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2000). combined: with the recent infall of a companion galaxy. is very attractive: it provides a natural explanation for the low metallicity observed in the star-forming knots near the centre. while it also explains why we observe concentrated star formation in star clusters or very luminous regions in a tightlv-wound rine-like structure surrounding the centre (e.g. van der Ixruit de να 1976. TOs4. P93. Meurer et al. | 2000), combined with the recent infall of a companion galaxy, is very attractive: it provides a natural explanation for the low metallicity observed in the star-forming knots near the centre, while it also explains why we observe concentrated star formation in star clusters or very luminous regions in a tightly-wound ring-like structure surrounding the centre (e.g., van der Kruit de Bruyn 1976, TG84, P93, Meurer et al. |
1995. 896. Conselice ct al. | 1995, S96, Conselice et al. |
2000. 1:02). coinciding with the end of the nuclear bar (C'onselice et al. | 2000, E02), coinciding with the end of the nuclear bar (Conselice et al. |
2000. but see DOO). but not inside this ring. | 2000, but see D00), but not inside this ring. |
The most luminous single star-forming region in NGC 3310 is the giant. “Jumbo” region (191151. TCG8S4. 896): it is roughly 10. more luminous than 30. Doraclus in. La (BLISI). but it is of very low metallicitv. Z0.1Z. (P93). | The most luminous single star-forming region in NGC 3310 is the giant “Jumbo” region (BH81, TG84, S96); it is roughly $10 \times$ more luminous than 30 Doradus in $\alpha$ (BH81), but it is of very low metallicity, $Z \sim 0.1 Z_\odot$ (P93). |
The Jumbo region contains several individual UV-hbright star clusters (Aleurer et al. | The Jumbo region contains several individual UV-bright star clusters (Meurer et al. |
1995. E02). | 1995, E02). |
The most luminous star-forming structure is the ring of regions. which produces some ofthe observed far-UV Dux of NGC 3310 (896). | The most luminous star-forming structure is the ring of regions, which produces some of the observed far-UV flux of NGC 3310 (S96). |
We have indicated. our cluster detections in Fig. 2.. | We have indicated our cluster detections in Fig. \ref{clustercoords.fig}, |
overlaid on our press release image of the galaxy (Winclhorst et al. | overlaid on our press release image of the galaxy (Windhorst et al. |
2001). | 2001). |
Numerous other. less [luminous star clusters and regions. with sizes [rom 10 pe for the most luminous ones to S| pe for the unresolved clusters in the background cise (see Conselice et al. | Numerous other, less luminous star clusters and regions, with sizes from 10 pc for the most luminous ones to $\lesssim 1$ pc for the unresolved clusters in the background disc (see Conselice et al. |
2000). are found scattered in the galaxys outer parts bevond Z7?z4 kpe (van der Ixruit de Bruvn 1976. BLISL. δα. MvDD95. S96) and as condensations in the “how and arrow” structure (e.g. Bertola Sharp 1984. SOG. INSOL): these clusters may have been produced by the aceretion event. or might be remnants of the progenitor companion galaxy (cl. | 2000), are found scattered in the galaxy's outer parts beyond $R \simeq 4$ kpc (van der Kruit de Bruyn 1976, BH81, TG84, MvDB95, S96) and as condensations in the “bow and arrow” structure (e.g., Bertola Sharp 1984, S96, KS01); these clusters may have been produced by the accretion event, or might be remnants of the progenitor companion galaxy (cf. |
BLISL. 896). | BH81, S96). |
All of the observational evidence points at very recent star formation in the star clusters and regions. and a time since the interaction of Ss10—107 vr (van der Ixruit. 1976. BUST. PCsdt. P93. SOG. E202): the optical colours ancl the equivalent widths of the Lla-bright) cireumnuclear sources are best reproduced. by a combination of a 2.5 Myr and an S Alvr-olcl population (cf. | All of the observational evidence points at very recent star formation in the star clusters and regions, and a time since the interaction of $\lesssim 10^7-10^8$ yr (van der Kruit 1976, BH81, TG84, P93, S96, E02): the optical colours and the equivalent widths of the $\alpha$ -bright circumnuclear sources are best reproduced by a combination of a 2.5 Myr and an 8 Myr-old population (cf. |
P93. DOO). | P93, D00). |
