASS

NIR JHK s data for the point sources within the Hii region (shown in Fig. 1) have been obtained from the 2MASS Point Source Catalog (PSC, Cutri et al. 2003). Sources with uncertainty 0.2 mag and quality flag 'AAA' in all the three bands were selected to ensure good quality data. The spatial resolution of 2MASS JHK s bands is ∼ 2 ′′ .

RESULTS

Identification and Classification of Young Stellar Objects

Full understanding of cluster properties and its star formation history requires proper identification of cluster members. In the absence of spectroscopic observations, one robust approach is to identify young stellar content using multiwavelength photometric observations. Since IC 1805 is located at a low Galactic latitude, the contamination from background/foreground stars may significantly affect the analyses at infrared bands (e.g. Getman et al. 2006;Roccatagliata et al. 2011). Hence, careful identification and classification of probable YSO members is essential. During their early phase, stars possess cirmcumstellar disks and hence show excess emission at IR wavelengths. The IR excess properties of YSOs can be used to identify and classify them. Class '0' sources emit most of their radiation in far infra-red (FIR) to submillimetre regime of the electromagnetic spectrum (André 1995). In the Class 'I' phase, YSOs radiate mostly in MIR-FIR wavelengths, whereas Class 'II' sources exhibit NIR-FIR excess emission. Class 'III' sources possess little or no excess in IR [START_REF] Lada | IAU Symposium[END_REF]Adams, Lada & Shu 1987;André, Ward-Thompson & Barsony 1993;André 1995) but show enhanced X-ray activities as explained in Section 3.1.4. However, sources which lack inner disks and show excess emission in MIR wavelengths are referred to as 'transition disk' sources (Muzerolle et al. 2006;Hernández et al. 2006). We used optical, Spitzer-IRAC-MIPS, 2MASS and X-ray datasets to identify and classify the YSOs in IC 1805. The various selection criteria adopted for identification and classification of those YSOs are given below.

YSOs selected by IRAC colours

Since YSOs occupy distinct positions on the IRAC colour-colour diagrams due to their different characteristics, the colour-colour diagrams are widely used to identify and classify YSO population in a star-forming region (Megeath et al. 2004;Allen et al. 2004). Gutermuth et al. (2009) developed a method consisting of a series of IRAC/MIPS colour criteria to identify and classify YSOs. Though the mid-IR and far-IR observations can penetrate deep into the thick layers of dust and gas and reveal the embedded YSO population, they also enhance the chances of inclusion of background contaminants in the YSO sample. For example, polycyclic aromatic hydrocarbon (PAH)-emitting galaxies, active galactic nuclei (AGN), unresolved blobs of shocked emission etc., may mimic YSOs and, thus, can contaminate our YSO sample. In order to weed out possible contaminants, we used the customized cuts in the Here, σ is the photometric uncertainty in respective colours. The rest of the sources may be Class III or field stars. However, it is to be noted that in a few cases, due to reddening and disk inclination effects, a Class II source could mimic the colours of a Class I source (Whitney et al. 2003;Gutermuth et al. 2009).

YSOs from 2MASS/IRAC

At the central region of W4, in the vicinity of IC 1805, the background emission at 5.8 and 8.0 µm is high compared to the emission at 3.6 and 4.5 µm. This high background emission is often observed in massive star forming regions (e.g., Indebetouw et al. 2007;Koenig et al. 2008;Ojha et al. 2011). Many of the sources detected in 3.6 and 4.5 µm may be missing in To identify the YSOs using IRAC and 2MASS data, we first created an extinction map based on 2MASS data following the method discussed in Panwar et al. (2014), and then derived the intrinsic colours of the YSOs by correcting for the extinction values read off from this map. To generate the extinction map, we divided the region into small cells and computed the colour excess in each cell using the relation

E = (H -K s ) obs -(H -K s ) int ,
where (H -K s ) obs is the observed median colour of the stars in a cell and (H -K s ) int is the intrinsic median colour of the control field stars. The colour excess ratios presented in Flaherty et al. (2007) have been adopted to calculate the extinction in K s band (A Ks ) for each cell by using the relation A Ks =1.82 E. Based on the extinction map, the mean A Ks in the cluster direction corresponds to A V ∼ 2.6 mag, which is comparable to the value reported in the literature (Joshi & Sagar 1983). We calculated the (K s -[3.6]) 0 and ([3.6] -[4.5]) 0 colours of the sources and used various colour cuts as described in Gutermuth et al. (2009) to select Class I and Class II sources. By this method, we obtained 161 Class II and 6

Class I candidates. Among these, based on the IRAC colours, 5 Class I sources have been classified as Class II (Sec. 3.1.1), and for further analyses we have considered them as Class II candidates. Out of these 167 Class I/II candidates 133 were identified by using IRAC colourc 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 11 colour diagram. Fig. 2b shows the 2MASS/IRAC colour-colour diagram for sources having K s , 3.6 and 4.5 µm detection.

YSOs from MIPS 24 µm

To identify the transition disk candidates, we used the 24 µm data in addition to 4. We checked these Chandra X-ray sources for the counterparts of our optical dataset with a less stringent matching radius of 3 arcsec (Pandey et al. 2013b) because of the large off-axis beam. The agreement between X-ray and optical positions is excellent in majority of the cases (∼ 80%), with offsets <1 ′′ . In a few cases, where there was more than one source within the matching radius, we considered the closest one as the best match. This search revealed possible optical counterparts to 403 (62 %) of the X-ray sources.

Fig. 3 shows V vs. (V -R) colour-magnitude diagram (CMD) for the stars in the area covered by X-ray observations. The MS from Girardi et al. (2002) for the solar metallicity is also plotted. Joshi & Sagar (1983) have found that in the direction of IC 1805, E(B -V ) = 0.7 -0.9 mag, indicating a small variation of extinction within the cluster in agreement with the A V = 2.2 -2.7 mag, estimated by Straižys et al. (2013) from the 39 members of the cluster. The MS is shifted for the adopted distance of 2.0 kpc and the minimum reddening E(B -V ) of 0.7 mag. Our approach of using the minimum reddening is consistent with

the studies of open clusters (e.g., Sung & Bessell 2004;Sharma et al. 2006;Pandey et al. 2013b). The optical counterparts of the X-ray sources are shown with red circles. In Fig. 3 c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 13 Marino, Micela & Peres (2000) have discussed that the coronal activity of X-ray active field stars such as foreground dwarf M star (dMe) and background giants could lead to misidentification as young stars. For example, in Fig. 3, a few X-ray emitting sources are falling to the left or close to MS. Based on spectroscopic observations of the MS stars, Massey, Johnson & Degioia-Eastwood (1995) have inferred that the cluster is likely to be younger than 5 Myr. Considering the young age of the cluster, these sources are unlikely to be part of the cluster. Similarly, extended extragalactic objects seen along the line of sight could also be misidentified as young stars (Getman et al. 2006). Since we have performed PSF photometry on our optical images, we expect that the contamination due to extended objects is minimal in our sample; however we can not rule out the possible presence of a few point-like extra-galactic objects such as quasars. The field star contamination in our sample is quantified in Sect. 3.1.5, but at least quantitatively, it appears to be consistent with the amount of foreground X-ray contamination expected in the direction of IC 1805.

To remove the X-ray active field stars from our X-ray YSO sample, we considered only those X-ray sources which are located in the empirical YSO zone on the V vs. (V -R)

CMD (e.g., Allen et al. 2012). To determine the YSO zone, we used the following process iteratively. We arranged the X-ray sources in 0.75 mag bins between V =15.0 to V =24 mag.

The median and standard deviation (σ) of the (V -R) colour of the YSOs in each bin were calculated and the YSOs having colours greater than 3σ from the median were considered as outliers. The median colour and σ were recalculated. The lower and upper limits of the YSO zone in each bin are then taken to be 1.25 σ and 2 σ around the median colour of the bin. These upper and lower limits are then fit with polynomials and the YSOs lying within these upper and lower bounds are taken to be probable members. As discussed in Allen et al. (2012), the lower limits of the magnitude bins are taken closer to the median value because the density of sources quickly increases in the direction perpendicular to the isochrone towards the MS. In particular, foreground dwarfs and background giants may disguise as YSOs by appearing more luminous for a given colour. We note that although the above statistical approach is likely to remove non-YSO X-ray sources, some actual YSOs may also be eliminated. For example, a few X-ray sources located near the MS (where majority of the field stars are populated), could be YSOs, although their chances of being so are low.

Spectroscopic observations are needed to ascertain the nature of such sources, because some of them could be YSOs with disk systems seen edge-on that appear under luminous for their colours because of scattering (Sicilia-Aguilar et al. 2005) or members that are significantly older than the mean cluster age.

Using the above approach, we obtained 337 X-ray emitting YSO candidates in the region covered by Chandra X-ray observations (for details; see the Section 3.2.2). Of these, 40 sources have shown the characteristics of Class I/II candidate YSOs on the basis of 2MASS

and Spitzer colours (discussed in Sections 3.1.1, 3.1.2, 3.1.3), so the rest 297 X-ray sources are considered as Class III candidates.

