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At this time the density is relatively high: our €O column density is 6107em7. while Lerpin Cernicharo (2000) derive a value of 7107em> and our LCN is 2.51011em=~ onlv. slightly over a factor of 4 less than the observed one: our LINC is. however. over 3 orders of magnitude underabundant with respect to observations.
At this time the density is relatively high: our CO column density is $6\times 10^{18}~{\rm cm^{-2}}$, while Herpin Cernicharo (2000) derive a value of $7\times 10^{18}~{\rm cm^{-2}}$ and our HCN is $2.5\times 10^{17}~{\rm cm^{-2}}$, only slightly over a factor of 4 less than the observed one; our HNC is, however, over 3 orders of magnitude underabundant with respect to observations.
Given the dillerences in the physical parameters adopted. any further comparison would. be of limited value.
Given the differences in the physical parameters adopted, any further comparison would be of limited value.
Lt should also be noted that following the ISO discovery of benzene in CRLGIS (Cernicharo et 2001b). Woocs οἱ al. (
It should also be noted that following the ISO discovery of benzene in CRL618 (Cernicharo et 2001b), Woods et al. (
2002) presented a chemical model of this proto-planetary nebula.
2002) presented a chemical model of this proto-planetary nebula.
They show that the physical conditions of 1015 are such to encourage an ellicient. production of benzene.
They show that the physical conditions of CRL618 are such to encourage an efficient production of benzene.
llowever. similarly to Cernicharo et (2001b) results. they model the dense inner torus ancl therefore we are unable to make a direct comparison with their chemical moclel.
However, similarly to Cernicharo et (2001b) results, they model the dense inner torus and therefore we are unable to make a direct comparison with their chemical model.
The voung PN NGC 7027 is molecule rich with a large number of species identified within it.
The young PN NGC 7027 is molecule rich with a large number of species identified within it.
Llasegawa&Kwok(2001) have observed several species and collated the results of previous work.
\citet{hasegawa&kwok01} have observed several species and collated the results of previous work.
In their table 3. they list measured fractional abundances.
In their table 3, they list measured fractional abundances.
Lf the emission. that. Llascgawa&Kwok(2001). detected. is regarded. as originating from a laree collection of clumps then their results can be directly compared with ours in Table 1.
If the emission that \citet{hasegawa&kwok01} detected is regarded as originating from a large collection of clumps then their results can be directly compared with ours in Table 1.
Adopting an age of around 1000 vr for the nebula and comparing with our results gives the following: CN is predicted to have a fractional abundance of 910+? and llasegawa&νοκ(2001) measure 1.7107: HCN is 610 (predicted): verses 1.210.7 (observed): and Collis 610. ‘vs Ld.107.
Adopting an age of around 1000 yr for the nebula and comparing with our results gives the following: CN is predicted to have a fractional abundance of $9\times 10^{-10}$ and \citet{hasegawa&kwok01} measure $1.7\times 10^{-8}$; HCN is $6\times 10^{-5}$ (predicted) verses $>1.2\times 10^{-9}$ (observed); and ${\rm C_2H}$ is $6\times 10^{-7}$ vs $1.1\times 10^{-8}$.
Given the simplicity of the model and. the inexactness of the comparison to be within around one order of magnitude of the observational results. for several species is) very encouraging (the exact age of the nebula or the filling factor of the clumps can casily vield uncertainties of an order of maenituce).
Given the simplicity of the model and the inexactness of the comparison to be within around one order of magnitude of the observational results for several species is very encouraging (the exact age of the nebula or the filling factor of the clumps can easily yield uncertainties of an order of magnitude).
To compare the results for the later stages with those of the Helix. we can use the last columns of Table 1 with the results of Bachilleretal.(1997).
To compare the results for the later stages with those of the Helix, we can use the last columns of Table 1 with the results of \citet{bachiller.et.al97}.
.. The mocel results for this evolved PN do not agree nearly so well with the observations as for NGC7027 and. CRLGIS.
The model results for this evolved PN do not agree nearly so well with the observations as for NGC7027 and CRL618.