‘This is consistent with(7) the detection of WR. features in the spectra of a few of the regions. including the Jumbo region. ancl of the NIR Ca LL triplet (Verlevich ct al. | This is consistent with the detection of WR features in the spectra of a few of the regions, including the Jumbo region, and of the NIR Ca II triplet (Terlevich et al. |
1990. P93). both | 1990, P93), both |
ol the relevant leptonic chanuels discussed in Section 5 is displayed as a fuuction of Higgs mass Lor our Chosen benchmark poiut in the L2HDM. | of the relevant leptonic channels discussed in Section \ref{sec:LHCSignatures} is displayed as a function of Higgs mass for our chosen benchmark point in the L2HDM. |
The SM results for the same processes are shown iu the left-pauel for comparison. | The SM results for the same processes are shown in the left-panel for comparison. |
The results in each panel correspond to an integrated luminosity of £=30fbI | The results in each panel correspond to an integrated luminosity of $\mathcal{L}=30~\mbox{~fb}^{-1}$. |
It is apparent [rom Fig. | It is apparent from Fig. |
9. that gq’—qq'h(h7)isone of the most promising detection channels for the chosen benchimark point in the L2HDM. as in the SM. | \ref{fig:SigPlot} that $qq'\rightarrow qq'h(h\rightarrow\tau^+\tau^-)$ is one of the most promising detection channels for the chosen benchmark point in the L2HDM, as in the SM. |
For this particular choice ol parameters. and Vio(ft) does not deviate drastically [rom D7;7SAL(1) (see Fig. 8)). | For this particular choice of parameters, and $\Gamma_{\mathit{tot}}(h)$ does not deviate drastically from $\Gamma_{\mathit{tot}}^{SM}(h)$ (see Fig. \ref{fig:GamTotPlot}) ), |
and consequently the overall siguilicauce level iu this chaunel is essentially uuchaueed from its SM value. | and consequently the overall significance level in this channel is essentially unchanged from its SM value. |
However. iu other regions of parameter space. drastic amplificatious cau occur: for example. the choice 0.3. tan;= 7) results in a amplification of the statistical siguificance for the same process by a factor of f. | However, in other regions of parameter space, drastic amplifications can occur: for example, the choice $\sin\alpha = 0.3$ , $\tan\beta=7$ ) results in a amplification of the statistical significance for the same process by a factor of $\sim 4$. |
It should also be noted that in the (sina=0.55.t tan=3) case. the significance levels [or both gg—fhr7 aud {ή—rr) also exceed So. | It should also be noted that in the $\sin\alpha = 0.55$, $\tan\beta=3$ ) case, the significance levels for both $gg\rightarrow h\rightarrow \tau\tau$ and $t\bar{t}h(h\rightarrow\tau\tau)$ also exceed $5\sigma$. |
The processes in which the Higgs decays to 1uuous are statistically less significant. but also provide strong evidence at the 3o. level with > 100 fb.+ of integrated luminosity. | The processes in which the Higgs decays to muons are statistically less significant, but also provide strong evidence at the $3\sigma$ level with $\gtrsim$ 100 ${\rm fb}^{-1}$ of integrated luminosity. |
ludeed. the evidence for such a Higgs bosou would be dramatic aud ummnistakable. | Indeed, the evidence for such a Higgs boson would be dramatic and unmistakable. |
Furthermore. once the Higgs is observed in any of the muonic chiauuels. the excellent invariaut-11nass resolution of the muon pairs can be used to determiue the value of 115, with a very high degree of precision. | Furthermore, once the Higgs is observed in any of the muonic channels, the excellent invariant-mass resolution of the muon pairs can be used to determine the value of $m_h$ with a very high degree of precision. |
While the significances in tliose chaunels which involve a leptonically-decayiug HiggsMD can poteutially be amplified in L2HDM. those iu other chanuels useful for the detection of a SM Higgs may be substantially suppressed. | While the significances in those channels which involve a leptonically-decaying Higgs can potentially be amplified in L2HDM, those in other channels useful for the detection of a SM Higgs may be substantially suppressed. |
This is illustrated in Fig. 10.. | This is illustrated in Fig. \ref{fig:SigPlotOther}, |