Data completeness and contamination

As in the c2d legacy project (Evans et al. 2003), the peak of the observed luminosity function can be assumed to estimate 90% completeness limit (see also Jose et al. 2016). We constructed the luminosity functions for each band by using the histogram method (see Fig. 4). The resultant completeness limits of the optical and IRAC data are given in Table 2.

c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 15

For cross-checking, the completeness limits of V and R bands were also estimated by the artificial star experiment by using ADDSTAR routine in IRAF (e.g., Pandey et al. 2013a;Jose et al. 2013). We randomly added artificial stars of various magnitudes into each image.

The luminosity distribution of artificial stars was chosen in such a way that more stars were inserted towards the fainter magnitude bins. The frames were then re-reduced by using the same procedure as used for the original frames (see Section 2.1). The ratio of the number of stars recovered to those added in each magnitude interval gives the completeness factor as a function of magnitude. We thus obtained >90% completeness of the photometry to be 22.5 and 22 mag in V and R bands, respectively, which agrees with that of the histogram analysis.

The estimation of the completeness census of the YSO sample obtained by various colour combinations is rather difficult as it is limited by several factors. For example, bright extended nebulosity in the IRAC bands significantly limits the point source detection. The YSO identification from the 2MASS/IRAC colours is limited by the sensitivity of the 2MASS survey, while that from the IRAC/MIPS colours suffers from the significant saturation in the IRAC 8 µm and MIPS 24 µm images caused by the central luminous sources as well as the bright nebulosity. Similarly, the variable reddening and stellar crowding characteristics across the region could also affect the local completeness limit. All these effects are difficult to quantify. Since 3.6 and 4.5 µm bands are the most commonly used colour combination for identifying Class II YSOs, we calculate the approximate completeness limit of our Class II YSOs by using the completeness limit of these two bands. The 90% photometry completeness limits of 3.6 and 4.5 µm bands are estimated to be ∼ 16 mag. Assuming a distance of 2.0 kpc for IC 1805 and an average extinction A V of ∼ 2.5 mag, we find the photometric completeness limit corresponds to an approximate stellar mass in the range 0.2-0.3 M ⊙ for a YSO of age ∼ 2-3 Myr (using the evolutionary tracks of Siess, Dufour & Forestini 2000).

The identification of Class III sources in our sample is primarily based on their detection in both the X-ray and optical bands. Only 62% of the X-ray sources have optical counterparts. This indicates that the completeness limit of Class III YSOs is primarily dictated by the sensitivity limit of the optical data. The 90% photometry completeness limits of V and R bands lead to the completeness limit of our Class III YSOs as ∼ 0.3 M ⊙ at the assumed distance, mean A V and age ∼ 2.5 Myr (using the evolutionary tracks of Siess, Dufour & Forestini 2000). Adopting different sets of evolutionary tracks would provide different values of stellar masses. However, for low-mass objects, the tracks of Siess, Dufour & Forestini (2000) are close to those of Baraffe et al. (1998). The agreement between masses of these two models is within 10-20%. In summary, we considered that our YSO sample is expected to be largely complete above 0.3 M ⊙ .

In the case of IC 1795, based on Franceschini et al. (2006) 2009), we have removed these contaminants from our catalog. Although Gutermuth et al. (2009) criteria for a region at ∼ 2 kpc may provide an overestimation of the contamination, but this would ensure the high reliability of our YSO sample. Our X-ray YSO sample is selected based on the X-ray and optical observations. The majority of the extragalactic contaminants and distant background stars are expected not to have optical (V-band) counterparts. Moreover optical images being at high spatial resolution and with our PSF photometry, extended objects like galaxies are unlikely to be the major contaminants in our sample. In the optical CMD, most of the likely members are located in the PMS zone and there are ∼ 40 X-ray sources located near the MS. As discussed in Sect. 3.1.4, these sources are likely foreground sources of the region, and seem to agree well with the expected foreground contamination in the direction of IC 1805. So overall, we expect that the likely number of X-ray contaminants in our YSO sample should be low.

Final catalog of YSOs

Our final catalog includes the YSO candidates identified from Spitzer IRAC/MIPS, 2MASS and X-ray datasets. In total, we have 297 Class III, 8 transition disk and 272 Class I/II candidates in the direction of W4 complex. However, there are only 101 Class I/II and 283 Class III candidates in the cluster (i.e., within ∼ 9 ′ ; see Section 3.2.1). The cluster radius is well within the common area covered by the X-ray, IRAC-MIPS as well as optical c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 17 In this work, we are more interested in studying the optical properties of the YSOs associated with the cluster. Unlike, NIR colours, the optical filters are more sensitive to the temperatures of late-type stars. Thus, in the low-extincted regions like IC 1805, the optical CMDs are expected to be more reliable in separating PMS members from the field stars and also better for the characterization of YSOs.

Properties of the Cluster IC 1805

Physical Extent of the Cluster

The radial extent is one of the important parameters to study the dynamical properties of clusters. To estimate the cluster radius of IC 1805, we applied the star count technique and assumed a spherically symmetric distribution of stars in the cluster. From the RDP, the cluster radius appears to be ∼ 9 ′ (5 pc).

For comparison, we also estimated the extent of the cluster (r cl ) using 2MASS data.

Within uncertainties, the r cl obtained from the YSO sample and 2MASS data are matching.

The RDP is parametrised with the empirical model of King (1962),

ρ(r) ∝ ρ 0 1 + r r c 2 ,
where r c is the core radius at which the surface density ρ(r) becomes half of the central density, ρ 0 . The best fit (solid line) to the radial density obtained by a χ 2 minimization technique is shown in Fig. 5.

0 0.5 1 1.5 2 (V-R)
Figure 6. V /(V -R) colour-magnitude diagram for the YSO candidates in the cluster region. The PMS isochrones from Siess, Dufour & Forestini (2000) and MS from Girardi et al. (2002), corrected for the adopted distance and reddening are also overplotted. The arrow is a reddening vector corresponding to A V = 3 mag. Diamonds and squares represent the Class II and Class I sources, respectively. Blue crosses represent the X-ray sources which are considered as YSOs (see Sec. 3.1.4 and Sec. 3.2.2) in the present study.

Colour-Magnitude Diagram of YSOs : Age and Mass Estimation

The ages and masses of the cluster members were estimated by comparing their locations on the CMD with PMS isochrones of various ages after correcting for the distance and extinction. The minimal differential extinction (∆ A V ∼ 0.6 mag) towards the cluster is an advantage to examine the position of the YSOs on the CMD in order to confirm their youth, to determine their optical properties and to study the star formation history of the region.

Fig. 6 shows V vs. (V -R) CMD for the YSOs in the cluster area. The MS from Girardi et al. (2002) as well as the PMS isochrones for 1, 5 and 10 Myr for the solar metallicity from Siess, Dufour & Forestini (2000), corrected for the adopted distance and reddening are also overplotted. The Class I and Class II sources (squares and diamonds, respectively)

within the cluster area are identified based on the Spitzer IRAC/MIPS photometry. The location of X-ray emitting (likely Class III) sources are shown with the crosses. The age distribution of the majority of these sources is in the range of 0.5 -7 Myr, concentrated around an age in the range of 2-3 Myr (mean age 2.5 ± 1.5 Myr). We find that masses for a majority of the YSOs are in the range of 0.2 M ⊙ -2.5 M ⊙ . As discussed in Sect. 3.1.4, our approach to select the Class III candidates might have removed some YSOs, which may affect the mean age of the cluster. Hence, we also estimated the age of the cluster using only the disk-bearing sources (i.e, sources that have shown excess emission in the 2MASS and IRAC bands). This method resulted an age ∼ 2.3 Myr for the cluster, which is comparable to the mean age estimated using all the YSOs.

Age spread is common in young clusters and is probably due to the combined effect of differential reddening, circumstellar extinction, variability, binarity and/or different evolutionary stages (Jose et al. 2017). However, the reddening vector is nearly parallel to the isochrones, hence differential reddening does not have much effect on the age estimation of YSOs. The circumstellar extinction may affect the age and mass determination of Class 0/I sources, but its effect should be minimal in optically visible PMS sources such as Class III YSOs. Though the exact role of circumstellar disks requires detailed SED modelling or spectroscopic observations, which is beyond the scope of this paper, it is worthwhile to mention that Jose et al. (2013) performed SED modeling on a set of optically visible PMS sources and found that SED based ages are largely in agreement with those estimated from the optical CMDs. Binary companion will apparently brighten the star, consequently CMD will yield a younger age. For example, in the case of equal-mass binaries, cluster is expected to show a sequence in the CMD which is shifted by 0.75 mag upwards. The effect of binarity can be seen more prominently in some older clusters (∼ 10 Myr) for which the isochrones are close together (e.g., NGC 7160; Sicilia-Aguilar et al. 2005) whereas in the case of young clusters like Tr 37 (∼ 4 Myr), which is similar to IC 1805, the isochrones are well separated, hence the effect of binarity on age estimation will be less compared to the natural age spread of the stars in clusters (see discussion in Sicilia-Aguilar et al. 2005). Similarly, though TTSs tends to show variability at optical bands (e.g., Herbst et al. 1994;Briceño et al. 2001;Rodríguez-Ledesma, Mundt & Eislöffel 2010;Lata et al. 2015), the variability tends to move the objects parallel to the isochrones, resulting in little effect on the age estimation of YSOs (Burningham et al. 2005). Despite of low differential reddening and presence of large number of disk-less sources in the cluster, some degree of age spread due to the combination of the above factors is expected, but since the observed spread in colour due to member YSOs is larger than their mean uncertainties due to photometric colours (see Fig 6), we expect that the YSOs are at different evolutionary stages.