Some results are not inconsistent with the observations: the moclel predictions do not exceed. the upper limits established. for
Some results are not inconsistent with the observations: the model predictions do not exceed the upper limits established for
units of Gaussian c. that the source is not spurious (Le. a background fluctuation).
units of Gaussian $\sigma$, that the source is not spurious (i.e. a background fluctuation).
This probability is the result of numerical simulations of random fields (without. sources). repeated for different wavelet transform scales; it takes into account the number of background peaks mismatched as sources by the detection algorithm. for a given signal to noise ratio threshold in wavelet space.
This probability is the result of numerical simulations of random fields (without sources), repeated for different wavelet transform scales; it takes into account the number of background peaks mismatched as sources by the detection algorithm, for a given signal to noise ratio threshold in wavelet space.
The mean number of spurious sources expected over an image is kept constant (0.4 spurious detections per field). varying the detection threshold.
The mean number of spurious sources expected over an image is kept constant (0.4 spurious detections per field), varying the detection threshold.
That is. the detectioi threshold depends on the mean background counts value on the image and on the scale at which the search is performed (see Lazzati et al. 1999)).
That is, the detection threshold depends on the mean background counts value on the image and on the scale at which the search is performed (see Lazzati et al. \cite{lazz1999}) ).
The mean threshold value in the overall catalog is ~4.20.
The mean threshold value in the overall catalog is $\sim 4.2\sigma$.
In principle. the error for each parameter found with the wavelet transform algorithm can be estimated from the covariance matrix.
In principle, the error for each parameter found with the wavelet transform algorithm can be estimated from the covariance matrix.
In practice. such errors are reliable only when the number of source and background counts are sufficiently high (z2x107 counts per pixel).
In practice, such errors are reliable only when the number of source and background counts are sufficiently high $\gtrsim 2 \times 10^{-2}$ counts per pixel).
Otherwise. the distribution of the wavelet transform coefficients becomes Poissonian rather than Gaussian (in the lowest scales).
Otherwise, the distribution of the wavelet transform coefficients becomes Poissonian rather than Gaussian (in the lowest scales).
If this Is the case. a better estimate of the errors can be given by means of basic statistics (such as the error on the number of counts Nis VN. for more details see Lazzati et al. 1999)).
If this is the case, a better estimate of the errors can be given by means of basic statistics (such as the error on the number of counts $N$ is $\sqrt{N}$, for more details see Lazzati et al. \cite{lazz1999}) ).
The errors reported here and in the catalog are the maximum between the covariance matrix and the statistical estimates by Lazzati et al. (1999)),
The errors reported here and in the catalog are the maximum between the covariance matrix and the statistical estimates by Lazzati et al. \cite{lazz1999}) ).
There is an additional source of errors on the absolute position determination of HRI sourees. the uncertainty on the boresight correction.
There is an additional source of errors on the absolute position determination of HRI sources, the uncertainty on the boresight correction.
If the alignment between the node and the telescope optical axis is perfect. on-axis Images are exactly in the center of the instrument.
If the alignment between the node and the telescope optical axis is perfect, on-axis images are exactly in the center of the instrument.
However. in the real situation there is always a finite misalignment that needs to be corrected.
However, in the real situation there is always a finite misalignment that needs to be corrected.
Unfortunately. uncertainties i the aspect solution (that describes the orientation of the telescope as a function of time) and errors in the alignment between the star trackers and the ROSAT X-ray Telescope introduce uncertainties in the boresight correction. of variable size for each observation.
Unfortunately, uncertainties in the aspect solution (that describes the orientation of the telescope as a function of time) and errors in the alignment between the star trackers and the ROSAT X-ray Telescope introduce uncertainties in the boresight correction, of variable size for each observation.
In practice. the systematic offset between accurately known optical positions and X-ray positions can be used to evaluate the extent of boresight correction uncertainties.
In practice, the systematic offset between accurately known optical positions and X-ray positions can be used to evaluate the extent of boresight correction uncertainties.
For the ROSAT HRI. the offset can be as large as 10". even if in most cases it will be much less (David et al. 1998)).
For the ROSAT HRI, the offset can be as large as $10''$, even if in most cases it will be much less (David et al. \cite{david1998}) ).