which shows the siguificance of discovery in each individual channel which contributes meauinelully to the discovery. potential of a SM Higgs bosou in the low to intermeciate-1nass region. both iu the SM Celt-hanucl pauel) and in the L2HDM | which shows the significance of discovery in each individual channel which contributes meaningfully to the discovery potential of a SM Higgs boson in the low to intermediate-mass region, both in the SM (left-hand panel) and in the L2HDM |
over (he boxcar limits. | over the boxcar limits. |
To convert our flux measurements from count rates to CGS units we extracted spectra whose limits were defined by (he start and stop time of the boxcars. | To convert our flux measurements from count rates to CGS units we extracted spectra whose limits were defined by the start and stop time of the boxcars. |
For each [αν measurement we extracted a spectrum for (he region of interest. a background spectrum and a response matrix. similar to what was done for (he spectral evolution analvsis. | For each flux measurement we extracted a spectrum for the region of interest, a background spectrum and a response matrix, similar to what was done for the spectral evolution analysis. |
Each spectrum was fit to a photolectrically absorbed blackbody using in order to measure the 220 keV flux in CGS units. | Each spectrum was fit to a photolectrically absorbed blackbody using in order to measure the 2–20 keV flux in CGS units. |
The 220 keV total (uence was determined by integrating our background-subtracted time series. | The 2–20 keV total fluence was determined by integrating our background-subtracted time series. |
If the burst had emitted all of its energy during the observation. the integrated burst profile would eventually plateau. | If the burst had emitted all of its energy during the observation, the integrated burst profile would eventually plateau. |
However. this is nol what we observed. | However, this is not what we observed. |
The integrated burst profile was still steadily rising even at the end ol our observation. indicating that our observation finished before catching the end of the burst. | The integrated burst profile was still steadily rising even at the end of our observation, indicating that our observation finished before catching the end of the burst. |
Thus. we can only set an upper-Iimit on the total 220 keV fluence: see Table 1.. | Thus, we can only set an upper-limit on the total 2–20 keV fluence; see Table \ref{ta:burst}. |
To convert our {lnence upper limit from counts to CGS units the exact procedure was followed as lor (he peak flux measurements. | To convert our fluence upper limit from counts to CGS units the exact procedure was followed as for the peak flux measurements. |
Magnetar candidates have been observed to be highly flux. variable. which is why we regularly monitor the pulsed {lis of this source (seeGavril&Ixaspi2004.foradetaileddiscussionofpulsedfluxcaleulationsForLE1048.1. 5937). | Magnetar candidates have been observed to be highly flux variable, which is why we regularly monitor the pulsed flux of this source \citep[see][for a detailed discussion of pulsed flux calculations for
\tfe]{gk04}. |
. The pulsed flux during the entire observation in which the burst occured was not significantly higher than in neighboring observations. | The pulsed flux during the entire observation in which the burst occured was not significantly higher than in neighboring observations. |
Llowever. in some ANPs and SGRs. short time-scale («&1000 s) abrupt changes in pulsed fIux have been observed in conjunction with bursts (e.g.Lentersetal.2003:elal.2004. 2005). | However, in some AXPs and SGRs, short time-scale $\ll
1000$ s) abrupt changes in pulsed flux have been observed in conjunction with bursts \citep[e.g.][]{lwg+03,wkt+04,wkg+05}. |
. Motivated by such observations we decided to search for short-term pulsed fIux enhancements around the time of the burst from5937. | Motivated by such observations we decided to search for short-term pulsed flux enhancements around the time of the burst from. |
. We broke the observation into LO intervals and caleulated the pulsed flux for each. | We broke the observation into 10 intervals and calculated the pulsed flux for each. |
In order to avoid having the burst spike biasing our pulsed [lux measurements we removed a 4 s interval centered on the burst peak. | In order to avoid having the burst spike biasing our pulsed flux measurements we removed a 4 s interval centered on the burst peak. |