We considered that the errors associated with the determination of age and mass are mainly of two kinds; random errors in photometry, and systematic errors due to different theoretical evolutionary tracks. We estimated the effect of random errors by propagating them to the observed estimations of V , (V -R) and E(V -R) by assuming a normal error c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 21 distribution and using Monte Carlo simulations (see e.g., Chauhan et al. 2009). Since we have used the evolutionary model by Siess, Dufour & Forestini (2000) for all the PMS stars, the age and mass estimations given in Table 4 should not be affected by systematic errors. A sample of Table 4 is given here and the complete table is available in the electronic version.

DISCUSSION

Mass function of the YSOs

Young clusters are the best objects to study the initial MFs. Since young low-mass stars have limited time to segregate from the cluster and/or loose mass through evolutionary processes, the variation of the MF gives clues about the physical conditions of the star formation process. The MF is defined as the number of stars per unit logarithmic mass interval, and is generally represented by a power law with a slope,

Γ = d log N(log m)/d log m,
where N(log m) is the distribution of the number of stars per unit logarithmic mass interval.

Most of the star clusters in the solar neighborhood have MF slopes near the Salpeter MF slope, i.e., Γ = -1.35.

Since our YSO sample is complete down to 0.3 M ⊙ , for the MF study we have taken only those YSOs which have masses in the range 0.3 M/M ⊙ 2.5. The mass distribution of our YSO sample has a best-fit slope, Γ = -1.23 ± 0.14 (see Fig. 7). Ninkov et al. (1995) studied the cluster IC 1805 and found a slope of Γ = -1.38 ± 0.19 for masses between 2.5 and 30 M ⊙ . A similar value of Γ (= -1.3 ± 0.2) has been reported by Massey, Johnson & Degioia-Eastwood (1995) for the massive stars of the W4 complex. Our YSO MF is consistent with those reported for other active star forming regions. Erickson et al. (2011) derived the MF of an extinction limited YSO sample (with masses >0.2 M ⊙ ) consistent with the field star MF. Kang et al. (2009) did not find any difference in the YSO MF slopes of the star-forming regions W51 A (Γ = -1.17 ± 0.26) and W51 B (Γ = -1.32 ± 0.26). We note that a few older Class III sources might have been excluded as contaminants due to our conservative selection criteria (see Sec. 3.1.4). Even then, we believe our result should represent the IMF for the cluster, for example, and low-mass isochrones are very close to each other.

Evolution of T-Tauri Stars in IC 1805

TTSs are low-mass stars (< 3M ⊙ ) which are contracting towards the MS. They are generally classified into weak-line TTSs (WTTSs) and classical TTSs (CTTSs) on the basis of the strength of the H α emission line (Strom et al. 1989). WTTSs exhibit a weak, narrow H α (EW 10 Å) in emission with no or little infrared excess, while CTTSs display a strong H α emission line (EW 10 Å) and large infrared excesses. Class II sources are generally considered to correspond in the optical category to CTTSs and Class III sources to WTTSs (e.g., Hartmann et al. 2005;Pandey et al. 2013b). The 'standard model' by Kenyon & Hartmann (1995) postulates that the CTTSs evolve to WTTSs by losing their circumstellar disks or at least their inner parts. In the Taurus region the WTTSs are systematically older than the CTTSs (Bertout, Siess & Cabrit 2007), but the statistical significance is low (Kenyon & Hartmann 1995;Hartmann 2001;Armitage, Clarke & Palla 2003). Contrary to the above results, there are studies which favour the assumption that CTTS and WTTS are coeval and possess indistinguishable stellar properties (Lawson, Feigelson & Huenemoerder 1996;Gras-Velázquez & Ray 2005). The coevality of CTTSs and WTTSs in a star-forming c 2017 RAS, MNRAS 000, 000-000 region can be explained by assuming that YSOs have intrinsically a wide range of disk masses and their accretion activity and/or the dispersal of the disk takes place in a correspondingly wide range of time scales (Furlan et al. 2006;Bertout, Siess & Cabrit 2007).

In the present work, we have derived the ages of a large sample of YSOs in the IC 1805 cluster, hence, it is worthwhile to attempt to address the problem of the evolutionary status of CTTSs and WTTSs. Fig. 8a shows the histograms of the ages of the CTTS and WTTS candidates within the cluster region in the mass range 0.3 -2.5 M ⊙ . This manifests that the CTTSs are relatively younger (mean age ∼ 2.3 Myr) than the WTTSs (mean age ∼ 2.6 Myr). The cumulative distribution of the ages of CTTSs and WTTSs are shown in Fig. 8b. We obtained the Kolmogorov-Smirnov (KS) test value of 11% that the two samples are drawn from the same population, which is marginally significant. Although the statistical significance is less, our result is in agreement with those of Bertout, Siess & Cabrit (2007) for the Taurus-Auriga T-association and of Chauhan et al. (2009) for the BRC regions, that WTTSs are relatively older than CTTSs, which supports the notion that CTTSs evolve into WTTSs. Here again, we would like to remind readers that these results are affected by the uncertainty associated with the membership of the stars (see Sec. 3.1.4), particularly the WTTSs (Spectroscopic information is needed to confirm the true nature of WTTSs and CTTSs).

Disk Fraction Variation in the Cluster

The disk fraction of YSOs in clusters decreases with ages (e.g., Haisch, Lada & Lada 2001;Hernández et al. 2008). Haisch, Lada & Lada (2001) quantitatively concluded that in young clusters an initially very high (80% -90%) disk fraction decreases to 50% in ∼ 3 Myr and reaches almost 10% at ∼ 5 Myr. Generally massive stars seem to lose their disks earlier than lower-mass stars. For example, most low-mass stars (spectral type K5 and later) lose their disks in 5-7 Myr (e.g., Sicilia-Aguilar et al. 2005), whereas for the intermediate-mass

Herbig Ae/Be stars the corresponding time scale is <3 Myr (e.g., Hernández et al. 2005).

In a young cluster the disk fraction depends also on the spatial position. Theoretical calculations predict that external UV radiation of high-mass stars can photoevaporate outer c 2017 RAS, MNRAS 000, 000-000

Low-mass young stellar population of the cluster IC 1805 25 disks only within 0.3 -0.7 pc (Johnstone, Hollenbach & Bally 1998;Adams et al. 2004).

Decrease in disk frequencies in the immediate vicinity of massive stars has been found in several massive clusters, e.g., NGC 2244, NGC 6611 and Pismis 24 (Balog et al. 2007;Mercer et al. 2009;Fang et al. 2012), suggesting rapid destruction of circumstellar disks in such environments. IC 1805 is a young cluster (age ∼ 2.5 Myr) with several low-mass (< 2.5 M ⊙ ) stars that have disks, thus the cluster represents a good site for studying the influence of massive stars on the evolution of disks.

In IC 1805, it is argued that the region is mainly powered by three O-type stars (Lefloch, Lazareff & Cast 1997). BD +60502 and BD +60501 are of spectral type O4.5III (Sota et al. 2014) and O7V (Sota et al. 2011), respectively and are located near the assumed cluster centre (see Fig. 9), whereas HD 15570 is an O4I star (Sota et al. 2011), located ∼ 4.7 ′ east of the cluster centre.

The radial velocities of BD +60502 and BD +60501 are -42 km/s (Kharchenko et al. 2007) and -52 km/s (Kharchenko et al. 2007), respectively, consistent with the molecular cloud velocity (-55 to -32 km/s) associated with the W3/W4 complex (Heyer & Terebey 1998).

The radial velocity of HD 15570 is -24.00 km/s [START_REF] Gontcharov | Low-mass young stellar population of the cluster IC 1805 31[END_REF], thus unlikely to belong to the W3/W4 complex, but seems consistent with that of an object in an inter-arm cloud, whose velocity lies between -32 to -18 km/s (Heyer & Terebey 1998). We conclude that the cluster is dominated by the O4.5III and O7V stars located near the cluster center.

To study the influence of these massive stars on the evolution of nearby low-mass stars, we divided the cluster area into five concentric annuli (each of width 1.8 ′ ). the ratio between CTTSs and CTTSs+WTTSs in each annular area) as a function of the projected radial distance from the cluster centre which is only 0.03 pc away from BD +60502.

In the Figure, the substantial uncertainties shown as solid lines originate from the low statistics of IR excess sources in each bin. We also checked the distribution of disk frequency by re-centering the analysis on the O7 star and find the disk frequency variation to be nearly the same. Although we see a dip in the disk-fraction vs. distance plot in the vicinity of the O-type star, due to large statistical error, we do not have a strong evidence of variation in the disk fraction within the cluster. However, we note that the disk fraction estimation can be affected by many factors, such as various methods used to identify YSOs, the variation in stellar density (e.g., Spezzi et al. 2015), sensitivity of different bands and the uncertainty in the membership of the older WTTs. For example, Spitzer observations of high mass star forming regions suffer from high, non-uniform nebulosity, primarily emission from PAHs at IRAC bands and is strong close to the massive stars near the center of rich clusters, where stellar crowding is high. In contrast Chandra's sensitivity is less affected by nebulosity and moreover it has several times better on-axis spatial resolution than Spitzer.