The 10 value is conservatively assumed as a fiducial value for the boresight offset.
The $10''$ value is conservatively assumed as a fiducial value for the boresight offset.
Usually. the statistical errors on the position from the detection algorithm. calculated as in the previous section. are much less than 10” and therefore can be neglected in a first approximation.
Usually, the statistical errors on the position from the detection algorithm, calculated as in the previous section, are much less than $10''$ and therefore can be neglected in a first approximation.
As each pointing is treated separately. not all of the catalog entries correspond to independent sources.
As each pointing is treated separately, not all of the catalog entries correspond to independent sources.
An estimate of the overall number of independent sources can be given compressing the catalog in a 10” radius (again. the fiducial error value indicated by the boresight uncertainty).
An estimate of the overall number of independent sources can be given compressing the catalog in a $10''$ radius (again, the fiducial error value indicated by the boresight uncertainty).
This procedure selects only a source for each 10” cone radius. on the basis of an autocorrelation of the position.
This procedure selects only a source for each $10''$ cone radius, on the basis of an autocorrelation of the position.
With this procedure. 20433 sources are left.
With this procedure, 20433 sources are left.
Obviously. the compression brings to the loss of sources truly close to each other.
Obviously, the compression brings to the loss of sources truly close to each other.
In the following. we will largely use the term to indicate a wavelength passband different from X-rays.
In the following, we will largely use the term to indicate a wavelength passband different from X-rays.
The limiting flux of a given catalog in a given band will be generically denominated fj.
The limiting flux of a given catalog in a given band will be generically denominated $\rm f_{offband}$.
For each of the BMW-HRI entries. cross correlations with GSC2 (McLean et al. 2005).
For each of the BMW-HRI entries, cross correlations with GSC2 (McLean et al. \cite{mclean2005}) ),
2MASS Second Data Release (Kleinmann et al. 1994)).
2MASS Second Data Release (Kleinmann et al. \cite{klein1994}) ),
IRAS Point Source Catalog (Beichman et al. 1988))
IRAS Point Source Catalog (Beichman et al. \cite{beich1988}) )
and FIRST Survey Catalog (White et al. 1997))
and FIRST Survey Catalog (White et al. \cite{white1997}) )
were performed by Panzera et al. (2003)).
were performed by Panzera et al. \cite{panz2003}) ).
Off-band catalog properties are summarized in Table Ι..
Off-band catalog properties are summarized in Table \ref{crossid}.
The adopted radius for the cross-correlation is 10” (see Section 2.1). assumed as positional X-ray uncertainty.
The adopted radius for the cross-correlation is $10''$ (see Section 2.1), assumed as positional X-ray uncertainty.
In fact. the positional uncertainties for GSC2. 2MASS SDR and FIRST catalogs are <0.5” (3c. MeLean et al. 2005)).
In fact, the positional uncertainties for GSC2, 2MASS SDR and FIRST catalogs are $<0.5''$ $3\sigma$, McLean et al. \cite{mclean2005}) ),
<0.5” (lo. see the 2MASS documentation for an extensive discussion). «0.5” (90%, MeMahon et al. 2001))
$<0.5''$ $1\sigma$, see the 2MASS documentation for an extensive discussion), $<0.5''$ $90\%$, McMahon et al. \cite{mcmahon2001}) )
respectively. so that they can be safely neglected for cross-identification purposes.
respectively, so that they can be safely neglected for cross-identification purposes.
An exception ts the case of IRAS PSC. for which the positional accuracy varies with source size. brightness and spectral shape and it is different in different directions. but it is usually better than 20” (see Beichman et al. 1988)).
An exception is the case of IRAS PSC, for which the positional accuracy varies with source size, brightness and spectral shape and it is different in different directions, but it is usually better than $20''$ (see Beichman et al. \cite{beich1988}) ),
so this last value has been used as cross-correlation radius with this catalog.
so this last value has been used as cross-correlation radius with this catalog.