A factor of 3.5 increase in pulsed flux can be seen in the tail of the burst (see Fig. 2)). | A factor of 3.5 increase in pulsed flux can be seen in the tail of the burst (see Fig. \ref{fig:flux}) ). |
This coupling between bursting activity and pulsed flux establishes that iis definitely the burst source. | This coupling between bursting activity and pulsed flux establishes that is definitely the burst source. |
Following the discovery of a new burst from the direction of5937.. we triggered observations of the source with imaging X-ray telescopes. | Following the discovery of a new burst from the direction of, we triggered observations of the source with imaging X-ray telescopes. |
The AXP was observed once with | The AXP was observed once with |
dominating their X-ray emission. | dominating their X-ray emission. |
In the jy calculation we consider sources in the redshift range 0.4<z«0.5 with SpacSOs (es the Se radio completeness limit) anc fx(0.5 SkeV)>110Moreslem2 (X-ray detection limit). | In the $j_X$ calculation we consider sources in the redshift range $0.4<z<0.5$ with $S_{1.4}>50\mu
\rm Jy$ (i.e. the $5\sigma$ radio completeness limit) and $f_X(\rm
0.5-8\,keV)>1\times10^{-16}\rm erg\,s^{-1}\,cm^{-2}$ (X-ray detection limit). |
Only 4 sources in the Dauer et al. ( | Only 4 sources in the Bauer et al. ( |
2002) sample satisfy these criteria and therefore the estimated jx. plottec in Figure 5. sullers from large uncertainties. | 2002) sample satisfy these criteria and therefore the estimated $j_X$ plotted in Figure \ref{fig_jx} suffers from large uncertainties. |
Despite the small number statistics the py of the CDE-N X-ray detecte radio sources at z20.5 is elevated compared to both loca spirals and z0.3 radio sources. | Despite the small number statistics the $j_X$ of the CDF-N X-ray detected radio sources at $z\approx0.5$ is elevated compared to both local spirals and $z\approx0.3$ radio sources. |
The significance leve of this result is low but it is interesting that to the firs approximation the jy of X-raydeleefed radio sources is in broad agreement with the stacking analysis results presente here. | The significance level of this result is low but it is interesting that to the first approximation the $j_X$ of X-ray radio sources is in broad agreement with the stacking analysis results presented here. |
Under the assumption that the faint radio population is indeed dominated by star-formation activity. the observe increase in ὃν al least to z0.3 SUgegests ab evolutionary rate of the form (1|z)* broadly consistent with the Peak-\ moclels of Ghosh White (2001). | Under the assumption that the faint radio population is indeed dominated by star-formation activity, the observed increase in $j_X$ at least to $z\approx0.3$ suggests an evolutionary rate of the form $(1+z)^{3}$ broadly consistent with the Peak-M models of Ghosh White (2001). |
This set of models is also in agreemen with the stacking analysis results of 2&0.5 spirals of Brand 6 al. ( | This set of models is also in agreement with the stacking analysis results of $z\approx0.5$ spirals of Brandt et al. ( |
2001b). | 2001b). |
Assuming that the faint radio population is dominate ov star-formation activity (as the present study. indicates) 1ο increase in jy with redshift is attributed to the evolution the global star-formation rate. | Assuming that the faint radio population is dominated by star-formation activity (as the present study indicates) the increase in $j_X$ with redshift is attributed to the evolution of the global star-formation rate. |
Using the empirica relation SPR.=δν107?Lx(0.5 2keV)M.ve+ (Ranalli οἱ al. | Using the empirical relation $\rm
SFR=2.2\times10^{-40}\,\times\,L_X(\rm 0.5-2\,keV)\,\,M_{\odot}\,yr^{-1}$ (Ranalli et al. |
2003) we estimate a SER density of 0.029+Ἑve+AMpe% for the 2=0.00.3 redshilt. bin at a median. redshift of 0.240 in good agreement with orevious studies selecting star-forming galaxies at UV. optical or racio wavelengths. | 2003) we estimate a SFR density of $0.029\pm0.007 \rm
M_{\odot}\,yr^{-1}\,Mpc^{-3}$ for the $z=0.0-0.3$ redshift bin at a median redshift of 0.240 in good agreement with previous studies selecting star-forming galaxies at UV, optical or radio wavelengths. |
Εις is demonstrated in Figure 6.. showing the SET. density as a function of redshift. | This is demonstrated in Figure \ref{fig_sfr}, showing the SFR density as a function of redshift. |
