ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
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.

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).

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).

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.

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).

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).

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).

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.

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.

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

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.

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.

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.