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la shows an extended bulee from which two spiral arms originate. | 1a shows an extended bulge from which two spiral arms originate. |
A small clongation is seen in the ceuter. that could be due to the presence of a sal bax (see below). | A small elongation is seen in the center, that could be due to the presence of a small bar (see below). |
Although this galaxy is classified as S(rx)b in the ΠΟ, MeLoeod Ricke (1995) already reported the preseuce of a bar with a 16 aresec radius; anc e-0.5 from their I& image. | Although this galaxy is classified as S(rs)b in the RC3, McLeod Rieke (1995) already reported the presence of a bar with a 16 arcsec radius, and $\epsilon$ =0.5 from their K image. |
It doesut show up in our shiarp-divided image in Fig. | It doesn't show up in our sharp-divided image in Fig. |
Eb. but is clearly detected in the plo of e and PA with radius (Fie. | 1b, but is clearly detected in the plot of $\epsilon$ and PA with radius (Fig. |
le) which is in agrecmen with Peletier et al. ( | 1e) which is in agreement with Peletier et al. ( |
1999). | 1999). |
The parameters we deduce for he bar from that plot are very simular ο those reported w MeLeod Rieke (1995). | The parameters we deduce for the bar from that plot are very similar to those reported by McLeod Rieke (1995). |
The sharp-divided nuage reveals the presence of a small bir. extending along PA =1397.. up to 7 arcsec Yon the ceuter. | The sharp-divided image reveals the presence of a small bar, extending along PA =, up to 7 arcsec from the center. |
It is not detected in the PA aud € plot as it is too weak to be apparent. | It is not detected in the PA and $\epsilon$ plot as it is too weak to be apparent. |
Other evidence for this clear bar is the curved dust pattern that surrounds the very central region in the broac baud HIST image (filter FGOGW) by Malkan et al. ( | Other evidence for this nuclear bar is the curved dust pattern that surrounds the very central region in the broad band HST image (filter F606W) by Malkan et al. ( |
1998). | 1998). |
The differcuce inage in Fie. | The difference image in Fig. |
1e shows that the overall fit is good except in the region where the arm contribution is irportant. | 1c shows that the overall fit is good except in the region where the arm contribution is important. |
The surface brightuess profile (iu agreement with that of MeLeod Rieke) is well ft by a | disk model. except iu the zone where the aru. contribution is iuportanut (Fies. | The surface brightness profile (in agreement with that of McLeod Rieke) is well fit by a $+$ disk model, except in the zone where the arm contribution is important (Figs. |
If aud le). | 1f and 1g). |
This galaxy shows rotation of the PA and twisting of the isophotes in the central region. | This galaxy shows rotation of the PA and twisting of the isophotes in the central region. |
Evidence for the existence of a bar inside the primary bar is preseuted in Fies. | Evidence for the existence of a bar inside the primary bar is presented in Figs. |
2b and 2e. where € shows two maxima or a rather constaut PA. | 2b and 2e, where $\epsilon$ shows two maxima for a rather constant PA. |
The inner bar is also evident in he broad band IST nuaee by Malkan et al. ( | The inner bar is also evident in the broad band HST image by Malkan et al. ( |
1998). | 1998). |
The region counecting the wo bars is visible iu Fig. | The region connecting the two bars is visible in Fig. |
2b as a thin curved elougatiou starting at the end of the iuncr (thicker) har. | 2b as a thin curved elongation starting at the end of the inner (thicker) bar. |
The difference nuage in Fie. | The difference image in Fig. |
2cM shows the two nested uw as well as the region where the spiral arms begin: the spiral arin to the north is wich brighter than its southern counterpart. | 2c shows the two nested bars as well as the region where the spiral arms begin; the spiral arm to the north is much brighter than its southern counterpart. |
Due to the bars and arms. the residuals are eh except in the very outer zoucs (Fie. | Due to the bars and arms, the residuals are high except in the very outer zones (Fig. |