So IRAC bands can lead to severe decrease in sensitivity of the disk-bearing YSOs compared to disk-less YSOs near the cluster center. Nonetheless, our result is in agreement with Roccatagliata et al. (2011), who found no variation of the disk fraction as a function of the distance from the high-mass stars in IC 1795, which has similar number of high mass stars and located in the same cloud complex.

Star Formation History

Figure 10 shows the Canadian Galactic Plane Survey (CGPS; Taylor et al. 2003) 1.4 GHz emission in the direction of the W4 bubble, superimposed on the WISE4 12 µm emission and Herschel5 in 250 µm emission. We have also marked the BRCs in the region (Sugitani, Fukui & Ogura 1991). The 12 µm WISE band contains 11.7 µm emission commonly attributed to PAH molecules excited in the photon dominated regions (PDRs) at the interface of Hii regions and their adjacent molecular clouds by the far-UV photons leaking from the Hii regions (Pomarès et al. 2009). Therefore, PAHs are good tracers of the warm PDR that surrounds the Hii region. PAH emission is also a good tracer of newly formed, embedded B-type star formation (Peeters, Spoon & Tielens 2004), as these stars heat the surrounding dust to high temperatures enough to excite the PAH bands and fine-structure lines. Figure 10 reveals a filamentary distribution of the 12 µm emission with its long axis in the east-west direction that bisects the Hii region and also in the PDRs at the periphery of the bubble. The dust emission detected by the Herschel 250 µm also has an elongated distribution similar to that of the 12 µm emission. Compared to the PAH and dust emission, the ionized gas appears to be distributed orthogonal to the long axis of the filamentary structure, with decreasing intensity away from the filament axis. The ionized gas seems to be bounded more towards the east and west directions than in the north and south. The overall morphology of the nebula looks bipolar. Such a bipolar morphology has been noticed in NIR to MIR bands in a few Hii regions and/or bubbles (see e.g., Saito et al. 2009;Samal et al. 2012;Mallick et al. 2013;Deharveng et al. 2012). Recently Herschel observations have clearly shown the presence of a cold dense neutral filament in bipolar Hii regions that bisects ionized lobes (e.g., Deharveng et al. 2015), similar to the one observed here.

Thus, the elongated nature of the 250 µm emission most likely represent the long axis of the parental filament.

It has been suggested that the stellar distribution in star forming regions is governed by the structure of the parental molecular cloud as well as how star formation proceeds in the cloud (e.g., Chen, Chen & Shu 2004; Sharma et al. 2006;Schmeja, Kumar & Ferreira 2008;Gutermuth et al. 2009;Samal et al. 2015;Jose et al. 2016). Although our survey does not sample all the YSOs of the W4 complex, it is good enough to provide clues on the star formation history. We generated the YSO stellar surface density map using the nearest neighborhood method (Gutermuth et al. 2005). At each sample position [i, j] in a uniform grid, we measured r N (i, j), the projected radial distance to the N th nearest star. The value of N is allowed to vary depending upon the desired smallest scale structures of interest.

We generated a map using N=20, which, after a series of experiments, was obtained as a good compromise between the resolution and signal-to-noise ratio of the map. In the stellar surface density distribution (see Fig. 10), the YSOs show a centrally concentrated clustering.

At the same time, the overall distribution has an elongation along the east-west direction, 2012) have suggested that young clusters in the W5 complex were also formed in a filamentary cloud. In the W4 complex, we could see a few BRC structures in the north and south directions facing the IC 1805 cluster (see Fig. 10). This could be due to the fact that as the ionized gas streams away in the low density side, it encounters a few small, dense clouds along its way, forming the BRC structures.

The YSO distribution along the filament axis of the W4 bubble suggests that some of them could have resulted from the fragmentation of the original filament as are the cases found in other filamentary systems (e.g., Jackson et al. 2010;Hacar & Tafalla 2011;Samal et al. 2015) and/or due to compression of dense gas of the primordial filament by an expanding Hii region. Though precise characterization (e.g., proper motion and age) of individual YSOs projected on the filament is needed to understand their origin, the overall morphology of the W4 Hii region and the distribution of the associated YSOs suggest that IC 1805 could have been formed in a filamentary cloud.

With the aim of unraveling the less known low-mass YSO population of the cluster IC 1805 which is rich in massive OB stars, we studied the region using optical, IR and X-ray datasets.

Our results suggest that despite the less favorable conditions for star formation around the high-mass stars, IC 1805 shows signs of low-mass star formation similar to other clusters in the solar vicinity. We identified and characterized a large number of low-mass YSO candidates in the region. Most of them have masses in the range 0.2 -2.5 M ⊙ and age 0.1 -5 Myr.

The slope of the MF in the mass range 0.3 M/M ⊙ 2.5 is found to be Γ = -1.23 ± 0.14, similar to the Salpeter MF. The mean age of the candidate YSOs is found to be ∼ 2.5 Myr.

The candidate WTTSs are found to be relatively older when compared to the CTTS candidates. We do not find strong evidence for disk dispersal due to massive stars. The spatial distribution of the YSOs, dust and gas in the Hii complex shows a filamentary distribution, which suggests the filamentary morphology of the parental molecular cloud.

Figure 1 .

 1 Figure 1. The DSS2-R image of the IC 1805 Hii region. The area covered by X-ray, optical and Spitzer-IRAC observations are shown with blue, green and red colours, respectively (also marked as '1', '2' and '3').

  then performed point-spread function (PSF) photometry on the selected sources using the ALLSTAR routine. These deep optical observations were obtained from the archive with the aim to characterize the low-mass members (in particular to obtain the age and mass of the low-mass young stellar population) of the cluster IC 1805. To calibrate these deep images, we obtained the V and R band observations of the central region of IC 1805 on 2012, October 26 by using a 2048 × 2048 pixel 2 CCD camera mounted at the f/4 Cassegrain focus of the 1.3-m Devasthal Optical Telescope (DOT) at Aryabhatta Research Institute of

Figure 2 .

 2 Figure 2. (a) IRAC ([3.6]-[4.5])/ ([5.8]-[8.0], (b) IRAC/2MASS, and (c) IRAC/MIPS colour-colour diagrams for the sources (black dots) in IC 1805. The error bars represent the 1-σ uncertainty in respective colours. Class I and Class II sources are shown with red square and blue diamond symbols, respectively. Transition disk candidates are shown with magenta triangles.

  [4.5]-[5.8] vs. [5.8]-[8.0] and [3.6]-[5.8] vs. [4.5]-[8.0] colour-colour spaces, as described by Gutermuth et al. (2009). After excluding the contaminants based on the Gutermuth et al. (2009) approach, we identified 238 YSO candidates. Of these, 4 sources satisfy the colour criteria of Class I sources ([4.5] -[5.8] > 0.7 & [3.6] -[4.5] > 0.7) and 234 sources (satisfying the following conditions) were classified as Class II candidates : [4.5] -[8.0] -σ > 0.5, [3.6] -[5.8] -σ > 0.35, [3.6] -[5.8] + σ (0.14/0.04) × (([4.5] -[8.0] -σ) -0.5) + 0.5, and [3.6] -[4.5] -σ > 0.15;

  Fig. 2a shows the [3.6] -[4.5] vs. [5.8] -[8.0] colour-colour diagram for all the uncontaminated IRAC sources. The Class I and Class II sources are shown as red squares and blue diamonds, respectively.

5. 8

 8 and 8.0 µm bands due to the enhanced nebulosity of the background. To ascertain the nature of those sources, we combined the IRAC 3.6 and 4.5 µm data with those of 2MASS H and K s bands. Gutermuth et al. (2009) used the intrinsic (K s -[3.6]) vs. ([3.6] -[4.5]) colour criteria to identify the sources which are detected in 3.6, 4.5 µm and possess counterparts in 2MASS bands.

  data. FollowingMuzerolle et al. (2006) andHernández et al. (2006), sources mainly with photospheric radiation in the IRAC or IRAC/2MASS classification but having excess emission at 24 µm ([5.8] -[24] > 2.5 and/or [4.5] -[24] > 2.5) are reconsidered as transition disk sources. We identified 8 such sources. Since heavily embedded Class II sources can be misclassified as Class I based on the IRAC colour criteria (see Section 3.1.1), those Class I sources with 24 µm data were re-checked to ensure that their SEDs rise from the IRAC to MIPS bands. All Class I candidates were considered as Class II, if they do not have [5.8]-[24] > 4 or [4.5] -[24] > 4. Out of 4 Class I sources identified in Section 3.1.1, 2 are re-classified as Class II, while the remaining 2 do not possess 24 µm measurements. Fig. 2c shows the IRAC/MIPS colour-colour diagram for the sources in the region. In Fig. 2, transition disk candidates based on the MIPS photometry are marked as magenta filled circles.3.1.4 YSOs from Chandra X-ray DataYSOs are often strong X-ray emitters compared to their main-sequence (MS) counterparts (e.g.,Feigelson & Montmerle 1999;Getman et al. 2005;Preibisch & Feigelson 2005;Güdel et al. 2007). Hence, X-ray surveys of star forming regions are used to uncover YSO populations. For the detection of diskless YSOs, i.e., Class III sources, X-ray observations complement the IR YSO sample to have a more complete census of stellar content of clusters (e.g.,Preibisch & Feigelson 2005; Pandey et al. 2013bPandey et al. , 2014)). The Massive Star-Forming Regions Omnibus X-ray Catalog (MOXC,Townsley et al. 2014) detected X-ray sources in a sample of massive star-forming regions, including IC 1805. They characterized the X-ray properties of 647 stars in the IC 1805 using Chandra/ACIS observations. ACIS-I consist of four 1024 × 1024 pixel 2 CCD chips that cover a 17 ′ × 17 ′ field of view, and covers the 0.5 -8.0 keV energy band with a spectral resolution of ∼ 150 eV at 6 keV and a PSF radius of 0 ′′ .5 within ∼ 2 ′ of the on-axis position, degrading to ∼ 6 ′′ at a 10 ′ off-axis angle.

