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All seismology targets of LRaOI are subject to a small downward trend of instrumental origin. | All seismology targets of LRa01 are subject to a small downward trend of instrumental origin. |
This detrended light curve is shown as the thick line in the upper panel of reflc.. | This detrended light curve is shown as the thick line in the upper panel of \\ref{lc}. |
Large variations occur. with peak-to-peak values of | Large variations occur, with peak-to-peak values of |
lice was observed toward stars behind the Serpeus (Eiroa Illodapp 1989) and Taurus dark clouds (Whittet et 119885. Sinith ct 11993. Murakwa et 22000). | ice was observed toward stars behind the Serpens (Eiroa Hodapp 1989) and Taurus dark clouds (Whittet et 1988, Smith et 1993, Murakawa et 2000). |
These studies indicate that lice is formed deep iu the clouds. at visual extiuctions Ay>3 magnitudes (the ice formation threshold). | These studies indicate that ice is formed deep in the clouds, at visual extinctions $A_{\rm V}>3$ magnitudes (the ice formation threshold). |
Solid CO was observed toward backeround stars as well (e.9.. Whittet et 11985. Clüar et 11991. 1995). | Solid CO was observed toward background stars as well (e.g., Whittet et 1985, Chiar et 1994, 1995). |
Its formation threshold is significautle larger =66-15 mae) due to the lower stblinmation temperature of solid CO. | Its formation threshold is significantly larger 6-15 mag) due to the lower sublimation temperature of solid CO. |
Studies of ices towiux background stars lave been liuüted to biuds below 5 bbecause of telluric alsorption and the fact that stellar fluxes drop rapid vowith increasing waveleneth. | Studies of ices toward background stars have been limited to bands below 5 because of telluric absorption and the fact that stellar fluxes drop rapidly with increasing wavelength. |
Observations of back:oround stars lave become increasinely feasible— with the(ISO) mission aud. in particular. with the launch of theTelescope (Werner ct 22001). | Observations of background stars have become increasingly feasible with the mission and, in particular, with the launch of the (Werner et 2004). |
With7$0.. solid aat 1.25 wwas observed toward two Taurus backeround stars (Whittet et 11998: Ντο et 2001). indicating that radiation. from nearby protostarsαι is not required to form this species. | With, solid at 4.25 was observed toward two Taurus background stars (Whittet et 1998; Nummelin et 2001), indicating that radiation from nearby protostars is not required to form this species. |
Recent observations withSpitzer detected the Dbbending mode at 15 ttoward background stars (Bergin et 22005). | Recent observations with detected the bending mode at 15 toward background stars (Bergin et 2005). |
ere we present observations of ices toward three background | Here we present observations of ices toward three background |
Redshift surveys of clusters extending to large radius reveal that the (rumpet-shapecl patters are ubiquitous (Hines et al. | Redshift surveys of clusters extending to large radius reveal that the trumpet-shaped patters are ubiquitous (Rines et al. |
2003: Rines Dialerio 2006). | 2003; Rines Diaferio 2006). |
The majority of massive clusters ave well-separated from foreground. ancl backeround structures in redshift space. | The majority of massive clusters are well-separated from foreground and background structures in redshift space. |
Rines Diaferio (2006) show that masses computed from the caustic technique within the virial radius correspond well wilh those computed [rom the x-ray aud from other dynamical techniques. | Rines Diaferio (2006) show that masses computed from the caustic technique within the virial radius correspond well with those computed from the x-ray and from other dynamical techniques. |
Diaferio. Geller Rines (2005) show that mass profiles derived Irom weak lensing agree well with those derived from the caustic technique. | Diaferio, Geller Rines (2005) show that mass profiles derived from weak lensing agree well with those derived from the caustic technique. |
One of the most interesting results of the application of the caustic technique is the evaluation of the mass contained within the infall region. | One of the most interesting results of the application of the caustic technique is the evaluation of the mass contained within the infall region. |