28). aud the bulee | disk fit is not very good (Fig. | 2g), and the bulge + disk fit is not very good (Fig. |
2f). | 2f). |
A big bar (Fig. | A big bar (Fig. |
30) aud three spiral arms are detected. among which the north west arm is the brightest and most detached (Fig. | 3e) and three spiral arms are detected, among which the north west arm is the brightest and most detached (Fig. |
3a). | 3a). |
At laree scales. the image lasa somewhat triuigulaur shape. | At large scales, the image hasa somewhat triangular shape. |
A πα galaxy is secu 35 aresce to the south aloug PA=168".. but no redshift is available for it. | A small galaxy is seen 35 arcsec to the south along , but no redshift is available for it. |
The object Q223710305. (Lluchra ct al. | The object Q2237+0305 (Huchra et al. |
1985). comprises a source quasar at a redshift of +=1.695 that is eravitationally lensed by a foreground galaxy with 2=0.0394 producing 4 resolvable images with separations of Mo" | 1985) comprises a source quasar at a redshift of $z=1.695$ that is gravitationally lensed by a foreground galaxy with $z=0.0394$ producing 4 resolvable images with separations of $\sim 1''$. |
Each of the 4 images are observed through the galactic bulge. which has a microlensing optical depth in stars that is of order unity (e.g. Went Falco. 1988: Schneider et al. | Each of the 4 images are observed through the galactic bulge, which has a microlensing optical depth in stars that is of order unity (e.g. Kent Falco 1988; Schneider et al. |
LOSS: Schmidt. Webster Lewis 1998). | 1988; Schmidt, Webster Lewis 1998). |
In addition. the proximity of the lensing galaxy. means that the effective transverse velocity may be high. vielding an expected microlensing event time-scale significantly shorter than that of other Iensed quasars. | In addition, the proximity of the lensing galaxy means that the effective transverse velocity may be high, yielding an expected microlensing event time-scale significantly shorter than that of other lensed quasars. |
The combination of these considerations make (223703052() the ideal object [from which to study microlensing. | The combination of these considerations make Q2237+0305 the ideal object from which to study microlensing. |
Indeed. Q2237|0305 is the only object in. which. cosmological nmücrolensing has been directlv. confirmed. (Irwin et.al 1989: Corrigan et.al 1991: Wozniak et al. | Indeed, Q2237+0305 is the only object in which cosmological microlensing has been directly confirmed (Irwin et.al 1989; Corrigan et.al 1991; Wozniak et al. |
2000a.b). | 2000a,b). |
Initially. this confirmation came in the form of a ~0.2 magnitude brightening of image A with a rise-time of ~26 davs (Corrigan et al. | Initially, this confirmation came in the form of a $\sim$ 0.2 magnitude brightening of image A with a rise-time of $\sim26$ days (Corrigan et al. |
1901). | 1991). |
Wambseanss. Paczvuski Schneider (1990) found that. assumitig à galactic transverse velocity of ~GO0kmsec.. this rise. could. be explained by microlensing due to stellar masses of a source having dimensions much (<0.01E Z2) smaller than the microlens | Wambsganss, Paczynski Schneider (1990) found that, assuming a galactic transverse velocity of $\sim 600km\,sec^{-1}$, this rise could be explained by microlensing due to stellar masses of a source having dimensions much $<0.01\,ER$ ) smaller than the microlens |
QSOs and ensures that both type-1 and type-2 QSOs span a similar luminosity We note that. for QSOs in the redshift range 7=0.4-1.5. the 3" size of the SDSS fibers encloses regions as large as 16- kpe diameter. and therefore samples a significant portion of the host galaxy in which star formation can take place. | QSOs and ensures that both type-1 and type-2 QSOs span a similar luminosity We note that, for QSOs in the redshift range $z$ =0.4-1.5, the 3” size of the SDSS fibers encloses regions as large as 16-26 kpc diameter, and therefore samples a significant portion of the host galaxy in which star formation can take place. |
In Fig. | In Fig. |