Note that: à) when two or more entries in the correlating catalog are found. only the brightest source parameters are reported (but in any case the number of found sources is given). b) the sky coverage of the off-band catalogs is usually not complete (see again Table 1))
Note that: a) when two or more entries in the correlating catalog are found, only the brightest source parameters are reported (but in any case the number of found sources is given), b) the sky coverage of the off-band catalogs is usually not complete (see again Table \ref{crossid}) ).
In particular. the preliminary. unpublished version of the GSC2 (GSC2.3) catalog used in. Panzera et al. (2003))
In particular, the preliminary, unpublished version of the GSC2 (GSC2.3) catalog used in Panzera et al. \cite{panz2003}) )
lacks coverage of the zones where a bright source caused an overexposure of the Schmidt plates.
lacks coverage of the zones where a bright source caused an overexposure of the Schmidt plates.
For the 2MASS and FIRST. more complete catalogs are now available (2MASS All Sky Data Release. FIRST Survey Catalog O3Aprl1 Version).
For the 2MASS and FIRST, more complete catalogs are now available (2MASS All Sky Data Release, FIRST Survey Catalog 03Apr11 Version).
Our final list has been checked with them.
Our final list has been checked with them.
The first sample selection was made on the basis of the following criteria: The reported numbers are catalog entries. Le. no compression was applied to eliminate multiple detections (see Section 3.24 below).
The first sample selection was made on the basis of the following criteria: The reported numbers are catalog entries, i.e. no compression was applied to eliminate multiple detections (see Section 3.2 below).
The fx we used is computed from the
The $\rm f_{X}$ we used is computed from the
stars (columns (6) and (8) in. Table 5) versus the cluster metallicity [Fe/H],.
stars (columns (6) and (8) in Table 5) versus the cluster metallicity $_K$.
The linear regression over the entire sample (solid line) yields a slope of 0.2040.06 mag dex. regardless of the adopted mixing-length parameter. while the zero-point of the relation changes from 0.9450.10 mag to 0.8240.10 mag with //H,,=1.5 and 2.0. respectively.
The linear regression over the entire sample (solid line) yields a slope of $\pm$ 0.06 mag $^{-1}$, regardless of the adopted mixing-length parameter, while the zero-point of the relation changes from $\pm$ 0.10 mag to $\pm$ 0.10 mag with $l/H_p$ =1.5 and 2.0, respectively.
However. the data given in Table 5 clearly shows that at constant metal content the RR ο luminosity depends on the cluster HB type: e.g.. the variables in NGC 7089 (HB=+0.96) are ~ 0.2 mag brighter than those in IC 4499, NGC 6934 and NGC 3201. which show a HB type from HB=+0.08 to +0.25. yet all these clusters have nearly the same metallicity.
However, the data given in Table 5 clearly shows that at constant metal content the RR $_{ab}$ luminosity depends on the cluster HB type: e.g., the variables in NGC 7089 (HB=+0.96) are $\sim$ 0.2 mag brighter than those in IC 4499, NGC 6934 and NGC 3201, which show a HB type from HB=+0.08 to +0.25, yet all these clusters have nearly the same metallicity.
This result is not a novelty since theoretical (see Paper IV and references therein) and observational studies (Lee Carney 1999: Clement Shelton. 1999; Alves. Bond Onken 2001) have already suggested that the RR Lyrae absolute magnitude depends on the cluster HB morphology and metal content.
This result is not a novelty since theoretical (see Paper IV and references therein) and observational studies (Lee Carney 1999; Clement Shelton 1999; Alves, Bond Onken 2001) have already suggested that the RR Lyrae absolute magnitude depends on the cluster HB morphology and metal content.
The comparison with field RR,, stars with [Fe/H]2.1.0 is shown in Fig.
The comparison with field $_{ab}$ stars with $\ge -$ 1.0 is shown in Fig.
15. where the absolute magnitudes of the field variables are determined by using Eqs. (
15, where the absolute magnitudes of the field variables are determined by using Eqs. (
8) and (9).
8) and (9).
It is quite clear that the linear M/y(RR)-[Fe/H] relation provided by Galactic globular clusters 1s not suitable to the most metal-rich ([Fe/H|x -0.7) field variables.