Rines Dialerio (2006) show that the mass within Ry. is 2.90.18 times the mass inside Iso). in excellent agreement with theoretical predictions of the ultimate masses of clusters (Busha οἱ al. | Rines Diaferio (2006) show that the mass within $_{turn}$ is $\pm$ 0.18 times the mass inside $_{200}$, in excellent agreement with theoretical predictions of the ultimate masses of clusters (Busha et al. |
2005). | 2005). |
We can explore the infall regions of clusters efficiently now because of the power of wide-lielcl spectrographs like Hectospec. | We can explore the infall regions of clusters efficiently now because of the power of wide-field spectrographs like Hectospec. |
But. when Hectospec was just "an idea”. the main scientific motivation was to map the large-scale structure of the universe to greater reclshilt (Geller 1994). | But, when Hectospec was just “an idea”, the main scientific motivation was to map the large-scale structure of the universe to greater redshift (Geller 1994). |
Clusters of galaxies are markers of the highest density regions of the universe. but they are a poor second (o seeing (he entire grand pattern of the "cosmic web.” | Clusters of galaxies are markers of the highest density regions of the universe, but they are a poor second to seeing the entire grand pattern of the “cosmic web.” |
The first large-area survey with Ilectospec is now underway. | The first large-area survey with Hectospec is now underway. |
We call it HectoMAD. | We call it HectoMAP. |
In 1939. the Great Wall was the largest structure known in the universe (Geller Tluchra 1989). | In 1989, the Great Wall was the largest structure known in the universe (Geller Huchra 1989). |
It still seems remarkable that the largest structure was as big as it could be to fit within the survey boundaries. | It still seems remarkable that the largest structure was as big as it could be to fit within the survey boundaries. |
The Sloan Digital Skv Survey contains the Sloan Great Wall. | The Sloan Digital Sky Survey contains the Sloan Great Wall. |
Estimates suggest that the Sloan Great. Wall is only ereater in extent than the CLA Creat Wall (Gott et al. | Estimates suggest that the Sloan Great Wall is only greater in extent than the CfA Great Wall (Gott et al. |
2005) even though the Sloan redshift survey is more than three limes as deep as the CLA slices. | 2005) even though the Sloan redshift survey is more than three times as deep as the CfA slices. |
The Sloan Great Wall is. nonetheless. a potential challenge to our understanding of the development of the cosmic web from Gaussian initial conditions (Sheth Dialerio 2011). | The Sloan Great Wall is, nonetheless, a potential challenge to our understanding of the development of the cosmic web from Gaussian initial conditions (Sheth Diaferio 2011). |
The obvious question is whether (he Sloan uncovered (he biggest structure: after all. ihe Sloan Great Wall could be bigger and still [it in the survey. | The obvious question is whether the Sloan uncovered the biggest structure; after all, the Sloan Great Wall could be bigger and still fit in the survey. |
One of the goals of our new survey. HectoMAD. is to begin to approach this question by carrving out a deep dense redshift survev over an area large enough to detect. "greater walls.” | One of the goals of our new survey, HectoMAP, is to begin to approach this question by carrying out a deep dense redshift survey over an area large enough to detect “greater walls.” |
IlectoMAP is a redshift survey of red galaxies (g—r>1 and r—i> 0.5) with SDSS Yoogpo<21 and vp,<22 covering a 50 square degree region of the skv with and 42.5°<doggy44. | HectoMAP is a redshift survey of red galaxies $g - r > 1$ and $r - i > 0.5$ ) with SDSS $_{petro} < 21$ and $_{fiber} < 22$ covering a 50 square degree region of the sky with $200^\circ < \alpha_{2000} < 250^\circ$ and $42.5^\circ < \delta_{2000} < 44^\circ$. |
We select galaxies from the SDSS. | We select galaxies from the SDSS. |
The complete survey. will | The complete survey will |
We now study how the colour-o relation (hereafter Colt) depends on the aperture within which the colour is defined. | We now study how the $\sigma$ relation (hereafter $\sigma$ R) depends on the aperture within which the colour is defined. |