5 we plot the [O II|/L,Ne V] ratio as a function of the measured X-ray column density. | \ref{oiinenh_sdss} we plot the [O II]/[Ne V] ratio as a function of the measured X-ray column density. |
The blue SDSS QSOs in the YO9 sample do show [O IIL]/[Ne V] ratios on average lower than type-2 QSOs. | The blue SDSS QSOs in the Y09 sample do show [O II]/[Ne V] ratios on average lower than type-2 QSOs. |
Some positive correlation. albeit with a large scatter. is Indeed seen between the [O II]/|[Ne V] ratio and the absorbing column density N;;. | Some positive correlation, albeit with a large scatter, is indeed seen between the [O II]/[Ne V] ratio and the absorbing column density $N_H$. |
When considering those objects with observed [O II]/[Ne V] >4. we found that only 2 out of 12 are not obscured. and half of them (6 objects) are likely obscured by CT absorption. | When considering those objects with observed [O II]/[Ne V] $>4$, we found that only 2 out of 12 are not obscured, and half of them (6 objects) are likely obscured by CT absorption. |
Conversely. there are no objects with Nj;>I0? among those with |O I]/[Ne V] <I. | Conversely, there are no objects with $N_H>10^{23}$ among those with [O II]/[Ne V] $<1$. |
The [ο IL/[Ne V] ratios measured on the SDSS type-2 and type-| QSO composites by and were also considered and found to bein good agreement with the averages measured in this work for obscured and unobscured QSOs. respectively (see Fig. 5)). | The [O II]/[Ne V] ratios measured on the SDSS type-2 and type-1 QSO composites by and were also considered and found to be in good agreement with the averages measured in this work for obscured and unobscured QSOs, respectively (see Fig. \ref{oiinenh_sdss}) ). |
We tried to compute the significance of the correlation between the [O II]/[Ne V| ratio and the logarithm of the column density. | We tried to compute the significance of the correlation between the [O II]/[Ne V] ratio and the logarithm of the column density. |
We note that it is difficult to deal with objects which are either unobscured or CT candidates. because they cannot be treated statistically as proper upper or lower limits on Nj;. since the gas column density can plausibly vary only within a bounded range (1.8. it cannot be zero or infinite). | We note that it is difficult to deal with objects which are either unobscured or CT candidates, because they cannot be treated statistically as proper upper or lower limits on $N_H$, since the gas column density can plausibly vary only within a bounded range (i.e. it cannot be zero or infinite). |
For simplicity we therefore assumed logN;; 220 for unobscured objects and logN;;224 for CT candidates. respectively (see Fig. 5)). | For simplicity we therefore assumed $N_H$ =20 for unobscured objects and $N_H$ =24 for CT candidates, respectively (see Fig. \ref{oiinenh_sdss}) ). |
The presence of a correlation has been estimated through the software package(?).. using the generalized Kendall’s r and the Spearman's p correlation tests. | The presence of a correlation has been estimated through the software package, using the generalized Kendall's $\tau$ and the Spearman's $\rho$ correlation tests. |
We found that the probability that the correlation is not present is only 2x1077 and 1x107. respectively. | We found that the probability that the correlation is not present is only $2\times10^{-4}$ and $1\times10^{-4}$, respectively. |
If the [O II] emission measured in type-2 QSOs is interpreted as entirely due to star formation. the median [O II luminosities of the [O III]- and [Ne V|-selected samples woulc correspond to star formation rates of=100 and =2007 respectively (using the relation by 2)). | If the [O II] emission measured in type-2 QSOs is interpreted as entirely due to star formation, the median [O II] luminosities of the [O III]- and [Ne V]-selected samples would correspond to star formation rates of $\approx 100$ and $\approx 200\; M_\odot$ /yr, respectively (using the relation by ). |
AGNThese values coulc decrease by up to a factor of ~2 if the contribution to the [OIL] emission ts significant(?). | These values could decrease by up to a factor of $\sim 2$ if the AGN contribution to the [OII] emission is significant. |