It is quite clear that the linear $M_V$ (RR)-[Fe/H] relation provided by Galactic globular clusters is not suitable to the most metal-rich $\ge -$ 0.7) field variables.
Conversely. we show in Fig.
Conversely, we show in Fig.
16 that over the whole metallicity range of [Fe/H]=-2.5 to ~ Oall the variables are nicely fitted by the quadratic relation We have shown that the value of the mixing length parameter influences the zero-point of the Pertod-Amplitude-Magnitude relation |Eqs. (
16 that over the whole metallicity range of $-$ 2.5 to $\sim$ 0 the variables are nicely fitted by the quadratic relation We have shown that the value of the mixing length parameter influences the zero-point of the Period-Amplitude-Magnitude relation [Eqs. (
5) and (6)] and consequently the M¥4\(RR)- [Fe/H] calibration (see Fig.
5) and (6)] and consequently the $M_V$ (RR)-[Fe/H] calibration (see Fig.
14).
14).
In order to constram the most appropriate value of the mixing-length parameter for globular cluster RR,» stars. we show in Fig.
In order to constrain the most appropriate value of the mixing-length parameter for globular cluster $_{ab}$ stars, we show in Fig.
17 the Vo magnitudes of RR,» stars in ω Cen (Piersimoni et al.
17 the $V_0$ magnitudes of $_{ab}$ stars in $\omega$ Cen (Piersimoni et al.
2007. in. preparation) versus the observed KCI.5),,; and Κ.ΟΊων parameters.
2007, in preparation) versus the observed $k(1.5)_{puls}$ and $k(2.0)_{puls}$ parameters.
By using Eqs.
By using Eqs.
from two distant gamma ray sources, and making the simple requirement that a passively evolving photon population not create an optical depth that is greater than 1.
from two distant gamma ray sources, and making the simple requirement that a passively evolving photon population not create an optical depth that is greater than 1.
In §3.2,, we will address the how these limits on the EBL can be translated into bounds on pop-III star formation.
In \ref{sec:highzsf}, we will address the how these limits on the EBL can be translated into bounds on pop-III star formation.
We find that a simple gamma-ray optical thinness criterion can significantly limit the high-redshift contribution to the observed EBL.
We find that a simple gamma-ray optical thinness criterion can significantly limit the high-redshift contribution to the observed EBL.
In Fig. 1,,
In Fig. \ref{fig:ebllims},
we show the limits placed on EBL contributions from high redshift, based on two different gamma-ray sources that have been recently observed with LAT.
we show the limits placed on EBL contributions from high redshift, based on two different gamma-ray sources that have been recently observed with LAT.
The first, GRB 080916C (Abdoetal.2009) was a bright gamma-ray burst with a measured redshift of z=4.35 (Greineretal. 2009), which was observed byFermi LAT and GBM on September 16, 2008.
The first, GRB 080916C \citep{abdo09a} was a bright gamma-ray burst with a measured redshift of z=4.35 \citep{greiner09}, which was observed by LAT and GBM on September 16, 2008.
The spectrum of the source at high energy was found to be continuous over nearly 6 orders of magnitude, and the highest energy event was a 13.6 GeV photon, corresponding to a rest-frame energy of 73 GeV. The other source is flat-spectrum radio quasar PKS 1502+106, which was detected during a flaring event at z=1.839 (Abdoetal.2010b).
The spectrum of the source at high energy was found to be continuous over nearly 6 orders of magnitude, and the highest energy event was a 13.6 GeV photon, corresponding to a rest-frame energy of 73 GeV. The other source is flat-spectrum radio quasar PKS 1502+106, which was detected during a flaring event at z=1.839 \citep{abdo10b}.
. The spectrum of this source was found to be most satisfactorily fit by a log-parabolic function peaking near 1 GeV. Abdoetal.(2010a) reported a highest energy photon of 48.9 GeV from this source, corresponding to a rest-frame energy of 139 GeV. The thick colored lines in the figure delineate the allowed present-day flux as a function of wavelength that does not produce an optical depth greater than 1 for the highest energy photons from each of these sources.