In this paper we apply stacking analvsis to study the mean X-rav properties of sub-mJdy radio galaxies using a 50kks pointing overlapping with a subregion of a deep and homogeneous radio survey reaching jd sensitivities (Phoenix Deep Survey). | In this paper we apply stacking analysis to study the mean X-ray properties of sub-mJy radio galaxies using a ks pointing overlapping with a subregion of a deep and homogeneous radio survey reaching $\mu$ Jy sensitivities (Phoenix Deep Survey). |
Multiwayelength UV. optical and. NIlt. photometric data are available for this field allowing photometric redshifts and spectral types (1.6. ellipticals. spirals) to be estimated: for all radio: sources brighter than #=21.5 mmag (total of 82). | Multiwavelength UV, optical and NIR photometric data are available for this field allowing photometric redshifts and spectral types (i.e. ellipticals, spirals) to be estimated for all radio sources brighter than $R=21.5$ mag (total of 82). |
The subsample of 2<21.5 mamag faint racio sources with spiral galaxy SEDs (total of 34. after excluding AGN dominated sources) is segregated into two redshift bins with a median of z=0.240 and 0.455 respectively. | The subsample of $R<21.5$ mag faint radio sources with spiral galaxy SEDs (total of 34, after excluding AGN dominated sources) is segregated into two redshift bins with a median of $z=0.240$ and 0.455 respectively. |
Stacking analvsis is used to study the mean X-ray properties of the sources in each redshift’ subsample. | Stacking analysis is used to study the mean X-ray properties of the sources in each redshift subsample. |
In the kkeV soft band a statistically significant signal is obtained for both the low and the high-z subsamples. | In the keV soft band a statistically significant signal is obtained for both the low and the $z$ subsamples. |
Stacking of the hard band counts vields a mareinally significant signal (2.60) for the z=0.455 subsample only. | Stacking of the hard band counts yields a marginally significant signal $2.6\sigma$ ) for the $z=0.455$ subsample only. |
We argue that the AGN contamination of our sample is minimal on the basis of the following arguments: Although the evidence above suggests that the observed X-ray emission ids dominated bv star-Lormation activity we cannot exclude. the possibility of contamination of the present sample by a small number of low-luminosit. or heavily obseurech AGN. | We argue that the AGN contamination of our sample is minimal on the basis of the following arguments: Although the evidence above suggests that the observed X-ray emission is dominated by star-formation activity we cannot exclude the possibility of contamination of the present sample by a small number of low-luminosity or heavily obscured AGN. |
Our main conclusions are summarised below: | Our main conclusions are summarised below: |
EIS scanned point 1 at 02:33:58 UT in the impulsive phase of the flare (see Figure 3)). | EIS scanned point 1 at 02:33:58 UT in the impulsive phase of the flare (see Figure \ref{flux_curve}) ). |
This point shows upllows in all emission lines except for the Πο II line. | This point shows upflows in all emission lines except for the He II line. |
Because of line blending we do not use the Ca XVII and Fe XXIV lines. | Because of line blending we do not use the Ca XVII and Fe XXIV lines. |
We also ignore the Fe XXIII line since it is very weak in this region. | We also ignore the Fe XXIII line since it is very weak in this region. |
We plot the line profiles and their filling curves in Figure 4.. | We plot the line profiles and their fitting curves in Figure \ref{profilep1}. |
Note that in the wavelength window of the Fe X 184.54 line. there exists the Fe XI 184.41 line: however. the latter is «quite distinguishable from. the former and therefore does not affect the fitting result (see also Figures 5. and 6)). | Note that in the wavelength window of the Fe X 184.54 line, there exists the Fe XI 184.41 line; however, the latter is quite distinguishable from the former and therefore does not affect the fitting result (see also Figures \ref{profilep2} and \ref{profilep3}) ). |
Some of the line profiles are fitted by double Gaussian components. | Some of the line profiles are fitted by double Gaussian components. |
We also caleulate the intensity ratio of the blue component to the stationary one. | We also calculate the intensity ratio of the blue component to the stationary one. |