. This finding ts consistent with the expectations from the AGN evolutionary sequence outlinec above. | This finding is consistent with the expectations from the AGN evolutionary sequence outlined above. |
We have presented a diagnostic diagram to identify heavily obscured. Compton-Thick AGN candidates at z~| based on the ratio between the 2-10 keV flux and the [Ne V]3426 emission line flux (X/NeV). | We have presented a diagnostic diagram to identify heavily obscured, Compton-Thick AGN candidates at $z\sim 1$ based on the ratio between the 2-10 keV flux and the [Ne V]3426 emission line flux (X/NeV). |
The diagnostic was calibrated on a sample of 74 local Seyfert galaxies and then applied to populations of type-1 and type-2 QSOs at different redshifts (from z0.1 to z=1.5) selected from the SDSS. | The diagnostic was calibrated on a sample of 74 local Seyfert galaxies and then applied to populations of type-1 and type-2 QSOs at different redshifts (from $z\sim 0.1$ to $z=1.5$ ) selected from the SDSS. |
The main results obtained in this work can be summarized as follows. | The main results obtained in this work can be summarized as follows. |
e The observed X/NeV ratio is found to decrease with increasing absorption: the mean. X/NeV ratio for unobscured Seyferts is about 400. about of local Seyferts with X/NeV<100 are obscured by column densities above πο. and essentially all objects with observed X/NeV «I5 are Compton-Thick. | $\bullet$ The observed X/NeV ratio is found to decrease with increasing absorption: the mean X/NeV ratio for unobscured Seyferts is about 400, about of local Seyferts with $<100$ are obscured by column densities above $10^{23}$ and essentially all objects with observed X/NeV $<15$ are Compton-Thick. |
ο We considered a sample of 83 blue type-1 QSOs and 2] [O IlI|-selected type-2 QSOs in the SDSS which have been observed in the X-rays and show significant [Ne V] detection. | $\bullet$ We considered a sample of 83 blue type-1 QSOs and 21 [O III]-selected type-2 QSOs in the SDSS which have been observed in the X-rays and show significant [Ne V] detection. |
It was verified that they follow the same X/NeV vs X-ray absorption trend which ts observed for local Seyferts. | It was verified that they follow the same X/NeV vs X-ray absorption trend which is observed for local Seyferts. |
Furthermore. SDSS type-2 QSOs classified either as Compton-Thick or Compton-Thin on the basis of their X/OIL ratio. would have been mostly classified in the same way based on the X/NeV ratio. | Furthermore, SDSS type-2 QSOs classified either as Compton-Thick or Compton-Thin on the basis of their X/OIII ratio, would have been mostly classified in the same way based on the X/NeV ratio. |
e The X/NeV diagnostic was used to investigate the obscuration of 9 SDSS obscured QSOs in the redshift range zc[0.85--1.31]. which is not accessible through [O III] selection. | $\bullet$ The X/NeV diagnostic was used to investigate the obscuration of 9 SDSS obscured QSOs in the redshift range $z=[0.85-1.31]$, which is not accessible through [O III] selection. |
The 9 objects were selected by means of their prominent [Ne V|3426 line (EW> 4A)). and ssnapshot observations for 8 of them were obtained (one object is from the archive). | The 9 objects were selected by means of their prominent [Ne V]3426 line $EW>4$ ), and snapshot observations for 8 of them were obtained (one object is from the archive). |
Based on the X/NeV ratio. complemented by X-ray spectral analysis. only 2 objects appear good Compton-Thick QSO candidates. | Based on the X/NeV ratio, complemented by X-ray spectral analysis, only 2 objects appear good Compton-Thick QSO candidates. |
However. when considering the 4. genuine. narrow-line objects only (FWHM of the Mell line <2000 km s! ). the efficiency in selecting Compton-Thick QSOs through the [Ne V] line is about (2/4). which is more similar. despite the large uncertainties. to what is achieved with [O LIE] selection(60- | However, when considering the 4 genuine narrow-line objects only (FWHM of the MgII line $\lesssim 2000$ km $s^{-1}$ ), the efficiency in selecting Compton-Thick QSOs through the [Ne V] line is about (2/4), which is more similar, despite the large uncertainties, to what is achieved with [O III] selection; ). |