Figure IS shows the Colt for our sample for rest-[rame 6r from the spectra. before (top) anc after (bottom) correcting for signal/noise. | Figure \ref{cvr-spec} shows the $\sigma$ R for our sample for rest-frame $g-r$ from the spectra, before (top) and after (bottom) correcting for signal/noise. |
Fitting the galaxies bv equation (S). but with AZ, replaced by loge vields with rms residuals of 0.0521 magnitudes. | Fitting the galaxies by equation (8), but with $M_r$ replaced by $\log\sigma$ yields with rms residuals of 0.0521 magnitudes. |
Note that colour evolution is strongly. detected. but it is ~0.06 0.072 less than in the corresponding CALR. | Note that colour evolution is strongly detected, but it is $\sim 0.06$ $0.07z$ less than in the corresponding CMR. |
Using gor and the (uncorrected or corrected): spectra-derived k-corrections vields with rms residuals of 0.0551. magnitudes (see Figure 19)). | Using $g-r$ and the (uncorrected or corrected) spectra-derived k-corrections yields with rms residuals of 0.0551 magnitudes (see Figure \ref{cvr-model}) ). |
This Colt relation is bluer than that based on the spectra. but the slope is almost unchanged. | This $\sigma$ R relation is bluer than that based on the spectra, but the slope is almost unchanged. |
Llowever. it has 0.19: less evolution. | However, it has $0.13z$ less evolution. |
Finally. usingfiber gr (Figure 20)) instead. gives Following Bernardi et al. ( | Finally, using $g-r$ (Figure \ref{cvr-fiber}) ) instead gives Following Bernardi et al. ( |
2005). if we plot the residual o[g rcolours magnitude and signal/noise corrected) tothe CM (9.rg(ggriep|Adp.co}. against the residual of a to the best-fit aAJ, relation. we see the two are correlated (Figure 21)). | 2005), if we plot the residual of $g-r$ colours magnitude and signal/noise corrected) to the CMR, $(g-r)_{rf}-\langle (g-r)_{rf}|M_r,
z\rangle$, against the residual of $\sigma$ to the best-fit $\sigma-M_r$ relation, we see the two are correlated (Figure \ref{resid1}) ). |
Phe correlation coellicient is 0.2094 with a slope 0.2033+ 0.0024. an intercept 0.0000c0.0002. and a standard deviation 0.0559 mag. | The correlation coefficient is 0.2994 with a slope $0.2033\pm 0.0024$ , an intercept $0.0000\pm 0.0002$, and a standard deviation 0.0559 mag. |
However. if we plot the residual of colours to the Colt. (οεν(dr)r]a.z against the residual of AZ, to the best-lt AZ,@ relation (Fieure 22))). we find a much lower correlation coelIicient 1.0231 and slope 0.0022+0.0004 (vith an intercept again )0000£0.0002 and stancard deviation 0.0560 magnitudes). | However, if we plot the residual of colours to the $\sigma$ R, $(g-r)_{rf}-\langle (g-r)_{rf}|\sigma,z \rangle$ against the residual of $M_r$ to the best-fit $M_r-\sigma$ relation (Figure \ref{resid2}) )), we find a much lower correlation coefficient 0.0231 and slope $0.0022\pm 0.0004$ (with an intercept again $0.0000\pm 0.0002$ and standard deviation 0.0560 magnitudes). |
Again we [ind that gr colour in L/SOs are more tundamentally dependent. on. velocity. dispersion than on Al. and furthermore. the absence of an AL, correlation in 10 second case implies that the Colt is the same for all jns in magnitude at a fixed velocity dispersion. | Again we find that $g-r$ colour in E/S0s are more fundamentally dependent on velocity dispersion than on $M_r$, and furthermore, the absence of an $M_r$ correlation in the second case implies that the $\sigma$ R is the same for all bins in magnitude at a fixed velocity dispersion. |
Hence our ctermination of the Colt in this Section should. not he gaignificantly biased by our use of a Iux-limited sample. | Hence our determination of the $\sigma$ R in this Section should not be significantly biased by our use of a flux-limited sample. |
1n both the CMIU and Colt. rest-frame colours from je spectra show stronger gr evolution 0.272) than model-magnitude colours with the "corrected? k-correction 0.132). | In both the CMR and $\sigma$ R, rest-frame colours from the spectra show stronger $g-r$ evolution $0.27z$ ) than model-magnitude colours with the `corrected' k-correction$0.13z$ ). |
This is an aperture alfect. due to the colour eracdients which also shift the zero-point of both the CALR | This is an aperture affect, due to the colour gradients which also shift the zero-point of both the CMR |