Figure 3 .

 3 Figure3. V /(V -R) colour-magnitude diagram for the stars in the region covered by X-ray observations (grey dots). The MS fromGirardi et al. (2002), corrected for the adopted distance and reddening are also overplotted. The average error in (V R) colour as a function of magnitude is shown in the left-hand side. Red circles represent the X-ray sources in the region.

  , at the low-luminosity end, one can notice a few X-ray sources scattered close to MS. Similar distribution of X-ray sources close to MS has also been observed in the CMD of other young clusters. For example, the distribution of X-ray sources on the V vs. (V -I) diagram of NGC 6530 (median age ∼ 2.3 Myr and mean E(B-V)=0.35 Prisinzano et al. 2005), shows 90% of PMS stars are within 0.3 -10 Myr, and ∼ 10% of the X-ray sources are outside this age range. Prisinzano et al. (2005) argued that the latter sources are mainly contaminants.

Figure 4 .

 4 Figure 4. Histograms for the sources detected in optical (V,R) and IRAC bands showing the limiting magnitude and completeness limit for each band.

Figure 5 .

 5 Figure 5. Radial density profile of the cluster IC 1805 based on the YSO sample (left panel) and 2MASS data (right panel). The continuous curve shows the least-squares fit of the King (1962) profile to the observed data points. The error bars represent ± √ N errors.

Figure 7 .

 7 Figure 7. The mass function of the candidate YSOs in the mass range (0.3 M/M ⊙ 2.5), derived from the optical data. The error bars represent ± √ N errors. The continuous line shows the least-square fit to the mass ranges described in the text. The value of the slope obtained is given in the figure.

Figure 8 .

 8 Figure 8. (a) Histograms of the age distribution of the Class II and Class III candidates, (b) Cumulative distributions of CTTSs and WTTSs in the cluster region as a function of stellar age.

Figure 9 .

 9 Figure 9. The upper panel shows the surface distribution of WTTSs (blue circles) and CTTSs (red circles) superposed on the Spitzer 4.5 µm image of the cluster. The concentric annuli have a width of 1 ′ .8. The massive stars are shown with diamond symbols and cluster centre with cross. The lower panel shows the CTTS fraction as a function of the radial distance from the cluster centre. The two filled squares represent the locations of the two massive stars near the cluster centre.

  Fig. 9 (upper panel) displays the distribution of CTTSs (red circles) and WTTSs (blue circles) in the mass range 0.3 -2.5 M ⊙ from the catalogue of YSOs for which we have mass estimations from the optical CMDs. Our aim is to examine the disk fraction variation of the low-mass stars within the cluster due to external influence, as massive stars can self-destroy their disk, because evolution of protoplanetary disks is faster around higher mass stars (e.g.Kennedy & Kenyon 2009;Yasui et al. 2014). Fig.9(lower panel) displays the fraction of low-mass CTTSs (i.e.,

Figure 10 .

 10 Figure 10. Colour composite image of IC 1805 region reproduced with CGPS 1.4 GHz emission (blue colour) superimposed on those of the WISE 12 µm emission (green colour) and Herschel 250 µm emission (red colour). BRCs associated with the region and cataloged in Sugitani, Fukui & Ogura (1991) are also marked in the figure. The red contours display the surface density distribution of the YSOs with the plus symbol representing the cluster centre.

  following the structure revealed by the 250 µm emission. Thus it appears that the formation of the cluster IC 1805 possibly started in a filamentary cloud, where the density was high along the axis of the filament and low towards the perpendicular directions. With time the O stars of the cluster developed an Hii region, which subsequently grew in size, first inside the dense filament, then opening out in a hole in the low density direction of the filament, resulting a large cavity in the north and south direction. Based on the distribution of young stars on the Herschel images, Deharveng et al. (

Table 1 .

 1 The zero-point constants, colour coefficients and extinction coefficients

	Parameters	Constants
	Zero-point constants
	c1 c2 Colour coefficients	-0.226 ± 0.016 3.114 ± 0.014
	m1 m2 Extinction coefficients	1.058 ± 0.022 0.099 ± 0.024
	Kv Kr	0.158 ± 0.015 0.109 ± 0.011
	the SA 98 standard field observations are given in Table

Table 2 .

 2 Completeness limit of the data in IRAC and optical bands.

	Wavebands	V	R	[3.6] [4.5] [5.8] [8.0]
	Completeness (in mag) 22.5	22	16	16	14	13.5
	V	R			[3.6]

Table 3 .

 3 YSOs from 2MASS, IRAC/MIPS and X-ray data for the cluster area. The entire table is available in an electronic form. YSOs within the cluster radius, whereas to get an idea on the large scale star-formation history, we have used all the YSOs identified in this work (discussed in Sect. 4.4). A sample list of the YSO candidates with their magnitudes in different bands for the cluster area is given in Table3. The entire table is available in electronic form only.

	Id	RA	DEC	J±eJ	H±eH	Ks ±eKs	[3.6]±	[4.5]±	[5.8]±	[8.0]±	[24]±
		(J2000)	(J2000)				e[3.6]	e[4.5]	e[5.8]	e[8.0]	e[24]
	1	38.240929	61.374401	13.63 ± 0.04	12.65 ± 0.02	11.99 ± 0.02	12.38 ± 0.01	10.34 ± 0.01	9.58 ± 0.01	8.57 ± 0.01	6.94 ± 0.02
	2	38.182259	61.469330	12.15 ± 0.02	11.65 ± 0.04	11.02 ± 0.03	9.58 ± 0.01	8.88 ± 0.01	8.03 ± 0.01	6.72 ± 0.01	4.15 ± 0.01
	3	38.464951	61.406620	14.31 ± 0.04	13.88 ± 0.05	13.65 ± 0.04	13.61 ± 0.01	13.39 ± 0.01	13.02 ± 0.02	12.48 ± 0.03	-
	.	.....	.....	.....	.....	.....	.....	.....	.....	.....	....
		observations and our main aim is to explore and characterize the young stellar population
		associated with the cluster IC 1805. To determine the cluster properties, such as age (see Sect
		3.2.2), mass function, disk-fraction and disk evolution of YSOs (discussed in sections 4.1,4.2,
		and 4.3), we have used

In the W4 complex, using Bayes classifier and combining several criteria in a probabilistic approach,

Broos et al. (2013)

defined the list of MYStIX Probable Complex Members (MPCM). Within 9 arcmin of the cluster center,

Broos et al. (2013)

identified 389 YSOs (Xcl flag=2). Although a direct comparison with

Broos et al. (2013)

is not possible because of different methods used to extract photometric magnitudes on the Spitzer images (aperture photometry by them versus PRF photometry by us), and in the approach of YSO classification and contamination removal (for details see

Broos et al. 2013

). Nonetheless we find reasonable agreement between both the works. For example, within the cluster area we obtained 384 sources, whereas

Broos et al. (2013)

have identified 389 sources and only ∼ 15% sources are missing in either of the catalog.

Table 4 .

 4 Magnitudes, age and mass of the YSOs in the cluster region. Complete table is available in an electronic form.

	Id	RA (J2000)	DEC (J2000)	V ± eV	R ± eR	Age ± error Age (Myr) Mass ± error in Mass (M ⊙ )
	1 2 3 ...	37.888229 37.895050 37.919338 ...	61.429279 61.506271 61.502380 ...	22.10 ± 0.02 21.28 ± 0.01 19.87 ± 0.01 ...	20.66 ± 0.01 19.79 ± 0.01 18.61 ± 0.01 ...	3.01 ± 0.65 1.20 ± 0.30 1.57 ± 0.57	0.31 ± 0.02 0.30 ± 0.02 0.57 ± 0.08

c 2017 RAS, MNRAS 000, 000-000

http://archive.spitzer.caltech.edu/ c

RAS, MNRAS 000, 000-000

http://irsa.ipac.caltech.edu/data/SPITZER/docs/mips/mipsinstrumenthandbook/ c 2017 RAS, MNRAS 000, 000-000

https://www.univie.ac.at/webda/ c 2017 RAS, MNRAS 000, 000-000

http://irsa.ipac.caltech.edu/applications/wise/

http://www.cosmos.esa.int/web/herschel/science-archive/ c 2017 RAS, MNRAS 000, 000-000

Zavagno A. et al., 2007, A&A, 472, 835 c 2017 RAS, MNRAS 000, 000-000

ACKNOWLEDGEMENTS

We are thankful to the anonymous referee for valuable comments. NP acknowledges financial support from the Department of Science & Technology, INDIA, through INSPIRE faculty award IFA12-PH-36 at University of Delhi. The work at NCU is supported by the National Science Council through the grant no. 102-2119-M-008-001. This publication makes use of data from the Two Micron All Sky Survey (a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/ California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation), archival data obtained with the Spitzer Space Telescope (operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the NASA).