70%:: ?)). e We verified that neither extinction nor anisotropy corrections on the [νο V] emission would affect. our conclusions and that the X/NeV diagnostic is therefore a good method to identify clean. despite not complete. samples of heavily obscured AGN. | $\bullet$ We verified that neither extinction nor anisotropy corrections on the [Ne V] emission would affect our conclusions and that the X/NeV diagnostic is therefore a good method to identify clean, despite not complete, samples of heavily obscured AGN. |
We discussed the possibility of applying the X/NeVdliagnostic to objects in sky areas with deep optical spectroscopy and X-ray coverage. | We discussed the possibility of applying the X/NeV diagnostic to objects in sky areas with deep optical spectroscopy and X-ray coverage. |
This will allow to identify Compton-Thick Seyferts at z1. re. those objects which are thought to be responsible for a large fraction of the “MISSIig" X-ray background. | This will allow to identify Compton-Thick Seyferts at $z\sim 1$, i.e. those objects which are thought to be responsible for a large fraction of the “missing" X-ray background. |
ο Finally. the optical emission line properties of [Ne V|-selected QSOs were compared with those of other SDSS populations of obscured and unobscured QSOs. | $\bullet$ Finally, the optical emission line properties of [Ne V]-selected QSOs were compared with those of other SDSS populations of obscured and unobscured QSOs. |
By restricting the analysis to objects in the same redshift (and luminosity) range z-[0.4-1.5]. we found evidence that the ratio between the [O I[]3727 and [Ne V|3426 luminosity mereases with obscuration. | By restricting the analysis to objects in the same redshift (and luminosity) range $z$ =[0.4-1.5], we found evidence that the ratio between the [O II]3727 and [Ne V]3426 luminosity increases with obscuration. |
This correlation is interpreted as evidence of | This correlation is interpreted as evidence of |
composition of the mixture (e.g. Pontoppidan et al. 2003)). | composition of the mixture (e.g. Pontoppidan et al.\cite{ponto_ch3oh}) ). |
Hence, we derived upper limits to the CH3OH column density of the 3.53 um feature. | Hence, we derived upper limits to the $_3$ OH column density of the 3.53 $\mu$ m feature. |
The upper limits to the CH3OH column densities are 1.2x1015 cm"? relative to H5O), 1.0x1015 cm? (42%)), and 3.1x107 cm? (larger than the H5O column) towards L1527, IRAS04302, and HK Tau, respectively. | The upper limits to the $_3$ OH column densities are $1.2 \times 10^{18}$ $^{-2}$ relative to $_2$ O), $1.0 \times 10^{18}$ $^{-2}$ ), and $3.1 \times 10^{17}$ $^{-2}$ (larger than the $_2$ O column) towards L1527, IRAS04302, and HK Tau, respectively. |
In some objects, such as IRAS04302, we can see absorption around 4.1 um, where the OD stretching mode of HDO is observed in the laboratory (Dartois et al. 2003)). | In some objects, such as IRAS04302, we can see absorption around 4.1 $\mu$ m, where the OD stretching mode of HDO is observed in the laboratory (Dartois et al. \cite{dartois03}) ). |
Since the detection of HDO ice is very rare (Teixeira et al. 1999)) | Since the detection of HDO ice is very rare (Teixeira et al. \cite{teixeira99}) ) |
and has not been confidently confirmed, we checked the response function carefully and confirmed that the ~4.1 jum feature is not caused by an artifact in the response function. | and has not been confidently confirmed, we checked the response function carefully and confirmed that the $\sim 4.1$ $\mu$ m feature is not caused by an artifact in the response function. |
Figure 5 shows the spectra of this wavelength region towards L1527, IRC-L1041-2, IRAS04302, and HV Tau, which are fitted with a model of the amorphous HDO feature at 10 K: a Gaussian profile peaking at 4.07 um with a full-width half maximum (FWHM) of 0.2 um (solid lines). | Figure \ref{HDO} shows the spectra of this wavelength region towards L1527, IRC-L1041-2, IRAS04302, and HV Tau, which are fitted with a model of the amorphous HDO feature at 10 K: a Gaussian profile peaking at 4.07 $\mu$ m with a full-width half maximum (FWHM) of 0.2 $\mu$ m (solid lines). |