suggests that feedback aux backgrouud effects may not be independent. aud may iu fact amplify oue another.) | suggests that feedback and background effects may not be independent, and may in fact amplify one another.) |
Star formation is performace at a fixed density threshold & 7, or Hyper=0.03 0m.5m7. asΗ | Star formation is performed at a fixed density threshold g $^{-3}$, or $n_{H,\text{crit}}=0.03$ $^{-3}$, as. |
Ν For easv comparison. we use the same set of mitia conditions that were designated galaxies/halos A. C. and E in(2007).. aud are so designate here as well. | For easy comparison, we use the same set of initial conditions that were designated galaxies/halos A, C, and E in, and are so designated here as well. |
All sinulatious were performed with Loa? SPI particles (correspoucding to a eravitational softening leugth of 0.25 kpc for the gas aud star particles. aud twice that for the dark matter particles: eas aud star particles have masses of the order 10°AL..). | All simulations were performed with $100^3$ SPH particles (corresponding to a gravitational softening length of 0.25 kpc for the gas and star particles, and twice that for the dark matter particles; gas and star particles have masses of the order $10^6 M_{\odot}$ ). |
As will be discussed below. galaxies A aud C eave the same qualitative results. while galaxy E was somewhat cdiffercut. due to its different merecr aud accretion historv. | As will be discussed below, galaxies A and C gave the same qualitative results, while galaxy E was somewhat different, due to its different merger and accretion history. |
Since Galaxy A was the most well-studied in(2007)... we choose here to focus on it. bringing in the other two ealaxies where relevant. ( | Since Galaxy A was the most well-studied in, we choose here to focus on it, bringing in the other two galaxies where relevant. ( |
FO UV. a later addition to the study. was run on galaxy A only.) | FG UV, a later addition to the study, was run on galaxy A only.) |
Throughout the paper. all distances are physical except where noted: the assumed cosmology is (Qa,ΟνὉΌλιoshi)=(0.3.0.7.0.2.0.86.0.65) as in(2007). | Throughout the paper, all distances are physical except where noted; the assumed cosmology is $(\Omega_M,\Omega_{\Lambda},\Omega_b/\Omega_M,\sigma_8,h)=(0.3,0.7,0.2,0.86,0.65)$ as in. |
. We naturally expect au inerease in the ionizing backeround to lead to more efficient gas heating at lüeh redshifts. but the consequences on galaxy formation cannot easily be predicted. | We naturally expect an increase in the ionizing background to lead to more efficient gas heating at high redshifts, but the consequences on galaxy formation cannot easily be predicted. |
Figure 2 shows the temperature distribution of the eas in the central 2 Mpe of the galaxy. A simulation (essentially the ligh resolution regiou of validity: the siuulatiouns. beige resampled from a larecr box. are nonperiodic) with the three backgrounds. at 2=5.2 (chosen to be where the Old UW backeround is nonzero but still substantially lower than New UV / FG UV) and ;—0. | Figure \ref{fig:Thist-all} shows the temperature distribution of the gas in the central 2 Mpc of the galaxy A simulation (essentially the high resolution region of validity; the simulations, being resampled from a larger box, are nonperiodic) with the three backgrounds, at $z=5.2$ (chosen to be where the Old UV background is nonzero but still substantially lower than New UV / FG UV) and $z=0$. |
At 25.2 the models are well separated: iu fact the eas in the Old UV simulation has just finished heating from a very cold (LOOKS) state (to which it had cooled by adiabatic expansion froin its initial state at 2=2L the cooling function iupleumienuted iu the code is primordial (I-TWe ouly) aud. cuts off at 104 K). and has just finished IL/IIeI reiouization. | At $z=5.2$ the models are well separated; in fact the gas in the Old UV simulation has just finished heating from a very cold (100K) state (to which it had cooled by adiabatic expansion from its initial state at $z=24$; the cooling function implemented in the code is primordial (H-He only) and cuts off at $10^4$ K), and has just finished H/HeI reionization. |