Ines Fayala

Oualid Limam

Ioannis Stefanou

EXPERIMENTAL AND NUMERICAL ANALYSIS OF REINFORCED STONE BLOCK MASONRY BEAMS USING GFRP REINFORCEMENT

Introduction

Natural stone structures represent the largest part of the construction heritage in the world such as bridges, civil and worship buildings or historical monuments. The natural stone is preferred for many reasons such as beauty, accessibility, hardness, durability, strength and sustainability. It is necessary to have a good understanding of the mechanical behaviour of stone structures. Their main characteristics are a high compressive strength and an almost null tensile strength due to the joints. Therefore, in historical structures, the use of stone is mostly restricted to members mainly working in compression (piers, arches, walls and vaults) [START_REF] Heyman | The stone skeleton: structural engineering of masonry architecture[END_REF]. Most historical buildings are built with masonry and it's well known that they are inadequately reinforced or unreinforced and pose a major seismic hazard in areas if they are not correctly designed. This heritage is accumulating damage due to deterioration of materials, repeated loading and exceptional events like earthquakes. Their preservation is a fundamental issue in the cultural life of modern societies by repairing and strengthening. In Tunisia, there are many masonry monuments damaged as for example, the Roman amphitheater of "El Jem" dating from the 3rd century which presents stone deterioration and many structural cracks in the vaults and the "Ribat" of Monastir damaged by the earthquake in October 2013 (Fig. 1) The reinforcement of masonry structures is one of the most frequently used practices in restoration of historical buildings in the purpose to enhance the resistance. It can be performed using steel bars, rings and/or composite materials [START_REF] De Felice | Assessment of the load-carrying capacity of multi-span masonry arch bridges using fibre beam elements[END_REF] especially in seismic areas [START_REF] Tao | Behaviour of a masonry arch bridge repaired using fibre-reinforced polymer composites[END_REF][START_REF] Omikrinemetalssi | Experimental study on reinforced stone beams[END_REF]. In the last two decades, composite materials, as FRP [START_REF] Kiss | FRP Strengthened Masonry Beams. Part I -Model[END_REF], have been increasingly considered for strengthening and repairing both modern and historic masonry constructions. FRP are excellent candidates for strengthening because of the high tensile strength they provide, their resistance in corrosion and in their easy handling. Several papers have addressed the strengthening of masonry with carbon and glass fibers reinforced composites [START_REF] Ehsani | Flexural behavior of masonry walls strengthened with composite fabrics[END_REF][START_REF] Schweglerg | Masonry construction strengthened with fiber composites in seismically endangered zones[END_REF][START_REF] Ehsani | Seismic retrofitting of masonry structures with composite materials: Research and field application[END_REF][START_REF] Triantafillou | Innovative strengthening of masonry monuments with composites[END_REF][START_REF] Schwegler | Earthquake resistance of masonry structures strengthened with fiber composites[END_REF][START_REF] Triantafillou | Strengthening of masonry structures using epoxy-bonded FRP laminates[END_REF] in the purpose to increase their ultimate capacity (strength and ductility). In the beginning, FRP were used for the reinforcement of concrete structures. FRP consist of high strength fibres embedded in a resin matrix. The fibres are usually Carbon (CFRP), Glass (GFRP) or Aramid (AFRP). They are available in sheets, strips, tendons, reinforcing bars. The available literature related to their use either as an internal reinforcement method or as an external strengthening method for reinforced concrete structural members is rich [START_REF] Theriault | Effects of FRP reinforcement ratio and concrete strength on flexural behavior of concrete beams[END_REF][START_REF] Abdalla | Evaluation of deflection in concrete members reinforced with FRP bars[END_REF][START_REF] Limam | RC two-way slabs strengthened with CFRP strips: experimental study and a limit analysis approach[END_REF][START_REF] Saikia | Strength and serviceability performance of beams reinforced with GFRP bars in flexure[END_REF]. The researchers carried out monotonic and cyclic tests on FRP reinforced concrete beams and investigated the effect of FRP reinforcement ratio and strength of concrete on the flexural resistance of the specimens. They concluded that the ultimate moment capacity of the beams increase as the reinforcement ratio increases and tested beams were considered to be safe in terms of deformability. Most research is done on masonry walls or concrete beams reinforced with industrial composite strips or wraps. Because of their high stiffness, industrial composites, when bonded to masonry produce under loading a strain incompatibility which induces a stress concentration and their premature debonding. We must focus on the specific features of the masonry structures because of the complexity of the monuments' geometry, cracks' causes and the importance of every intervention on historical structures. They can be subjected to a local tension under their weight or under horizontal inertial loads induced by a seismic action. When the tensile strength is reached, cracks appear in the planes perpendicular to the principal tensile stress direction. Starting from extension of approaches proposed for concrete structures, the objective of this paper is to present experimental and numerical tests performed on masonry beams reinforced with economic hand-made GFRP strips which are externally bonded on their tensile face. The externally bonded (EB) technique was preferred to the Near Surface Mounted (NSM) technique because in the latter technique, performing grooves to introduce reinforcements may cause cracking planes through the thickness of the structure. The used composites are not only a low-cost technology but they present also the mechanical advantage of strain compatibility with masonry with a Young's modulus close to that of masonry but a tensile strength many times higher than the masonry. The choice of spaced strips with predefined shapes presents an advantage relevant to the hygroscopic compatibility in comparison with the wraps. The aesthetic aspect can be a drawback but if it's compared to other alternative technologies, EB strips are in accordance with some ICOMOS recommendations (removability, minimum repair and respect of the original conception, durability and compatibility…) [START_REF] Lourenço | Analysis and restoration of ancient masonry structures: guidelines and examples[END_REF]. Furthermore, these composites are available and commonly used in the local small boats industry. So, in the purpose to assess the tensile performance of masonry strengthened by composite indirectly through the beams flexural testing, eight beams were constructed using stone masonry units and tested in four-point bending with an effective span of 0.824 m between the supports. These tests intend to reproduce artificially a tensile stress state in a masonry structure and to verify the efficiency of the EB composite strips strengthening technique to increase its bearing capacity by increasing the local tensile strength at its boundary. Load-displacement curves have been produced from recorded data, then a discussion of experimental and numerical results is proposed. The tests are not experimental prototypes of real masonry structures. Our final purpose is the validation of the proposed numerical model on the tested beams to be used for numerical analysis of ancient masonry structures. Numerical models can be used afterwards to compare the behaviors of strengthened and non-strengthened real masonry structures subjected to a local tensile stresses at their boundaries. Concerning the reinforcement of the structure showed in Fig. 1, GFRP strips can be used with adapted geometry and bonded externally to the masonry following the principal tensile stresses. In practice, cracks should be repaired first with epoxy injection.

Experimental program

Description of the beams

The tested specimens consisted of 8 stone masonry beams manufactured according to the sketch of Fig. 2. Their average length was 1.024 m with joints approximately 6 mm thick. All the tested beams have two courses of stone masonry units. Every beam has been made by 10 prismatic stone blocks (with the dimensions of 100 mm width, 100 mm height and 200 mm length) which were assembled using hydraulic lime mortar joints. Some of the beams were reinforced with GFRP strips at the traction side (Fig. 3). After the mortar hardening (28 days), the preparation of the surface for the reinforcement was carried out by sanding and dusting in order to ensure good adhesion between the strip and the stones. The GFRP strips were externally bonded to the masonry beams (see dry assembling in Fig. 3b) by using a structural thixotropic glue (sikadur_30), which is a mixture of epoxy resin and special fillers designed for gluing structural reinforcements. Seven days after, the glue reaches a good resistance according to the technical notice. After this period of time, beams were overturned with a kind of formwork in order to be tested in bending. The specimens are classified into four groups as shown in Tab.1: one group was not strengthened and it was kept as a reference and the rest three groups were reinforced externally on the bottom with GFRP strips of different thicknesses.

Materials properties

Experimental measures have been carried out to determine the Young's modulus, the compressive strength and the tensile strength of the stone and the mortar as well as the tensile properties of the composite reinforcement. The materials characteristics are summarized in Tab. 2 and Tab. 3 .