The spectrum of L1527 is not fitted well with this Gaussian; the absorption has a peak at longer wavelength (~4.13 wm) and a narrower band width, which resembles an annealed, rather than an amorphous, HDO feature (Dartois et al. 2003)). | The spectrum of L1527 is not fitted well with this Gaussian; the absorption has a peak at longer wavelength $\sim 4.13$ $\mu$ m) and a narrower band width, which resembles an annealed, rather than an amorphous, HDO feature (Dartois et al. \cite{dartois03}) ). |
We fitted the L1527 spectrum with a Gaussian peaking at 4.13 um and FWHM of 0.1 um (dashed line in Figure 5)). | We fitted the L1527 spectrum with a Gaussian peaking at 4.13 $\mu$ m and FWHM of 0.1 $\mu$ m (dashed line in Figure \ref{HDO}) ). |
Although IRAS04302 has the deepest and smoothest absorption around 4.1 um, this result should be taken with caution. | Although IRAS04302 has the deepest and smoothest absorption around 4.1 $\mu$ m, this result should be taken with caution. |
The spectrum also shows a broad absorption at 4.5—4.6 pm, which cannot be fitted well by the absorptions of CO and XCN (see §4.6). | The spectrum also shows a broad absorption at $4.5-4.6$ $\mu$ m, which cannot be fitted well by the absorptions of CO and XCN (see 4.6). |
These two absorptions (4.1 um and 4.5—4.6 pm) could be related; there could be an alternative explanation, rather than HDO, CO, and XCN. | These two absorptions (4.1 $\mu$ m and $4.5-4.6$ $\mu$ m) could be related; there could be an alternative explanation, rather than HDO, CO, and XCN. |
The features themselves are, however, robust. | The features themselves are, however, robust. |
We have two independent data sets of IRAS04302, and the two broad absorptions appear in both data sets. | We have two independent data sets of IRAS04302, and the two broad absorptions appear in both data sets. |
We integrated the fitted spectra to derive the column density of HDO (Table 2)). | We integrated the fitted spectra to derive the column density of HDO (Table \ref{column}) ). |
The band strength was assumed to be 4.3x cm molecule! (Dartois et al. 2003)). | The band strength was assumed to be $4.3 \times 10^{-17}$ cm $^{-1}$ (Dartois et al. \cite{dartois03}) ). |
Considering the small bumps and hollows that deviate from the HDO feature and the above discussion of IRAS04302 spectrum, the HDO column densities should be interpretedwith caution. | Considering the small bumps and hollows that deviate from the HDO feature and the above discussion of IRAS04302 spectrum, the HDO column densities should be interpretedwith caution. |
However, at facevalue, the HDO/H20 ratio ranges from 2 (L1527) to 22 (IRAS04302). | However, at facevalue, the $_2$ O ratio ranges from 2 (L1527) to 22 (IRAS04302). |
Except in the case of L1527, the ratios are much higher than those obtained in the previous observations and theoretical works: HDO/H5O «x 3 (Dartois et al. 2003,, | Except in the case of L1527, the ratios are much higher than those obtained in the previous observations and theoretical works: $_2$ O $\le$ 3 (Dartois et al. \cite{dartois03}, , |
Parise et al. 2003,, | Parise et al. \cite{parise03}, , |
Parise et al. 2005,, | Parise et al. \cite{parise05}, , |
Aikawa et al. 2005)). | Aikawa et al. \cite{aikawa05}) ). |
tum Land un) am assuming an isothermal wind (ds/dq = 0)ts(48) For the isothermal CAK75 stellar wind I implement a set of parameters corresponding to an O5É star. namely: and I use the line force parameters also used by CAIT5. namely: To study the existence of steady solutions. I first consider the function. | and in turn Since I am assuming an isothermal wind $ds/dq = 0$ ) For the isothermal CAK75 stellar wind I implement a set of parameters corresponding to an O5f star, namely: and I use the line force parameters also used by CAK75, namely: To study the existence of steady solutions, I first consider the $h$ function. |
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