As expected. the New UV model. which has been eutirelv. reionized Gucliding Well) since 28. has siguificantly more hot eas, and including X-ray heating pushes most of the eas up to temperatures siguificautly above 10! The FC UV anode! lies iu between Old UV and New K.UV. which is expected since it has a higher iuteusity than the former but a softer spectra than the latter. | As expected, the New UV model, which has been entirely reionized (including HeII) since $z\simeq 8$, has significantly more hot gas, and including X-ray heating pushes most of the gas up to temperatures significantly above $10^4$ K. The FG UV model lies in between Old UV and New UV, which is expected since it has a higher intensity than the former but a softer spectrum than the latter. |
FO UV has been ἩΠΟ reiouized since 28 but has not vet reionized ΠΟΠ. | FG UV has been H/HeI reionized since $z\simeq 8$ but has not yet reionized HeII. |
At ;—0. however. the Old UV. New UV. and FO UV models are csscutially identical: iucreasiug the iteusity or changing the spectrum of the ionizing backeround at high redshift has no effect at the preseut (in other words. the cooling time for all gas is shorter than 11 Car. siuce it has forgotten the extra heat from +> 2.1). | At $z=0$, however, the Old UV, New UV, and FG UV models are essentially identical; increasing the intensity or changing the spectrum of the ionizing background at high redshift has no effect at the present (in other words, the cooling time for all gas is shorter than 11 Gyr, since it has forgotten the extra heat from $z>2.4$ ). |
This is consistent with(2006)... who found that the effects of a UW backerouud which is suddenly turned off begim to dissipate after 0.3 of a IIubble time: since we are essentially turning off the vadiation at 5;=2.L we would expect a similar convergence. | This is consistent with, who found that the effects of a UV background which is suddenly turned off begin to dissipate after $\sim0.3$ of a Hubble time; since we are essentially turning off the radiation at $z=2.4$, we would expect a similar convergence. |
Moreover. found that well (22 Gyr) after reiouization. the IGAL equilibrium temperature approaches a value that depends only ou the spectral shape aud not the intensity of the ionizing background. so even if Ola UV and New UV had differeut intensitics all the wav to τ=0 we wouldut expect a sjenificaut difference. | Moreover, found that well $\gtrsim 2$ Gyr) after reionization, the IGM equilibrium temperature approaches a value that depends only on the spectral shape and not the intensity of the ionizing background, so even if Old UV and New UV had different intensities all the way to $z=0$ we wouldn't expect a significant difference. |
Adding N-avs (to all epochs). on the other laud. sienificautly heats the eas. especially the coldest gas: the neu gas temperature rises from 1.3«10° K to 1.5«10° I. The addition of N-ravs also produces a arecr total mass of eas due to reduced star formation in sinall svsteius. as we will see in the next section. | Adding X-rays (to all epochs), on the other hand, significantly heats the gas, especially the coldest gas: the mean gas temperature rises from $1.3\times 10^5$ K to $1.5\times 10^5$ K. The addition of X-rays also produces a larger total mass of gas due to reduced star formation in small systems, as we will see in the next section. |
This arger reservoir of cool (L0t10? IK) eas (cf. | This larger reservoir of cool $10^4 - 10^5$ K) gas (cf. |
the lower xuiel of Fie. 2)) | the lower panel of Fig. \ref{fig:Thist-all}) ) |
helps to prolong aud eublauce the epoch of star formation m nissbve galaxics. | helps to prolong and enhance the epoch of star formation in massive galaxies. |
The wariu-hot gas nass (WIT lob?οT«10* R3 is larecr iu the N-rav case than in the cases without N-ravs. | The warm-hot gas mass (WHIM; $10^{4.5}<T<10^7$ K) is larger in the X-ray case than in the cases without X-rays. |
Of course. we are especially interested in the eas which was collapsed aud virialized in deuse halos. | Of course, we are especially interested in the gas which has collapsed and virialized in dense halos. |
Figure 3 shows the temperature spectrun of high-density gas (p>200p. where pis the mean barvouic density of the universe) at +=0 for the four backgrounds aud the three ealaxies. ( | Figure \ref{fig:Thist-high} shows the temperature spectrum of high-density gas $\rho > 200\overline{\rho}$, where $\overline{\rho}$ is the mean baryonic density of the universe) at $z=0$ for the four backgrounds and the three galaxies. ( |