The spectrum of this source was found to be most satisfactorily fit by a log-parabolic function peaking near 1 GeV. \citet{fermiEBL} reported a highest energy photon of 48.9 GeV from this source, corresponding to a rest-frame energy of 139 GeV. The thick colored lines in the figure delineate the allowed present-day flux as a function of wavelength that does not produce an optical depth greater than 1 for the highest energy photons from each of these sources.
These contours were created by considering a large number of hypothetical r-EBL SEDs, and excluding those for which the optical depth exceeded unity.
These contours were created by considering a large number of hypothetical r-EBL SEDs, and excluding those for which the optical depth exceeded unity.
These results are sensitive to the nature of the SED types considered - for this exercise we have used blackbody spectra at a variety of temperatures and normalizations.
These results are sensitive to the nature of the SED types considered - for this exercise we have used blackbody spectra at a variety of temperatures and normalizations.
Because the optical depth from pair-production interactions (Eq. 10))
Because the optical depth from pair-production interactions (Eq. \ref{eq:opdep}) )
depends on the integrated spectrum falling within the allowed energy range, the actual bound is dependent on the assumed type of spectral feature; a more sharply-peaked SED would be less constraining.
depends on the integrated spectrum falling within the allowed energy range, the actual bound is dependent on the assumed type of spectral feature; a more sharply-peaked SED would be less constraining.
Our use of simple blackbody spectra avoids the introduction of any parameters into the results at this point.
Our use of simple blackbody spectra avoids the introduction of any parameters into the results at this point.
A couple of sample blackbody spectra are shown on this plot, each at a normalization that would produce optical depth unity for the highest energy photon detected from GRB 080916C. This plot also shows the r-EBL SED models described in 82.3 that we employ in the next section, assuming a star formation rate density (SFRD) of 0.2 yr! Mpc? from z=15 to z= 6.
A couple of sample blackbody spectra are shown on this plot, each at a normalization that would produce optical depth unity for the highest energy photon detected from GRB 080916C. This plot also shows the r-EBL SED models described in \ref{sec:spectmod} that we employ in the next section, assuming a star formation rate density (SFRD) of 0.2 $^{-1}$ $^{-3}$ from $z=15$ to $z=6$ .
We have also plotted the EBL prediction from Fernandezetal.(2010),, when using a star-formation efficiency of f.= 0.01.
We have also plotted the EBL prediction from \citet{fernandez10}, when using a star-formation efficiency of $f_* = 0.01$ .
Differences in the shape of the prominent peak created by emission are due to the fact that star-formation is not constantin their model, and decreases with redshift.
Differences in the shape of the prominent peak created by emission are due to the fact that star-formation is not constantin their model, and decreases with redshift.
Our results heredemonstrate the power of high
Our results heredemonstrate the power of high
‘This is the first occasion that evcle-to-evcle changes in he pulsation period ofa Cepheid can be followed.
This is the first occasion that cycle-to-cycle changes in the pulsation period of a Cepheid can be followed.
The OC analysis ofKepler data of V11I54CCve is published in a separate paper (Derckas ct al.
The $O-C$ analysis of data of Cyg is published in a separate paper (Derekas et al.
2011. in prep.).
2011, in prep.).
Here we study he long-term behaviour of the pulsation period.
Here we study the long-term behaviour of the pulsation period.
The computed. times of maxima were calculated. [rom he period fitted to theNepler data.
The computed times of maxima were calculated from the period fitted to the data.
One ο6 point was derived. for the mic-epoch of annual sections for each available photometric time series taken from the iterature.
One $O-C$ point was derived for the mid-epoch of annual sections for each available photometric time series taken from the literature.
Besides theAvepler maxima and available CCD ancl photoelectrie observation we publish for the first time VISACCve cla from digitized Harvard. plates and eve estimations from Sternberg Astronomical Institute. (SAL) photographic plates.
Besides the maxima and available CCD and photoelectric observation we publish for the first time Cyg data from digitized Harvard plates and eye estimations from Sternberg Astronomical Institute (SAI) photographic plates.
Phe passband of these observations is close to Johnson D.
The passband of these observations is close to Johnson $B$.