Masonry units

The masonry units are stone blocks originating from a quarry used for the restoration of "El Jem" amphitheater located in Monastir (Tunisia). The stone blocks were hewed into prismatic specimens with dimensions (100 x 100 x 200 mm) to build beams and (40 x 40 x 160 mm) for testing purposes. The type of rock is sandstone which is a sedimentary rock. According to the literature [START_REF] Chalhoub | Apports des méthodes d'homogénéisation numériques à la classification des massifs rocheux fracturés[END_REF][START_REF] Chalhoub | Massifs rocheux -Homogénéisation et classification numériques[END_REF][START_REF] Mestat | Lois de comportement des géomatériaux et modélisation par la méthode des éléments finis[END_REF], the Poisson's ratio varies from 0.2 to 0.35. The Young's modulus was determined for three cylindrical stone specimens (100 mm diameter and 200 mm height) by means of the longitudinal electronic extensometer as illustrated in Fig. 4. The extensometer used can measure the deformation of a cylindrical or prismatic specimen. The measuring rod was attached to the test specimen and after adjusting, stop levers were released so that the jaws became free to measure sensors displacements. Cyclic loading tests were carried out on the stone specimens under displacement control at a constant rate of 0.01 mm/sec by means of electromechanical compression machine. The maximum force applied is relative to the third of the material compressive strength in order to remain in the elastic range. The average value of Young's modulus was found equal to 17±0.5 GPa for the stone specimens. The Young's modulus was determined from the strain-stress curve after linear fitting of experimental data (see Fig. 5). The Poisson's ratio for the stone was assumed to be 0.2. Compressive and flexural strength tests were carried out on three prismatic stone specimens of dimensions 40 x 40 x 160 mm as shown in Fig. 6 in accordance with standards tests EN 196-1. The average strength was equal to 8.1±0.4 MPa, corresponding to the peak load under static monotonic loading) and 4.1±0.1 MPa (obtained by the elastic bending peak tensile stress) for compression and traction respectively (Tab.2).

Mortar

The choice of lime mortar is based on its use in almost all historical monuments. A lime to sand ratio (by volume) 1:3 was adopted for this work, according to literature [START_REF] Moropoulou | Composite materials in ancient structures[END_REF] which specifies that the lime mortar encountered in ancient structures has a binder to aggregate ratio ranging between 1:4 and 1:1. This ratio is also provided by the document "Specific use of natural hydraulic lime" that complies with specific recommendations of ICOMOS (International Council on Monuments and Sites) [START_REF] Labesse | Specific use of natural hydraulic lime[END_REF]. The lime used is an artificial hydraulic lime with strength class 10. A sand equivalent test was performed to verify the good quality of the used sand (SE = 84%). The mortar consistency was satisfactory in accordance with standards tests [START_REF]2: bulk sampling of mortars and preparation of test mortars[END_REF][START_REF]Methods of test for mortar for masonry -Part 3: determination of consistence of fresh mortar[END_REF] performed in order to determine the quantity of mixing water required. Compressive and flexural strength tests carried out on prismatic limebased mortar specimens (40 x 40 x 160 mm) gave respectively a mean strength equal to 6.2±0.3 MPa and 1.8±0.2 MPa as shown in Tab. 2. The determination of the Young's modulus was made by mean of the same longitudinal electronic extensometer already used previously and following the same procedure (see Fig. 7). The average value of Young's modulus found for mortar specimens is 1.2±0.1 GPa and the Poisson's ratio is assumed equal to 0.3.

GFRP strips

GFRP strips were used as reinforcement for stone masonry beams. There are three types of strips considered for this study: 2 layers, 3 layers and 4 layers strips. The manufacturing process consists in alternating resin and chopped glass strand mat (450 g/m 2 ), taking care to release the air bubbles by a roller for every layer (see Fig. 8). Glass support was employed as a working basis so the upper surface (which is the rough one) is the contact surface with the masonry beam. After three days of drying at room temperature, GFRP plates were cut, using a guillotine-shears, into strips (90 x 982 mm) for reinforcing and test specimens (25 x 250 mm) to determine tensile properties (see Fig. 9 for specimens and testing procedure) in accordance with ISO 527-4 [START_REF]Plastics -Determination of tensile properties -Part 4: Test conditions for isotropic and orthotropic fibre-reinforced plastic composites[END_REF] as depicted in Fig. 10. The tensile strength is measured in only one direction since the composite is considered as homogeneous material seeing that it was manufactured using a chopped glass strand mat. Tab. 3 summarizes the details of the tested FRP specimens corresponding to the curves drawn in Fig. 11 which shows an elastic-brittle behaviour of the composites.

Test setup, instrumentation and procedure

Four-points bending tests were performed on each beam until failure. The experimental setup including the loading and the support arrangement is illustrated in Fig. 12. The beams were simply supported by placing rollers in the middle of each end-block which gives an effective span of 0.824m. When overturned and placed on the supports, beams were, first, subjected naturally to bending under their own weight. Then, the four points bending load test was performed. So, the considered LVDT recorded only beams deflection when they were subjected to four points bending load. The load was applied via two other steel rollers positioned in the top of the beam in order to apply a line load at 362 mm from the edge; so that the distance between these two loading steel rollers is 300 mm. The beams were tested monotonically until failure by a displacement controlled electromechanical flexion machine at a constant rate of 0.01 mm/sec. The load was applied with a hydraulic jack with a capacity of 150 kN and distributed with a stiff metallic box girder. Data recording was started before the girder is set, so that the load induced by the girder is taken into account. The instrumentation was set up to measure the deflection at mid-span of the beam. A steel measure device was adopted to facilitate the mid-span deflection measurement. This device is simply placed on the upper side of the beam so that the supports coincide with the beam supports on the bottom. Two linear variable differential transducers (LVDTs) were used to measure the mid-span deflection and set on the both sides of the measure device.

Experimental results and discussion

To illustrate the global behaviour of the beams, the average load-deflection curves of the tested beams are presented in Fig. 13. A variance in the results is observed between two similar tested beams. The experimental scatter seems to be due to the manual manufacturing methodology of the masonry. In fact, for the tested beams, the peak load increases from 3 kN for the un-reinforced beams to a peak load ranging from 10 kN (for the beam B 2-1 strengthened with 2 layers) to 14 kN (for the beam B 4-2 strengthened with 4 layers). Furthermore, in order to visualize the experimental tendency, four average curves are depicted in Fig. 13 corresponding to two similar specimens each. As it can be noted in this figure, the experimental initial stiffness doesn't follow the expected increasing when the number of layers increases. Theoretically, when bonding thin strips having a Young's modulus E = 6 GPa close to the masonry one, the second moment of area should increase slightly which is followed by a small increasing of the initial stiffness. This small increasing is not detected automatically by the tests when the number of layers increases. However, the average curves (Fig. 13) show a tendency of the initial stiffness increasing within the strengthening and, as expected, the beams behave globally in a very similar way in the intial elastic part. In particular, Figure 14 shows the failure modes of masonry beams. Fig. 14a shows the failure mode of unreinforced beam: it is a bending mode which results in a vertical split at mid-span. Figures 14b, 14c, 14c' and14d show the failure mode of the GFRP reinforced beams which is a step-like, shear type crack that follows the masonry joints. After the midspan initial micro-cracks, a crack starts at the support, goes through the FRP-block interface and then it follows the mortar joints. The FRP strips remain intact. In the particular case of beam B 3-1, a crack starts from the support first by the composite debonding and then cracks propagates from the lower face in the span of constant moment (Fig. 14c). In all the other cases the debonding was observed at the end of testing. Concerning the debonding of the composite strip (Fig. 14e) from the stone surface, the observation of the interface after failure shows that cracking goes through the stone in the vicinity of the interface. Fig. 13 shows four diagrams of the average values of failure load for the beams B 1, B 2, B 3 and B 4 (load-deflection diagrams). It can be noted that experimental results show a linear behaviour for a little deflection below 0.1 mm and a non-linear behaviour above 0.1 mm with crack propagation. It can be noted also (see Tab.4) that the reinforcement with 4 layers FRP increases about 4 times the failure load and confirms the FRP contribution to increase the tensile stress of masonry structures. By comparing the failure load of beams group B 2 (beams reinforced with 2 layers) with B 4 (beams reinforced with 4 layers), we observe an increase of the failure load by about 30%. As presented in Fig. 13, a change of slope appears in curves B 2, B 3 and B 4 at a deflection between 0.2 mm and 0.4 mm. This decreasing of stiffness is the effect of micro-cracks in the masonry under local tension at midspan on the lower face. Above 0.4, the remaining part of the curves describes a second failure corresponding to propagating macro-cracks starting close to the support. So, the first failure to be considered for the modeling of the reinforcement corresponds to a state where the maximum tensile stress reaches the tensile strength in the masonry at midspan observed for a deflection equal to 0.2 mm.

Finite Element modeling

The objective of this section is to develop a Finite Element (FE) model in order to predict the structural behaviour of stone masonry structure reinforced with GFRP. The finite element model is shown in Fig. 15. The material properties used for the analysis were based on the results of laboratory tests shown previously in Tab.2.

Loading and boundary conditions of the FE model

The numerical analysis herein is a three dimensional-nonlinear finite element analysis of the experimental set-up by the means of the commercial FE program ABAQUS. Fig. 15 shows a typical three-dimensional solid model containing the geometry, boundary conditions, the load and the mesh used in this study. The loading and boundary conditions are designed to simulate the experimental behavior of the tested beams. The weight of the beam wasn't considered since the deflection recording began when the beam was already overturned and the deflection due to the weight wasn't recorded. During analysis, a vertical displacement is imposed at the upper side of the beam. The model is supported and restrained from out of the loading plan translation, using support (roller) lines. Full bond is assumed between FRP and masonry.