We note again that these sunulatious mclude uo optical depth effects aud therefore overestimate the UV dus that virialized regions sec.) | We note again that these simulations include no optical depth effects and therefore overestimate the UV flux that virialized regions see.) |
We soe that adding early UV does not affect the amount of cold (104K) deuse gas. aud has uucertaiu effect ou the hot (109I&) dense gas. jucreasing it slightle in galaxy €. mating uceheible change in galaxy E aud decreasing it somewhat in galaxy A (although uot for FO UV). | We see that adding early UV does not affect the amount of cold $10^4$ K) dense gas, and has uncertain effect on the hot $10^6$ K) dense gas, increasing it slightly in galaxy C, making negligible change in galaxy E and decreasing it somewhat in galaxy A (although not for FG UV). |
Adding an X-ray backeround substantially increases the hot deuse eas for ealaxics A and C. while having negligible effect on galaxy E We explain these differences by reference to the merger histories of the three galaxies. | Adding an X-ray background substantially increases the hot dense gas for galaxies A and C, while having negligible effect on galaxy E. We explain these differences by reference to the merger histories of the three galaxies. |
As explored in(2007).. galaxy A las a merger of mass ratio 6.5:1 at Doc0.6 (6 Gar ago). ealaxy € has a ierecr of mass ratio 3.5:dl at 2zιδ Cr ago). while galaxy E has no significant merger events after an equalimass merecr at 5c1.5 (10 Cir ago). | As explored in, galaxy A has a merger of mass ratio $6.5:1$ at $z\approx0.6$ (6 Gyr ago), galaxy C has a merger of mass ratio $3.5:1$ at $z\approx 0.8$ (7 Gyr ago), while galaxy E has no significant merger events after an equal-mass merger at $z\approx 1.5$ (10 Gyr ago). |
Since the gas-to-star ratio is a strouely declining function of halo mass. especially when there is significant ioniziug radiatiou to keep the low-deusity gas m sinall halos hot. we expect iu New UV|X for the accretion of smaller halos at later times to add more hot eas compared to accreting larger halos at carlicr times. aud the earlier the eas is added to the dense central galaxy. the more of it cau cool aud form stars in situ. | Since the gas-to-star ratio is a strongly declining function of halo mass, especially when there is significant ionizing radiation to keep the low-density gas in small halos hot, we expect in New UV+X for the accretion of smaller halos at later times to add more hot gas compared to accreting larger halos at earlier times, and the earlier the gas is added to the dense central galaxy, the more of it can cool and form stars in situ. |
Ou the other hand. the mergers in the New UV (no X) case involve large amounts of colder gas (see uext paragraph). which mineles with the existing eas aud forms stars quickly. thus paradoxically resulting m less hot eas at the preseut for Halo A. Fieve 1. shows the accretion rate of gas onto the central physical LOkpe of ealaxy A: galaxies € and E had qualitatively simular results. | On the other hand, the mergers in the New UV (no X) case involve large amounts of colder gas (see next paragraph), which mingles with the existing gas and forms stars quickly, thus paradoxically resulting in less hot gas at the present for Halo A. Figure \ref{fig:cgasacc-A100} shows the accretion rate of gas onto the central physical 10kpc of galaxy A; galaxies C and E had qualitatively similar results. |
We see munediatelv that New UV and New UV|X have significautlv higher accretion compared to Old UV. especially in the last [Gyr (2 <0.3) (and in the no X-ray case. the peak at ~8 Civr ago. correspouding to the merecr event mcutioned above): Εν UV has slehthy enhanced accretion. | We see immediately that New UV and New UV+X have significantly higher accretion compared to Old UV, especially in the last 4Gyr $z<0.3$ ) (and in the no X-ray case, the peak at $\sim 8$ Gyr ago, corresponding to the merger event mentioned above); FG UV has slightly enhanced accretion. |
This is | This is |
in this work. and for guidance on its implementation. | in this work, and for guidance on its implementation. |