The ratio and the Doppler velocity for the blue component are listed in Table 2.. | The ratio and the Doppler velocity for the blue component are listed in Table \ref{velp1}. |
From the results we can see that most of the lines show obvious blueshifts. | From the results we can see that most of the lines show obvious blueshifts. |
The upflow velocity increases with temperature [rom several tens of kms ! to the highest one of 116 kms τν the intensity ratio of the (wo components. derived [rom the double Gaussian fitting. has also an increasing tendency. | The upflow velocity increases with temperature from several tens of km $^{-1}$ to the highest one of 116 km $^{-1}$; the intensity ratio of the two components, derived from the double Gaussian fitting, has also an increasing tendency. |
For the Fe XIII line. the intensity of the blue component is just 1.72 times that of the stationary component. | For the Fe XIII line, the intensity of the blue component is just 1.72 times that of the stationary component. |
However. the ralio increases lo 9.95 [or the Fe NWI line. | However, the ratio increases to 9.95 for the Fe XVI line. |
This indicates that the blue components are dominant over (he stationary components for most of Che lines. especially (he lines with hieher formation temperatures. | This indicates that the blue components are dominant over the stationary components for most of the lines, especially the lines with higher formation temperatures. |
We also measure the average Doppler velocity over the 9 EIS pixels around point 1in the preflare aud post-impulsive phases and plot the value against the temperature in Figure 7.. | We also measure the average Doppler velocity over the 9 EIS pixels around point 1 in the preflare and post-impulsive phases and plot the value against the temperature in Figure \ref{velp1p2}. |
The same is done for point 2. | The same is done for point 2. |
Note that the velocities corresponding to the blue components from the double Gaussian fitting are marked with the plus svimbols. | Note that the velocities corresponding to the blue components from the double Gaussian fitting are marked with the plus symbols. |
We find that most of (he lines show significant blueshifts or blue-shifted components in the impulsive phase: while the shifts are trivial in (he preflare and post-inipulsive phases. | We find that most of the lines show significant blueshifts or blue-shifted components in the impulsive phase; while the shifts are trivial in the preflare and post-impulsive phases. |
In contrast to the case of point 1. point 2 shows obvious cownflows. | In contrast to the case of point 1, point 2 shows obvious downflows. |
EIS scanned point 2 ab 02:38:14 UT. slightly before the GOES soft X-ray peak time at 02:42 UT. | EIS scanned point 2 at 02:38:14 UT, slightly before the GOES soft X-ray peak time at 02:42 UT. |
We fit all the line profiles at point 2 using a single Gaussian function (shown in Figure 5)). | We fit all the line profiles at point 2 using a single Gaussian function (shown in Figure \ref{profilep2}) ). |
Different from point 1. point 2 shows strong emission in high temperature lines (e.g.. the unblended Fe XXIII line). | Different from point 1, point 2 shows strong emission in high temperature lines (e.g., the unblended Fe XXIII line). |
Youngetal.(2007b) reported that the Ca ANVIL line completely dominates the other lines in large flares. | \cite{youn07b} reported that the Ca XVII line completely dominates the other lines in large flares. |
Therefore. we can ignore the blendings of the Ca XVII line with ihe O V line and the Fe XI line. | Therefore, we can ignore the blendings of the Ca XVII line with the O V line and the Fe XI line. |
The blending of the Fe XXIV line with the Fe XI line can | The blending of the Fe XXIV line with the Fe XI line can |
and MOG. the disagreement of the TER slope of the Sa spirals should be interpreted in caution because the incompleteness of the Sa spirals is huge. especially iu he low huuimositv cud. which will artificially bias the slope to lower value. | and M06, the disagreement of the TFR slope of the Sa spirals should be interpreted in caution because the incompleteness of the Sa spirals is large, especially in the low luminosity end, which will artificially bias the slope to lower value. |
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