Constitutive model

Bricks and mortar are geomaterials, which are commonly described by Coulomb failure criterion. Nevertheless, the numerical treatment of Coulomb yield surfaces can provoke numerical problems. Therefore, for the numerical analysis performed in this section, the stone blocks and mortar joints are considered to obey to a Drucker-Prager (DP) criterion, with associate flow rule. This approach brings considerable advantages from the computational point of view due to its robustness and convergence. It is worth mentioning that both Drucker-Prager and Coulomb criteria might not be appropriate to describe the mechanical behavior of a geomaterial, in general. Plasticity theory itself might be also inadequate to take into account the detailed behavior of a geomaterial and especially in the shear-extension regime, i.e. at low confinement. However, they seem to be reasonable modeling choices given the few parameters needed to define these criteria and their evolution (e.g. hardening). Notice that Drucker-Prager is a criterion that it is often used in soil and rock mechanics at low confinements (see for example Davis & Salvadurai [START_REF] Davis | Plasticity in Geomechanics[END_REF]). For a further discussion on this topic we refer to I.Stefanou et al. [START_REF] Stefanou | Three dimensional homogenization of masonry structures with building blocks of finite strength: A closed form strength domain[END_REF]). The yield surface of DP is expressed as follows:

𝐹 = 𝑞 -𝑝 tan β -𝑑 Eq.( 1) Where 𝒑 = 𝟏 𝟑 𝒕𝒓𝝈 is the mean stress, 𝒒 = √ 𝟑 𝟐 𝒕𝒓 𝒔: 𝒔 the shear stress intensity, 𝒔 = 𝝈 -𝒑𝑰 is the deviatoric stress tensor, d and β are the DP criterion parameters, which can be expressed in terms of the Coulomb friction angle and cohesion (see I.Stefanou et al. [START_REF] Stefanou | Three dimensional homogenization of masonry structures with building blocks of finite strength: A closed form strength domain[END_REF]).

The choice of the appropriate model [START_REF] Abaqus | ABAQUS Analysis user's manual[END_REF] to be used depends largely on the analysis type, the kind of material, the experimental data available for calibration of the model parameters, and the range of mean stress values that the material is likely to experience. The accuracy with which the DP criterion can fit these test data is limited by the fact that it assumes linear dependence of the shear stress intensity with the mean stress. Moreover, numerical problems might arise in the case of the DP yield surface due to the non-differentiability of the failure surface F at p=0. In order to avoid such numerical problems and to better capture the mechanical behavior at shear-extensional regime (low mean stress) a regularized version of the DP criterion is often used in practice. In Figure 18 we present this regularized version of DP, which is often called hyperbolic DP in the literature. The yield function of the hyperbolic DP criterion is given by the following equation:

𝑭 = √𝒍 𝟎 𝟐 + 𝒒 𝟐 -𝒑 𝒕𝒂𝒏 𝜷 -𝒅 ′ Eq. (2)
Where

𝑙 0 = 𝑑′| 0 -𝑝 𝑡 | 0 𝑡𝑎𝑛 𝛽, l0 is equivalent to a cohesion, 𝑑 ′ | 0 is the initial value of 𝑑 ′ , 𝑝 𝑡 | 0 is the initial hydrostatic tension strength of the material  d' 2 𝑡𝑎𝑛𝛽
and d' is the cohesion of the material.

Materials properties for the FE analysis

The experimental results were used to define the model input parameters. The stone block and the mortar are modeled as isotropic materials, using respectively a Young's modulus of 17 GPa and 1.2 GPa, and respectively a Poisson's ratio of 0.20 and 0.30.

In the purpose to measure the influence of hydrostatic pressure on the properties and thus, to conclude on the yield criterion to be adopted for materials characterization, the previous compressive and tensile tests (Fig. 6) presented in section 2.2 were used leading to different states of hydrostatic pressure [START_REF] Radfar | Simulation of concrete cover separation failure in FRP plated RC beams[END_REF]. The idea of using compression tests and flexural tensile tests results from the works of Chataigner [START_REF] Chataigner | Conception et dimensionnement d'un ancrage de hauban plat en matériaux composites[END_REF] and Wang [START_REF]Towards an improved understanding of strength and anisotropy of cold compacted powder[END_REF]. Linear fittings were performed afterwards as shown in Fig. 16. This gives the identified cohesion and friction angle presented in Tab. 5. In F.E. modelling, DP material model is defined for the stone and the mortar using the friction angle β, initial tension strength pt|0 and hardening curves for different strain rates. For the latter, a light hardening was adopted starting from the yield tensile strength with a slight slope. Calculation parameters are summarized in Tab.5. Masonry is brittle in tension and ductile in compression. Furthermore, experimental tests show that the composite remains intact after beams failure, this is explained by the high tensile strength of composite, which is 370 times higher than the tensile strength of the mortar joint (Tab. 2). The experimental tests show also that the masonry beams strengthened with GFRP strips behave like the over reinforced concrete beams under positive bending where the lower reinforcement remains elastic and where the upper concrete exhibits ductility under compression. In the framework of prestressed and reinforced concrete modeling in the serviceability state, the crack opening is generally limited by a positive strain threshold. In the same context, our numerical analysis is stopped when the tensile stress reaches the tensile strength of masonry in mortar. The model was valid until the field of principal stress reaches the tensile strength of 0.26 MPa. This threshold corresponds to the first failure observed between 0.2 -0.4 mm. The GFRP are assumed elastic materials as the elastic yield is not reached here (since the FRP strips weren't affected when masonry beams failed). Typically, the fibers show no ductility, so the stressstrain behavior of most FRP's can be taken as linear elastic to failure [START_REF] Shrive | The use of fiber reinforced polymers to improve seismic resistance of masonry[END_REF]. Furthermore, the composite is assumed perfectly bonded to the masonry. This simplified hypothesis can describe the peeling-off failure mode through the analysis of stresses in masonry but in the vicinity of the composite. This failure mode corresponds to a debonding of the composite with a rock material which still bonded in it as observed in Fig. 14e. The masonry remains always the weaker element of the interface due to a higher strength of the epoxy.

Comparison between numerical and experimental results

Fig. 17 presents the total load as function of mid-span deflection for beams B 1 to B 4.The effect of plasticity appears through a global non linear structural behavior. It can be noted that the FRP reinforcement contributes to a slight increasing of the initial global stiffness, as expected. In the framework of the elastoplastic analysis, the effect of FRP is more noticeable on the bearing capacity. The analysis of the failure mode revealed that the weakest part of the structural system was the mortar joint where the tensile strength in the stone-mortar interface is reached. This is in accordance with the experimental tests. The structural behavior is stopped at a step corresponding to the maximum tensile stress in the stone-mortar interface reaching its limit of 0.26 MPa (see Tab.2). Because of the brittleness of masonry in tension, the process was stopped at this step. It is noted also that the longitudinal tensile strength of the composite strip is not reached in any of the simulations performed. Fig. 17 presents a comparison between numerical and experimental results. The discontinuous vertical line corresponds to the step at which the maximum of the field of principal tensile stresses defined in all joints reaches its maximum value of 0.26 MPa.

As presented in Fig. 17 the present numerical model reproduces the experimental tendency. It can be noted that the numerical analysis show an increasing of the initial stiffness within the strengthening.

Furthermore, the numerical analysis predicts an increasing of the numerical failure load from 3.8 kN for the un-reinforced beam to 9.1 kN for the beam strengthened with 2 layers and to 11 kN for the beam strengthened with 4 layers. The numerical failure load is defined as the applied load corresponding to a local tensile strength reached in masonry. For the strengthened beams, these failure loads remain in safety when comparing them to the experimental ones as depicted in Tab.6. This indicates that this simplified numerical approach is suitable to estimate the bearing capacity of the strengthened structures. Moreover, we notice an increase of the failure load when the reinforcement is increased. The experimental results of the failure deflection obtained for B 3 and B 4 can be explained as follow : the reinforced beam doesn't fail when the mortar reach its tensile strength because crack propagation continues through the stone block until failure. Fig. 18 presents the shear strain field. This strain is concentrated in the horizontal joint between blocks and the vertical joints near the supports. This is in agreement with the failure mode observed experimentally and depicted in figure 14 in which we can see cracks appearing in the aforementioned joints. Furthermore, numerical results are in agreement with the relative sliding observed between blocks (see Fig. 14b).

Conclusion

The above-discussed results reveal interesting aspects of using economic hand-made GFRP strips in the reinforcement of masonry structures. In this paper four samples of reinforced stone beams have been tested. Their structural behaviour has been investigated and the results of these investigations have shown that the idea of developing reinforced stone structures by composites in a similar way to reinforced concrete structures is promising.

A 3D nonlinear numerical model using the finite element commercial package, ABAQUS, has been developed to predict the flexural load capacity and structural stiffness of beams. Based on the work described in this paper, the following conclusions are drawn: -All strengthened beams exhibited a higher load capacity and a more ductile behavior, compared with the un-reinforced beams.

-The analysis of the failure mode revealed that the weakest part of the structural system was, as expected, the mortar joint.

-There was an acceptable agreement between experimental and numerical results.

-The Drucker-Prager criterion in his hyperbolic form was suitable for the FE modelling analysis of masonry structures. Concerning the design of masonry structures reinforced by FRP, the authors think that the present experiments give a reliable database for the calibration of a numerical model.