We acknowledge helpful comments from the anonymous referee that improved this paper. | We acknowledge helpful comments from the anonymous referee that improved this paper. |
J.M.M. gratefully. acknowledges funding from the Guest Observer program. | J.M.M. gratefully acknowledges funding from the Guest Observer program. |
EMC gratefully acknowledges support provided by NASA through the Chandra Fellowship Program. grant number PF8-90052. | EMC gratefully acknowledges support provided by NASA through the Chandra Fellowship Program, grant number PF8-90052. |
RCR acknowledges STFC for financial support. | RCR acknowledges STFC for financial support. |
of our sources exhibit strong spatial variation in line intensity ratios. | of our sources exhibit strong spatial variation in line intensity ratios. |
The parameters of the four main lines(NgH*,HNC,,HCO*, HCN)) at their respective positions of maximum integrated intensity are listed in Table 4.. | The parameters of the four main lines, ) at their respective positions of maximum integrated intensity are listed in Table \ref{mainlines}. |
The large size of our data set requires that we set a high level of significance when searching for features to avoid many false positives. | The large size of our data set requires that we set a high level of significance when searching for features to avoid many false positives. |
With 187 sources, each with 16 lines and 31x31 pixels in each map we are searching for line detections in nearly 3 million spectra. | With 187 sources, each with 16 lines and $\times$ 31 pixels in each map we are searching for line detections in nearly 3 million spectra. |
A ba detection criteria should produce one false positive per 1.7 million measurements (for a perfectly normal distribution). | A $\sigma$ detection criteria should produce one false positive per 1.7 million measurements (for a perfectly normal distribution). |
We consider this to be an acceptably small level of contamination, and refer to a 5c detection as a robust detection. | We consider this to be an acceptably small level of contamination, and refer to a $\sigma$ detection as a robust detection. |
Additional selection criteria combined with a lower detection threshold could be used to search for additional weak lines. | Additional selection criteria combined with a lower detection threshold could be used to search for additional weak lines. |
For instance, to improve the completeness of ddetections we could adopt a lower o threshold while constraining the search to locations with significant fflux. | For instance, to improve the completeness of detections we could adopt a lower $\sigma$ threshold while constraining the search to locations with significant flux. |
stripped dwarf elliptical galaxy). and most ultraluminous X-ray sources (οἱ. | stripped dwarf elliptical galaxy), and most ultraluminous X-ray sources (cf. |
Table 2 of Nell. Ulvestad. Campion [2003]. and references therein). | Table 2 of Neff, Ulvestad, Campion [2003], and references therein). |
It is of interest to compare the Gl source (ο various relatives of pulsars as well. | It is of interest to compare the G1 source to various relatives of pulsars as well. |
For instance. Gl is within the wide range of both luminosity aud racdio/N-rav ratio observed for pulsar wind nebulae (DWNs) (Frail&Scharringhausen1997).. less luminous (han the putative PWN in M81 (Bietenholz.Bartel.&Rupen2004).. but considerably more Iuminous than standard pulsars or anomalous X-ray pulsars (Ilalpernetal.2005). | For instance, G1 is within the wide range of both luminosity and radio/X-ray ratio observed for pulsar wind nebulae (PWNs) \citep{fra97}, less luminous than the putative PWN in M81 \citep{bie04}, but considerably more luminous than standard pulsars or anomalous X-ray pulsars \citep{hal05}. |
. The 8.4 ζωα luminosity of GI is similar to that ofthe magnetar SGR 1806—20 about 10 days alter its outburst in late 2004. and the lack of a 4.9 GHz detection would be consistent with the facing ol SGR 1306—20 two months after the outburst (Gaensleretal.2005). | The 8.4 GHz luminosity of G1 is similar to that ofthe magnetar SGR $1806-20$ about 10 days after its outburst in late 2004, and the lack of a 4.9 GHz detection would be consistent with the fading of SGR $1806-20$ two months after the outburst \citep{gae05}. |
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